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+source,target
+ In Section 6 we discuss the properties of the overdensities. and the colours of the galaxies within the 3 protocluster candidates.," In Section \ref{discussion} we discuss the properties of the overdensities, and the colours of the galaxies within the 3 protocluster candidates."
+ Fluxes are calibrated using the Vega magnitude scale unless noted otherwise., Fluxes are calibrated using the Vega magnitude scale unless noted otherwise.
+ A flat ACDM cosmology is assumed throughout. with Όλι=0.3. Q4—0.7 and Hy=70kms |.," A flat $\Lambda$ CDM cosmology is assumed throughout, with $\Omega_{\rm M}=0.3$, $\Omega_{\Lambda}=0.7$ and $_0$ $^{-1}$."
+" The distance scale at z=24 isKKpe/"".. or ~ eco-movingMpc/""."," The distance scale at $z=2.4$ is, or $\sim$ co-moving."
+. Distances are given in co-moving units. unless stated otherwise.," Distances are given in co-moving units, unless stated otherwise."
+ Six HzRGs with redshifts between 2.28 and 2.55 were selected from the compendium of ?.., Six HzRGs with redshifts between 2.28 and 2.55 were selected from the compendium of \citet{MileydeBreuck2008}.
+ The targets were chosen based on their distributions in right ascension. declination. and redshift. and their bright radio luminosity «ΩΜΗ;L089 . and not on any prior information about surrounding galaxy overdensities.," The targets were chosen based on their distributions in right ascension, declination, and redshift, and their bright radio luminosity $_{500 \rm{MHz}}>10^{28.5}$ $^{-1}$ ), and not on any prior information about surrounding galaxy overdensities."
+ Co-ordinates of the selected targets and the control field are given in Table Ι.., Co-ordinates of the selected targets and the control field are given in Table \ref{tab:obs}.
+ The 6 HzRG fields and a blank control field were observed in service mode using the High Acuity Wide field K-band Imager (HAWK-I: 23) on the ESO Very Large Telescope (VLT) YEPUN telescope in Paranal. Chile during the period April-September 2008.," The 6 HzRG fields and a blank control field were observed in service mode using the High Acuity Wide field K-band Imager (HAWK-I; \citealt{Kissler-Patig2008}) ) on the ESO Very Large Telescope (VLT) UT4-YEPUN telescope in Paranal, Chile during the period April--September 2008."
+" HAWRK-I is a near-infrared camera comprising of four Hawaii-2 2048x2048 pixel detectors separated by a gap of 7-5"".", HAWK-I is a near-infrared camera comprising of four Hawaii-2 2048x2048 pixel detectors separated by a gap of $\sim$.
+ The camera spans 7.5x aaremin with a pixel scale of 0.106 aresee per pixel., The camera spans $7.5\times7.5$ arcmin with a pixel scale of 0.106 arcsec per pixel.
+ The telescope pointing was optimised so that the HZRG were placed close to the centre of the HAWK-I field of view. but at least aaremin away from the chip gaps. and bright stars did not fall within the field of view.," The telescope pointing was optimised so that the HzRG were placed close to the centre of the HAWK-I field of view, but at least arcmin away from the chip gaps, and bright stars did not fall within the field of view."
+ Each target was observed through theJ.. aand ‘filters and total exposure times for all targets are provided in reftab:obs..," Each target was observed through the, and filters and total exposure times for all targets are provided in \\ref{tab:obs}."
+ The telescope dithered every mmins which resulted in a I aaremin-wide cross-shaped region on the final mosaic where he image depth is shallower., The telescope dithered every mins which resulted in a $\sim$ arcmin-wide cross-shaped region on the final mosaic where the image depth is shallower.
+ During the observing period the anti-reflection coating of he Dewar window of HAWK-I was damaged resulting in small cross-shaped patterns on the exposures., During the observing period the anti-reflection coating of the Dewar window of HAWK-I was damaged resulting in small cross-shaped patterns on the exposures.
+ These patterns. caused by the spider of the secondary mirror. rotated as the telescope racked the targets as HAWK-I is situated on the Nasmyth focus.," These patterns, caused by the spider of the secondary mirror, rotated as the telescope tracked the targets as HAWK-I is situated on the Nasmyth focus."
+ These patterns were not adequately removed by the background subtraction step unless the telescope had not moved far between successive exposures., These patterns were not adequately removed by the background subtraction step unless the telescope had not moved far between successive exposures.
+ Thus the integration time between dithers was reduced from mmins to mmin for all images obtained after May 2008., Thus the integration time between dithers was reduced from mins to min for all images obtained after May 2008.
+ The data were reduced using the ESO/MVM 01 data reduction pipeline optimised for the reduction of our HAWK-I data., The data were reduced using the ESO/MVM \citep{Vandame2004} data reduction pipeline optimised for the reduction of our HAWK-I data.
+ The usual near-infrared reduction steps were taken. including dark subtraction. flat-field removal. harmonising the gain of the four detectors. removal of fringing. sky-subtraction. creation of bad pixel and weight maps. calculating the relative astrometry between the chips and the absolute astrometry.," The usual near-infrared reduction steps were taken, including dark subtraction, flat-field removal, harmonising the gain of the four detectors, removal of fringing, sky-subtraction, creation of bad pixel and weight maps, calculating the relative astrometry between the chips and the absolute astrometry."
+ USNO-BI catalogues (2) were used to calculate the guess astrometric solutions for the images., USNO-B1 catalogues \citep{Monet2003} were used to calculate the first-guess astrometric solutions for the images.
+ A catalogue was then compiled from the objects detected in the iimages and used to calculate the astrometric solution of the / and H images., A catalogue was then compiled from the objects detected in the images and used to calculate the astrometric solution of the $J$ and $H$ images.
+ The accuracy of the relative astrometry between the 3 images of each target is typically within ppixel. but the absolute astrometry is limited by the USNO-BI catalogue which has an accuracy of aaresec. equivalent to 2 HAWK-I pixels.," The accuracy of the relative astrometry between the 3 images of each target is typically within pixel, but the absolute astrometry is limited by the USNO-B1 catalogue which has an accuracy of arcsec, equivalent to 2 HAWK-I pixels."
+ After the sky background was removed from each exposure by MVM in a two step process. low level large-scale variations were still seen across each quadrant and most notably between the quadrants of the reduced images.," After the sky background was removed from each exposure by MVM in a two step process, low level large-scale variations were still seen across each quadrant and most notably between the quadrants of the reduced images."
+ Therefore the sky background of the tinal images was measured and subtracted using a local background estimator with SExtractor LOCAL)., Therefore the sky background of the final images was measured and subtracted using a local background estimator with SExtractor ).
+ The data were flux calibrated with 2MASS catalogues. using 13. I5.mmag stars within the target fields.," The data were flux calibrated with 2MASS catalogues, using $13-15.5$ mag stars within the target fields."
+ The calibration was checked5 using 13[4+ mmag stars within standard star fields taken within hhours of the observations., The calibration was checked using $13-14$ mag stars within standard star fields taken within hours of the observations.
+ Typically the zero-points determined from both methods agreed within the uncertainties (~0.05 mmag)., Typically the zero-points determined from both methods agreed within the uncertainties $\sim0.05$ mag).
+ The end products of this reduction process are a science image containing the reduced data of the target. and an effective exposure-time Image. which is an exposure-time map that has been normalised to account for differences in sensitivity between the four detectors of HAWK-I and the chip gaps.," The end products of this reduction process are a science image containing the reduced data of the target, and an effective exposure-time image, which is an exposure-time map that has been normalised to account for differences in sensitivity between the four detectors of HAWK-I and the chip gaps."
+ The 3 colour images of a target were convolved to match the lowest resolution image using the packagePSFMATCH., The 3 colour images of a target were convolved to match the lowest resolution image using the package.
+ The convolution kernels were optimised so that the stellar growth curves (created by median combining at least 20 bright and unsaturated stars in each image) converged to within at and beyond a 1” radius., The convolution kernels were optimised so that the stellar growth curves (created by median combining at least 20 bright and unsaturated stars in each image) converged to within at and beyond a $\arcsec$ radius.
+" The 56 image depths given in reftab:obs were measured by placing 2"" diameter apertures at multiple random positions,", The $\sigma$ image depths given in \\ref{tab:obs} were measured by placing $\arcsec$ diameter apertures at multiple random positions.
+ Each of the four HAWK-I detectors contain 32 amplifiers., Each of the four HAWK-I detectors contain 32 amplifiers.
+ Cross-talk between the amplifiers produce a series of artefacts arranged horizontally with respect to each star in the field at regular 64 pixel intervals (2).., Cross-talk between the amplifiers produce a series of artefacts arranged horizontally with respect to each star in the field at regular 64 pixel intervals \citep{finger2008}.
+ These artefacts appear crater-like and although are produced for every object in the field. they are only detectable above the noise if the star is brighter than approximately /=14.5 mmag (although this is strongly dependent on the seeing).," These artefacts appear crater-like and although are produced for every object in the field, they are only detectable above the noise if the star is brighter than approximately $J=14.5$ mag (although this is strongly dependent on the seeing)."
+ The artefacts Were most pronounced in the deep images., The artefacts were most pronounced in the deep images.
+" 2""x|"" rectangular regions were masked at 64 pixels intervals from each star brighter than J14.5 across the entiredetector quadrant.", $2\arcsec\times1\arcsec$ rectangular regions were masked at 64 pixels intervals from each star brighter than $J\simeq14.5$ across the entiredetector quadrant.
+ Cross-talk was greatly reduced in programs observed after ESO semester 82., Cross-talk was greatly reduced in programs observed after ESO semester 82.
+ Bright stars and nearby galaxies were also masked because they cover a significant amount of area in some fields and therefore, Bright stars and nearby galaxies were also masked because they cover a significant amount of area in some fields and therefore
+As the waves need more time to reach the eas located παμοι away from the ceuter. the heating rate rises at xogressively later times for more distant aunuli.,"As the waves need more time to reach the gas located further away from the center, the heating rate rises at progressively later times for more distant annuli."
+ Ouce the first wave has reached a given distance. viscous beating )ocomies comparable to the cooling rate.," Once the first wave has reached a given distance, viscous heating becomes comparable to the cooling rate."
+ This is cousisteut with heating rate predictions made by Fabianotal.(2003a).. also assuniug Spitzer viscosity.," This is consistent with heating rate predictions made by \citet{fab03a}, also assuming Spitzer viscosity."
+ We also note hat dissipating waves of ereater initial amplitude in our simulations would eive even more heating to offset cooling., We also note that dissipating waves of greater initial amplitude in our simulations would give even more heating to offset cooling.
+ Tuterestinely. the average ratio of heating to cooling secus o be relatively stable as a function of time.," Interestingly, the average ratio of heating to cooling seems to be relatively stable as a function of time."
+ We have also computed the vohune-inteerated heating aud cooling rates and found that their ratio couverges to a value of the order of a few., We have also computed the volume-integrated heating and cooling rates and found that their ratio converges to a value of the order of a few.
+ However. the balance of heating aud cooling is uot automatic as it depends ou the choice of parameters (0.8. ACN power and deusitv eradieut in the intracluster mediunu) aud here feedback may plav a role.," However, the balance of heating and cooling is not automatic as it depends on the choice of parameters (e.g., AGN power and density gradient in the intracluster medium) and here feedback may play a role."
+ Note that the curves display a pronounced periodic behavior., Note that the curves display a pronounced periodic behavior.
+ This reflects the imutenuittenev of the ceutral source. with on- and off-states of 1.5«10* vears;," This reflects the intermittency of the central source, with on- and off-states of $1.5\times 10^{7}$ years."
+ This is cousisteut with the observational estimates based on observations of ripples iu the Perseus cluster (Fabianetal...2003a.b).," This is consistent with the observational estimates based on observations of ripples in the Perseus cluster \citep{fab03a,fab03b}."
+. We performed a series of numerical experiments to mvoestieate whether a single ACN outburst can generate waves for which the dissipation rates could offset local radiative cooling rates., We performed a series of numerical experiments to investigate whether a single AGN outburst can generate waves for which the dissipation rates could offset local radiative cooling rates.
+ These simulations demonstrated that. whereas secondary waves ecucrated by the interaction of the rising bubble with the surrounding iutraclustor medi are clearly present. the viscous heating associated with a single outburst is insufficient to balance radiative cooling.," These simulations demonstrated that, whereas secondary waves generated by the interaction of the rising bubble with the surrounding intracluster medium are clearly present, the viscous heating associated with a single outburst is insufficient to balance radiative cooling."
+ This suggests that the ripples observed in the Perseus cluster can be interpreted as beiug due to the AGN duty cvcle. ie.. they trace ACN activity The work doue by the expanding cavities on the ambieut medium is limited to a modest fraction of the οποιον injected by the ACN.," This suggests that the ripples observed in the Perseus cluster can be interpreted as being due to the AGN duty cycle, i.e., they trace AGN activity The work done by the expanding cavities on the ambient medium is limited to a modest fraction of the energy injected by the AGN."
+ If the cavities are approximately in pressure balance with their surroundings. the iucrease in the enerey of the ambicut eas. dU.=d(PV)/(1). is related to the work done. (IT=PdV. by dU~AVA(> 1) ," If the cavities are approximately in pressure balance with their surroundings, the increase in the energy of the ambient gas, $dU = d(PV)/(\gamma - 1)$, is related to the work done, $dW = P \, dV$, by $dU \simeq dW/ (\gamma -1)$ ."
+The first law of thermodynamics then implies that (ITx-- where dQ is the heat injected into the cavity.," The first law of thermodynamics then implies that $dW \simeq {\gamma - 1
+\over \gamma} dQ$, where $dQ$ is the heat injected into the cavity."
+ This meaus that. depending on the effective value of 5 (which can range between 1/3 and 5/3). 25—LO% of the energv input can be trausterred to the züubient medi (see also. c.e.. Chirazov ct al.," This means that, depending on the effective value of $\gamma$ (which can range between 4/3 and 5/3), $25-40\%$ of the energy input can be transferred to the ambient medium (see also, e.g., Churazov et al."
+ 2001)., 2001).
+ The fraction of the iuput power transterred to the ICAL will be laveer if the cavities are overpressured., The fraction of the input power transferred to the ICM will be larger if the cavities are overpressured.
+" The fraction of this work that goes iuto acoustic enerev. as opposed to other types of disturbance (οι, e-modes or internal waves). depends on the timescale of pressure fluctuations. as well as detailed structure of the cavityICAL interface."," The fraction of this work that goes into acoustic energy, as opposed to other types of disturbance (e.g., g-modes or internal waves), depends on the timescale of pressure fluctuations, as well as detailed structure of the cavity–ICM interface."
+ We expect the production of sound waves to be efficicut wheu the AGN duty evcle is of the same order as the sound crossing time at the cavity radius. or shorter.," We expect the production of sound waves to be efficient when the AGN duty cycle is of the same order as the sound crossing time at the cavity radius, or shorter."
+ This condition is satisfied for our choscn duty exele of 3«10* vr., This condition is satisfied for our chosen duty cycle of $3 \times 10^7$ yr.
+ In addition to work done by in situ expansion of the cavities. a roughly comparable amount of euergv is transferred to the surrounding medimm. iu the form of kinetic and eravitational potential energv. as the cavities vise through the backgound pressure eradieut.," In addition to work done by in situ expansion of the cavities, a roughly comparable amount of energy is transferred to the surrounding medium, in the form of kinetic and gravitational potential energy, as the cavities rise through the backgound pressure gradient."
+ The latter is the generic mechanisin appealed to bv Beechuan (2001) and Ruszkowski Beechuan (2002) iu thei discussion of “effervescent heating”.," The latter is the generic mechanism appealed to by Begelman (2001) and Ruszkowski Begelman (2002) in their discussion of “effervescent heating""."
+ Enerey injected in this way can also be converted to heat through viscous dissipation., Energy injected in this way can also be converted to heat through viscous dissipation.
+ Our sinulatious map the total viscous dissipation rate. aud do not distinguish between dissipation of sound waves and other kinds of motion.," Our simulations map the total viscous dissipation rate, and do not distinguish between dissipation of sound waves and other kinds of motion."
+ Note. however. that sound waves have lareer propagation speeds than other modes. aud therefore should progressively dominate the energeties at radii well outside the Although we have devised a specific model iu which the viscous dissipation rate of sound waves roughly balances local radiative cooling. such a balance may not be a universal property of ACN heating iu cluster cores.," Note, however, that sound waves have larger propagation speeds than other modes, and therefore should progressively dominate the energetics at radii well outside the Although we have devised a specific model in which the viscous dissipation rate of sound waves roughly balances local radiative cooling, such a balance may not be a universal property of AGN heating in cluster cores."
+ The distribution of sound cucrey dissipation is larecly determined by the radial structure of the model., The distribution of sound energy dissipation is largely determined by the radial structure of the model.
+ The sound dissipation leneth for a fixed waveleneth L(A.1) for the paraincters in our simulations decreases from the center to the edge of the simulated region. mainly due to the decrease in density (Fabian et al.," The sound dissipation length for a fixed wavelength $L(\lambda, r)$ for the parameters in our simulations decreases from the center to the edge of the simulated region, mainly due to the decrease in density (Fabian et al."
+ 2003)., 2003).
+ For the paralcters chosen in our simulation the characteristic dissipation leusth near the outer edge of the exid is of order the size of the simulation region. implying that the dissipation is spread over a volune that far exceeds that of the bubbles. aud ach of the acoustic energy goes iuto heating.," For the parameters chosen in our simulation the characteristic dissipation length near the outer edge of the grid is of order the size of the simulation region, implying that the dissipation is spread over a volume that far exceeds that of the bubbles, and much of the acoustic energy goes into heating."
+ The steady rate of production of acoustic energv then leads to a rough balance between heating aud cooling. eiven the adopted deusity aud temperature profile.," The steady rate of production of acoustic energy then leads to a rough balance between heating and cooling, given the adopted density and temperature profile."
+ These conditious may not be satisfied in all clusters., These conditions may not be satisfied in all clusters.
+ As the sound dissipation leusth is proportional to the square of the period of the sound waves. more frequent outbursts should lead to more ceutrally concentrated dampine.," As the sound dissipation length is proportional to the square of the period of the sound waves, more frequent outbursts should lead to more centrally concentrated damping."
+ Uowever. the dissipation rate does uot depend on the ACN intermittency period as such. since the pressure pulses ecuerated by the bubbles are likely to be far from sinusoidal aud will contain a wide rauge of frequencies.," However, the dissipation rate does not depend on the AGN intermittency period as such, since the pressure pulses generated by the bubbles are likely to be far from sinusoidal and will contain a wide range of frequencies."
+" The dispersion of the waves as they propagate sugecst that the ""effective"" wavelength will increase with r. au effect that will partially counteract the decrease of L(A.1) with +."," The dispersion of the waves as they propagate suggest that the “effective"" wavelength will increase with $r$, an effect that will partially counteract the decrease of $L(\lambda, r)$ with $r$."
+" Where the velocity field has small-scale structure or where the damping rate is much lareer than the Spitzer rate, acoustic waves (as well as eravity aud internal waves) can be dissipated much closer to the sites where they are eecuerated."," Where the velocity field has small-scale structure or where the damping rate is much larger than the Spitzer rate, acoustic waves (as well as gravity and internal waves) can be dissipated much closer to the sites where they are generated."
+ Distributed heating would then occur oulv after he bubbles had penetrated most of the cluster., Distributed heating would then occur only after the bubbles had penetrated most of the cluster.
+ This is the situation envisaged by Begehuan (2001) and Ruszkowski Beechuan (2002) in the effervescent heating scenario., This is the situation envisaged by Begelman (2001) and Ruszkowski Begelman (2002) in the effervescent heating scenario.
+ This orm of heating may be occumiug concurrently with the argc-scale acoustic heating iu Perseus. aud may dominate he heating in other clusters (e.g. those with smaller acoustic energev ecucration due to the intermittency xoperties of the central AGN).," This form of heating may be occurring concurrently with the large-scale acoustic heating in Perseus, and may dominate the heating in other clusters (e.g., those with smaller acoustic energy generation due to the intermittency properties of the central AGN)."
+ Note that viscosity may welp the bubbles penetrate to large distances without excessive LUNI., Note that viscosity may help the bubbles penetrate to large distances without excessive mixing.
+ We stress that our two-dimensional simulations do not accurately represent the behavior of three-dineusional acoustic heating in several respects., We stress that our two-dimensional simulations do not accurately represent the behavior of three-dimensional acoustic heating in several respects.
+ In three dimensions, In three dimensions
+small effect.,small effect.
+ For example. ife =1. changing Z bv a factor 5 would only change rias by ~20%.," For example, if $\epsilon=-1$, changing $Z$ by a factor 5 would only change $r_{\rm max}$ by $\sim20\%$."
+ The existence of a cutolf in the distribution bevond some size is naturally expected in the above picture. given the existence of some minimum density ys in the regions of the galaxies where SNe explode. and in view of the decrease Of Puuas With raciius.," The existence of a cutoff in the distribution beyond some size is naturally expected in the above picture, given the existence of some minimum density $\rho_{\rm min}$ in the regions of the galaxies where SNe explode, and in view of the decrease of $\rho_{\rm max}$ with radius."
+ At some radius. gas will equal fiiia and the integral in Eq.," At some radius, $\rho_{\rm max}$ will equal $\rho_{\rm min}$ and the integral in Eq."
+ 13. will therefore become zero., \ref{dNdr} will therefore become zero.
+ In other words. there will be nowhere in the MCSs a region with a density low enough to permit a SNR of that size that is still in its bright Sedov phase.," In other words, there will be nowhere in the MCs a region with a density low enough to permit a SNR of that size that is still in its bright Sedov phase."
+ A deficit of SNRs at small racii. as observed in M33 ane the MCs. is also expected in this scenario.," A deficit of SNRs at small radii, as observed in M33 and the MCs, is also expected in this scenario."
+ Before the onse of the Sedov stage. faster shock velocities will lead to fewer objects observed. in the bins with the smallest. radii.," Before the onset of the Sedov stage, faster shock velocities will lead to fewer objects observed in the bins with the smallest radii."
+" This onset happens at ages (sizes) that depend on the details of the ejecta structure. as well as the ambient density (see87in""Truelove&Melxee 1999).. but for most SNRs it should occur around a few hundred vears (a few pc). which is consisten with the deficits that we have observed in M33 and the MICs."," This onset happens at ages (sizes) that depend on the details of the ejecta structure, as well as the ambient density \citep[see \S~7 in][]{truelove99:adiabatic-SNRs}, but for most SNRs it should occur around a few hundred years (a few pc), which is consistent with the deficits that we have observed in M33 and the MCs."
+ This regime alfects only a small number of objects in the ALCs. and does not impact any of the arguments made above. so we will ignore it for the remainder of the paper.," This regime affects only a small number of objects in the MCs, and does not impact any of the arguments made above, so we will ignore it for the remainder of the paper."
+ We have also ignored the deviations from the standard evolutionary picture that can be introduced by the shape of he circumstellar medium excavated by the SN progenitors., We have also ignored the deviations from the standard evolutionary picture that can be introduced by the shape of the circumstellar medium excavated by the SN progenitors.
+ Baclenesοἱal.(2007) showed that most Type la SNRs with known ages have sizes that are consistent with an interaction witha uniform ambient medium. but no such study has been done for CC SNRs.," \citet{badenes07:outflows} showed that most Type Ia SNRs with known ages have sizes that are consistent with an interaction with a uniform ambient medium, but no such study has been done for CC SNRs."
+ The fast stellar outflows expected from he more massive CC SN progenitors will modify the sizes ofa few individual SNlis at certain stages of their evolution (c.g..Dwarkaclas2005.2007).. but most of the objects that we consider here are too large to be expanding in even the most extreme wind-blown cavities.," The fast stellar outflows expected from the more massive CC SN progenitors will modify the sizes of a few individual SNRs at certain stages of their evolution \citep[e.g.,][]{dwarkadas05:SNR-Bubbles_1D,dwarkadas07:SNRs_Bubbles_WR}, but most of the objects that we consider here are too large to be expanding in even the most extreme wind-blown cavities."
+ As long as the bulk of he SNRs in the sample spend most of their lifetimes in the Seclov stage. this should not allect our scenario.," As long as the bulk of the SNRs in the sample spend most of their lifetimes in the Sedov stage, this should not affect our scenario."
+ Similarly. he fact that some SNRs evolve inside superbubbles (e.g.AlacLow&AleCray1988) is naturally incorporated. into our picture superbubbles merely become one more of the actors driving the density distribution in the interstellar medium.," Similarly, the fact that some SNRs evolve inside superbubbles \citep[e.g.][]{maclow88:superbubbles} is naturally incorporated into our picture – superbubbles merely become one more of the factors driving the density distribution in the interstellar medium."
+ Incidentally. the size distribution of superbubbles also relates to the properties of the interstellar medium. as shown bv Oev&Clarke(1997). [or several nearby. galaxies. including the SAIC.," Incidentally, the size distribution of superbubbles also relates to the properties of the interstellar medium, as shown by \citet{oey97:superbubble_sizes} for several nearby galaxies, including the SMC."
+ As we have seen. a uniform SNR. size distribution can be understood as he result of SecOV expansion. οςmbined with a transition to the radiative pdase al an age hat depends on the local cknsbÜv. provided that the clensiv of the gas in the interstelar medium [οlows a clistribution close to a power law with an index. of 1. Pidp~p7.," As we have seen, a uniform SNR size distribution can be understood as the result of Sedov expansion, combined with a transition to the radiative phase at an age that depends on the local density, provided that the density of the gas in the interstellar medium follows a distribution close to a power law with an index of $-1$, $dP/d\rho\sim \rho^{-1}$."
+ In this Section. we tes this hypothesis by examining hree indirect tracers of gas censity in the Aagellanie Cloucs: HIE column density: star-formation rate (SER) based on resolved stellar populations: and Ho emission-ine surface briginess.," In this Section, we test this hypothesis by examining three indirect tracers of gas density in the Magellanic Clouds: HI column density; star-formation rate (SFR) based on resolved stellar populations; and $\alpha$ emission-line surface brightness."
+ These tracers are wel suited for our goals because they are valid over a wide range of densities. and the necessary data are available [rom public surveys that cover the whole extent of the Clouds. as deseribed in detail below.," These tracers are well suited for our goals because they are valid over a wide range of densities, and the necessary data are available from public surveys that cover the whole extent of the Clouds, as described in detail below."
+ We have taken the surface brightness of HE 21 cm line emission in the ALCS from the maps of Iximetal.(2003) and Stanimiroviectal.(1999)... which combine single-dish Parkes and aperture-svnthesis APCA data to. probe both small and laree scales in the LAIC and the SAIC. respectively.," We have taken the surface brightness of HI 21 cm line emission in the MCs from the maps of \citet{kim03:LMC_HI_Parkes_ATCA} and \citet{stanimirovic99MNRAS.302..417S}, which combine single-dish Parkes and aperture-synthesis ATCA data to probe both small and large scales in the LMC and the SMC, respectively."
+ “Phe 21 cem emission is optically thin. so the surface brightness is directly. proportional to the HIE column density.," The 21 cm emission is optically thin, so the surface brightness is directly proportional to the HI column density."
+ Since the LMC possesses a fairly face-on (inclinationPee357. 2001). wellorderecl LIE clisk. the column density should. in turn. be roughly. proportional to the volume density p.," Since the LMC possesses a fairly face-on \citep[inclination $i\sim 35^{\circ}$ , well-ordered HI disk, the column density should, in turn, be roughly proportional to the volume density $\rho$."
+ Iximal.(2007) report that the LIL column density distribution ofet individual “clouds” of neutral hvdrogen in the LMC follows à log-normal form. rather than a power law.," \citet{kim07:HI_Clouds_LMC} report that the HI column density distribution of individual “clouds” of neutral hydrogen in the LMC follows a log-normal form, rather than a power law."
+" However. their figure 13 suggests that. above a low cutoll of 2107""em.7. the distribution does behave as a power law of slope l.over at least an order of magnitude."," However, their figure 13 suggests that, above a low cutoff of $2\times 10^{20}~{\rm cm}^{-2}$, the distribution does behave as a power law of slope $\sim -1$ , over at least an order of magnitude."
+ ‘To re-examine this. in Figure 5 we show the cillerentia distribution of HIE column density in the individual sized pixels of the Iximetal.(2003). LMC map 2007)..," To re-examine this, in Figure \ref{HIHist} we show the differential distribution of HI column density in the individual beam-sized pixels of the \citet{kim03:LMC_HI_Parkes_ATCA} LMC map \citep[as opposed to
+the cumulative plot for ``clouds'' shown in][]{kim07:HI_Clouds_LMC}."
+" We see that the LL column in the LMC does follow an index l power Law fairlv well. between a column of 3.107""em ?2and 61073em.7."," We see that the HI column in the LMC does follow an index $-1$ power law fairly well, between a column of $3\times 10^{20}~{\rm cm}^{-2}$ and $6\times 10^{21}~{\rm
+ cm}^{-2}$."
+ The observed low cutoll in the column density is unavoidable because of the integration through the disk and over the beam size (every line of sigh is basically sampling the densest regions at that point)., The observed low cutoff in the column density is unavoidable because of the integration through the disk and over the beam size (every line of sight is basically sampling the densest regions at that point).
+ In regions with low density. the LE atoms might be ionized. as in the swarm ionized” phase of the interstellar medium (Ferriere2001).. so the tracer may become loss reliable there.," In regions with low density, the H atoms might be ionized, as in the “warm ionized” phase of the interstellar medium \citep{ferriere01:ISM}, so the tracer may become less reliable there."
+ ]t is quite plausible that. in the regions where SNe actually explode. the underlving distribution of densities also reaches a minimum as we recall. a minimum density. is required in order to reproduce the observed upper cutoll in SNIt size rut this does not necessarily correspond with the lower threshold in the LIE distribution.," It is quite plausible that, in the regions where SNe actually explode, the underlying distribution of densities also reaches a minimum – as we recall, a minimum density is required in order to reproduce the observed upper cutoff in SNR size – but this does not necessarily correspond with the lower threshold in the HI distribution."
+" To illustrate this. we have calculated the mean LIE columns in each of the spatial ""cells"" defined by Llarris&Zaritsky(2004) and Harris&Zaritsky(2009) that contain SNRs (see 5.2 below for adescription of the cells). which we display with the horizontal rulers in Figure 5.."," To illustrate this, we have calculated the mean HI columns in each of the spatial “cells” defined by \citet{harris04:SMC_SFH} and \citet{harris09:LMC_SFH} that contain SNRs (see \ref{sec:SFR} below for adescription of the cells), which we display with the horizontal rulers in Figure \ref{HIHist}."
+ We note that. in the LMC. these cells have average column values between ο10720em72 (close to. but higher: than the low LIE cutoll) and 61072lem2τι although most SNRs appear clustered around 21073em7.," We note that, in the LMC, these cells have average column values between $5\times10^{20}~{\rm cm}^{-2}$ (close to, but higher than the low HI cutoff) and $6\times 10^{21}~{\rm cm}^{-2}$, although most SNRs appear clustered around $2\times 10^{21}~{\rm cm}^{-2}$."
+ As shown in Figure 5..5.. the cistribution of HIE column densities in the SAIC is flat over the same range. although the rise ancl fall from the plateau happen at the same densities as the rise and fall of the powerlaw in the LMC.," As shown in Figure \ref{HIHist}, the distribution of HI column densities in the SMC is flat over the same range, although the rise and fall from the plateau happen at the same densities as the rise and fall of the powerlaw in the LMC."
+ While we do not know the reason for this. we speculate that it may. be related to an SAIC geometry. that is elongated along our line of sight. and the integration ellect that results.," While we do not know the reason for this, we speculate that it may be related to an SMC geometry that is elongated along our line of sight, and the integration effect that results."
+ The actual “depth” of the SAIC. whether just a few kpe or as much as 20 kpe. is debated. (Llatzicimitriou&Llawkins 2000)... but is likely at least a few times larger than that of the nearly [ace-on. LMC.," The actual “depth” of the SMC, whether just a few kpc or as much as 20 kpc, is debated \citep{hatzidimitriou89:SMC_structure,harris04:SMC_SFH,subramanian09:LMC_SMC_Depth}, , but is likely at least a few times larger than that of the nearly face-on LMC."
+ Such an integration elfect would explain why the SAIC SNRsare found at HE column, Such an integration effect would explain why the SMC SNRsare found at HI column
+Dynamics auc cor'espouding Laclative signatwes of nou-spherical relativistic shocks remains an linportaut unresolved issues tu studies on Gatuna Ray Bursts (GRBs).,Dynamics and corresponding radiative signatures of non-spherical relativistic shocks remains an important unresolved issues in studies on Gamma Ray Bursts (GRBs).
+ Since GRBs produce uarrowlv collimated οullows that evolve lateralN. uuderstandiug the overall dyuamies - both theoretical aud in terns of agreement between cliΠοιοί numerical results - is imperative to the iuterpretation of the 'oadbaud observatious of GRBs ??..," Since GRBs produce narrowly collimated outflows that evolve laterally, understanding the overall dynamics - both theoretical and in terms of agreement between different numerical results - is imperative to the interpretation of the broadband observations of GRBs \cite{Rhoads99,Frail01}."
+ Presently. there aο two coimpeti ews ο he lateral evolution of the relativistic outflows.," Presently, there are two competing views on the lateral evolution of the relativistic outflows."
+ Theoretically. it is twpically argued t je lateral evolution of the flow proceeds with relativistic velocities ?).(2).. (seealso? )..," Theoretically, it is typically argued that the lateral evolution of the flow proceeds with relativistic velocities \citep{PiranReview}, \citep[see also][]{2011arXiv1102.5618W}."
+ This view Οἱαςicted by he results of numerical simulations that show very little lateral evolΠιο in t alivistic reelme ?TTT?..," This view is contradicted by the results of numerical simulations that show very little lateral evolution in the relativistic regime \cite{2004ApJ...601..380C,2009ApJ...698.1261Z,2010A&A...520L...3M,2011arXiv1105.2485V}."
+ In this Letter we argue hat thisdisagreemen results L'om the incorrect theoretical assumptions about the lateral evolution of the flow., In this Letter we argue that this disagreement results from the incorrect theoretical assumptions about the lateral evolution of the flow.
+ Wjab is linj»ortant for the interpretation of observations is the evolution of a curved stock., What is important for the interpretation of observations is the evolution of a curved shock.
+ Previously t aeral evolution of the nor-spherical shocks was incorrectly treated as a f‘ee lateral expansion LO vacuttn (e.g..?.Eq.5) 2s , Previously the lateral evolution of the non-spherical shocks was incorrectly treated as a free lateral expansion into vacuum \citep[\eg][Eq. 5]{2011arXiv1102.5618W}. ..
+"The assumption of the lateral expausion witl Iie souud speed results 1 lagramophone-type"" p‘oliles aud. slowing down of the ejeca.", The assumption of the lateral expansion with the sound speed results in a “gramophone-type” profiles and slowing down of the ejecta.
+ This has drastic iinplicatious for the uuderlviug light curves (eg??)..," This has drastic implications for the underlying light curves \citep[eg][]{2000ApJ...541L...9K,2003ApJ...592..390P}."
+" Iu fact the dynamics of tle nou-sphlierical shocks is uore subtle: the correct treatineut. as we argue below. is cousistent with slow lateral evolution ""eel in uuimerical simulations."," In fact the dynamics of the non-spherical shocks is more subtle; the correct treatment, as we argue below, is consistent with slow lateral evolution seen in numerical simulations."
+ Evolution of strong ion-spherical shocks is a well studies problems iu fluid dynanics., Evolution of strong non-spherical shocks is a well studies problems in fluid dynamics.
+ The two fuudamental works that have laid the founcatiou for nou-sphlerical (two-dimeusional) shocks.," The two fundamental works that have laid the foundation for non-spherical (two-dimensional) shocks,"
+was located in the inter-binary system.,was located in the inter-binary system.
+ However. further work (eg Juett Chakrabarty 2005) showed that for individual sources. the Ne/O ratio showed evidence for variability from epoch to epoch. which they attributed. to source variability.," However, further work (eg Juett Chakrabarty 2005) showed that for individual sources, the Ne/O ratio showed evidence for variability from epoch to epoch, which they attributed to source variability."
+ This implied that the abundances could not be used to determine the composition of the mass donating star., This implied that the abundances could not be used to determine the composition of the mass donating star.
+ There is a clear similarity between the neutron star UCBs described by Juctt et al and RX 1914|24., There is a clear similarity between the neutron star UCBs described by Juett et al and RX J1914+24.
+ 1n each observation of RN J1914|24. there is clear evidence that the absorption component has an overabundance of neon.," In each observation of RX J1914+24, there is clear evidence that the absorption component has an overabundance of neon."
+ For the reasons outlined in 85.1. we rule out AA 1914|24 being an isolated neutron star.," For the reasons outlined in \ref{ins}, we rule out RX J1914+24 being an isolated neutron star."
+ Since all the known neutron star UCBs have X-ray emission extending up to many 10s of keV we also rule out an accreting neutron star UCB model., Since all the known neutron star UCBs have X-ray emission extending up to many 10's of keV we also rule out an accreting neutron star UCB model.
+ We cannot rule out that a neutron star is in a binary system where a secondary star was not filling its Roche Lobe., We cannot rule out that a neutron star is in a binary system where a secondary star was not filling its Roche Lobe.
+ In this scenario. an N-rav. bright svstemi would have to be powered by Ul.," In this scenario, an X-ray bright system would have to be powered by UI."
+ Lt is highly. unlikely that the line of sight absorption to RN 121914|24 has a chance enhancement of neon., It is highly unlikely that the line of sight absorption to RX J1914+24 has a chance enhancement of neon.
+ LH is much more likely that this overabundance is concentrated in the binary svstem., It is much more likely that this overabundance is concentrated in the binary system.
+ Juett Chakrabarty (2005) noted that for some neutron star UCBs the Ne/O abundance varied. from epoch to epoch and hence the observations could not be used to determine the abundance of the secondary. mass-donating star. in the binary svstem.," Juett Chakrabarty (2005) noted that for some neutron star UCBs the Ne/O abundance varied from epoch to epoch and hence the observations could not be used to determine the abundance of the secondary, mass-donating star, in the binary system."
+ In the case of RN 1914|24 there is clear evidence for a significant. over-abundance of neon in the absorption Component at each epoch., In the case of RX J1914+24 there is clear evidence for a significant over-abundance of neon in the absorption component at each epoch.
+ At this stage it is not clear if this over-abuncance is due to circumbinary materia eft over from a previous stage in the binary formation »Focess or can give us a direct. insight into the chemica composition of the secondary star (if accretion is occurring)., At this stage it is not clear if this over-abundance is due to circumbinary material left over from a previous stage in the binary formation process or can give us a direct insight into the chemical composition of the secondary star (if accretion is occurring).
+ What are the implications of our findings reearcing he X-rav luminositv of RA .1914]|24?, What are the implications of our findings regarding the X-ray luminosity of RX J1914+24?
+7 Stecehs et a (2006) discuss the extinction and distance estimates to UX J1914]24., Steeghs et al (2006) discuss the extinction and distance estimates to RX J1914+24.
+ While the distance is rather uncertain. it is likely that it is greater than 1 kpe.," While the distance is rather uncertain, it is likely that it is greater than $\sim$ 1 kpc."
+ We can rule ou he lower estimates (Lx~10%? Hor a distance of kpe) which were derived. using a low temperature thermal plasma model., We can rule out the lower estimates $L_\mathrm{X}\sim10^{33}$ for a distance of 1 kpc) which were derived using a low temperature thermal plasma model.
+ Taking the unabsorbed bolometric [uxes derived using the blackhocky with absorption component with variable abundances ancl assuming a distance of 1 kpe we find Ly—2.105.L6.107?," Taking the unabsorbed bolometric fluxes derived using the blackbody with absorption component with variable abundances and assuming a distance of 1 kpc we find $L_\mathrm{X}=2\times10^{34} -
+1.6\times10^{35}$."
+ DallOsso ct al (2007) made a detailed. investigation of the Ul model in the context of RN J1914|24 and RX JOSOG|15., Dall'Osso et al (2007) made a detailed investigation of the UI model in the context of RX J1914+24 and RX J0806+15.
+ They. predicted that for low luminosities. Lx105 the asvnchronism between the orbit. and he magnetic star in RA 1914|24 would have to be a~1.9 0.98. where a=wjfe. andy is the rotation frequency of the primary star and cy is the orbital frequency.," They predicted that for low luminosities, $L_\mathrm{X}\sim10^{33}$ , the asynchronism between the orbit and the magnetic star in RX J1914+24 would have to be $\alpha\sim0.9-$ 0.98, where $\alpha=\omega_{1}/\omega_{o}$, and $\omega_{1}$ is the rotation frequency of the primary star and $\omega_{o}$ is the orbital frequency."
+ Unless UN. J1914124 was located at a distance significantly. less han Ikpce. we can rule these low values of asynchronism.," Unless RX J1914+24 was located at a distance significantly less than 1kpc, we can rule these low values of asynchronism."
+ bor ο] luminosities (Lx=107 13). Dall'Osso et al (2007) predicted that an asvnchronism of a few was required (o~ 4).," For high luminosities $L_\mathrm{X}=10^{34-35}$ ), Dall'Osso et al (2007) predicted that an asynchronism of a few was required $\alpha\sim$ 4)."
+ For an observed period of 569 sec. à=5.10 gives a predicted: period. of 760-300 sec.," For an observed period of 569 sec, $\alpha=2-10$ gives a predicted period of $\sim$ 60-300 sec."
+ There is no evidence for power at these periods in the power spectra of the X-ray ight curves (Ramsay ct al 2006)., There is no evidence for power at these periods in the power spectra of the X-ray light curves (Ramsay et al 2006).
+ Until now the nature of the emission source that powers the X-ray spectrum of RA J1914124 has been far from clear., Until now the nature of the emission source that powers the X-ray spectrum of RX J1914+24 has been far from clear.
+ In this paper we have shown that it can be well modelled using a simple blackbody model. with an absorption component which has non-solar abuncances. in particular. an enhancement of neon.," In this paper we have shown that it can be well modelled using a simple blackbody model with an absorption component which has non-solar abundances, in particular, an enhancement of neon."
+ Since the X-ray light curves of RX 1914|24 and RX JOSOG|15 are practically identical. it suggests that their X-ray emission source is the same.," Since the X-ray light curves of RX J1914+24 and RX J0806+15 are practically identical, it suggests that their X-ray emission source is the same."
+ The fact wt their X-pav spectra were apparently different (albeit both being soft) was therefore. perplexing., The fact that their X-ray spectra were apparently different (albeit both being soft) was therefore perplexing.
+ Our result. showing. that the emission. source is the same for both RN 1914|24 and RX JOSOG|15 is therefore very attractive., Our result showing that the emission source is the same for both RX J1914+24 and RX J0806+15 is therefore very attractive.
+ Indeed. their temperatures are virtually identical owe obtain a mean value of AL~67 eV for RN JL914|24 compared to kl65 eV for RX JOSOG|15 Csracl οἱ al 2003)., Indeed their temperatures are virtually identical – we obtain a mean value of $kT\sim67$ eV for RX J1914+24 compared to $kT\sim65$ eV for RX J0806+15 (Israel et al 2003).
+ The cillerence between the X-ray spectrum. of RX J1914]24 and RA JOSOG|15 is that the absorption component of IUX J1914|24 has enhanced neon abundance., The difference between the X-ray spectrum of RX J1914+24 and RX J0806+15 is that the absorption component of RX J1914+24 has enhanced neon abundance.
+were acceptable.,were acceptable.
+ Such spectra contrast with the cool optically thick coronae obtained when modeling high huuinositv neutron star binaries with (e.g.7).," Such spectra contrast with the cool, optically thick coronae obtained when modeling high luminosity neutron star binaries with \citep[e.g.][]{disalvo00}."
+ The disk blackbody fit wasacceptable (4? /dof = 110/117): however the inner disk temperature (kta) was 3.20.1 keV. which is too hot for a black hole high state.," The disk blackbody fit wasacceptable $\chi^2$ /dof = 140/147); however the inner disk temperature $T_{\rm in}$ ) was $\pm$ 0.4 keV, which is too hot for a black hole high state."
+ Hence. we reject this fit ou physical grounds. even though it is a statistically acceptable fit to the data.," Hence, we reject this fit on physical grounds, even though it is a statistically acceptable fit to the data."
+ The two component inodel vielded Αι = I«1075402 and D — ME12. as with. the simple. power aw quodol.," The two component model yielded $N_{\rm H}$ = $\times 10^{21}$ and $\Gamma$ = 1.2, as with the simple power law model."
+" Furthermore. NSPEC was unable o estimate the uncertainties in the blackbody xuineters,"," Furthermore, XSPEC was unable to estimate the uncertainties in the blackbody parameters."
+ We therefore couchide that αν hermal componcut is too marginal to be detected: rence. NBOs2 is unlikely to be in the high accretion rate state for neutron star LAINBs.," We therefore conclude that any thermal component is too marginal to be detected; hence, XB082 is unlikely to be in the high accretion rate state for neutron star LMXBs."
+ We oxeseut the spectra and best fit two componcut nodel iu Fig. 3.., We present the spectrum and best fit two component model in Fig. \ref{4bhspec}.
+ We conclude that NBUS2 exhibited nou-thermal enission approximated by a power law., We conclude that XB082 exhibited non-thermal emission approximated by a power law.
+ However. active galactic unclei (ACN) also exhibit similar Cluission spectra. hence we calculated the probability of a coincident ACN from the 2.10 keV huuinositv function provided bv ?..," However, active galactic nuclei (AGN) also exhibit similar emission spectra, hence we calculated the probability of a coincident AGN from the 2–10 keV luminosity function provided by \citet{moretti03}."
+" The observed 210 keV flux of NDUS2 derived from the best fit power law model was 2.041007 erg en? hk the probability of an ACN this bright existiue within 1"" of one of the Bs GCs is +2 «10. ©.", The observed 2–10 keV flux of XB082 derived from the best fit power law model was $\times 10^{-12}$ erg $^{-2}$ $^{-1}$; the probability of an AGN this bright existing within $''$ of one of the 428 GCs is $\sim$ $\times$ $^{-6}$.
+ We therefore conclude that NDOS2 is a candidate black hole LMXND., We therefore conclude that XB082 is a candidate black hole LMXB.
+ If we assume P=1.l. then the black hole would require a mass ~30 ML. for a 1550 keV Iuuinositv <0.1 Lepp. consistent with the low state (following?)..," If we assume $\Gamma$ =1.4, then the black hole would require a mass $\sim$ 30 $_{\odot}$ for a 15–50 keV luminosity $\la$ 0.1 $L_{\rm EDD}$, consistent with the low state \citep[following][]{tang10}."
+ This is rather more massive than the Galactic stellar mass black holes (?).. but consistent with the priuary in the dyvuunicallv confirmed black hole | WolfRavet binary ICIO N-L (07).," This is rather more massive than the Galactic stellar mass black holes \citep{ozel10}, but consistent with the primary in the dynamically confirmed black hole + Wolf-Rayet binary IC10 X-1 \citep{silverman08}."
+ We note that the observed spectrum of NBOs2 (D ~1.2) is more often associated with ligh mass A-rav binaries (ΠΑΛΙΑΟΚ) with Be donors (seec.g.T.forarecent review): e.g. 7. examined NAIANewton observations of 11. Be NBs in the Suiall Magellanic Cloud. aud found the photon iudex distribution to be strongly peaked at D — 1.00. with a standard devition of 0.16.," We note that the observed spectrum of XB082 $\Gamma$ $\sim$ 1.2) is more often associated with high mass X-ray binaries (HMXBs) with Be donors \citep[see e.g.][ for a recent review]{reig11}; e.g. \citet{haberl04} examined XMM-Newton observations of 11 Be XBs in the Small Magellanic Cloud, and found the photon index distribution to be strongly peaked at $\Gamma$ = 1.00, with a standard devition of 0.16."
+ Therefore NDUS2 could be a IIMXD superposed on the elobular cluster., Therefore XB082 could be a HMXB superposed on the globular cluster.
+ However. we note that ποσα Be NBs tend to be transient. N-rav sources. as they have lone. eccentric orbits and oulv accrete for a short time when the neutron star is closest to the donor: fiuthermore. this cussion spectrum is ecucrally confined to the range ~107+ 107 eve 1 (?.andreferenceswithin)...," However, we note that known Be XBs tend to be transient X-ray sources, as they have long, eccentric orbits and only accrete for a short time when the neutron star is closest to the donor; furthermore, this emission spectrum is generally confined to the range $\sim$ $^{34}$ $^{37}$ erg $^{-1}$ \citep[][and references within]{reig11}."
+" Persisteuthy bright Be NBs exist. but tend to have N-ray Imniuosities < 10%"" ere +."," Persistently bright Be XBs exist, but tend to have X-ray luminosities $\la$ $^{35}$ erg $^{-1}$."
+ Tf NBOs2 is a neutron star IIMXD. it is unlike any know thus far.," If XB082 is a neutron star HMXB, it is unlike any know thus far."
+" NBI53 was observed in 75 ACTS and 15 URC observations. and appears to be persistently bright: its 0.3.10 keV huninosityv varied by a factor ~3up to ~2.1< 10°"" ere st."," XB153 was observed in 75 ACIS and 45 HRC observations, and appears to be persistently bright; its 0.3–10 keV luminosity varied by a factor $\sim$ 3 up to $\sim$ $\times 10^{38}$ erg $^{-1}$."
+ It was also observed in the 2002 NNMM-Newton observation with 60 ks eood time., It was also observed in the 2002 XMM-Newton observation with 60 ks good time.
+ The best fit power kuv to the NATALNewton pu spectrum vielded Αι = 8.541.0< 1029 atom ?. aud P= 1.6240.01: 42 /dof = 129/tlLaud Los.i10 = LOTEO.OL SIO ore ," The best fit power law to the XMM-Newton pn spectrum yielded $N_{\rm H}$ = $\pm$ $\times 10^{20}$ atom $^{-2}$, and $\Gamma$ = $\pm$ 0.04; $\chi^2$ /dof = 429/444 and $L_{0.3-10}$ = $\pm$ $\times 10^{38}$ erg $^{-1}$."
+Fitting a cluission mocel vielded a hot. optically thin corona with uncoustraimed kT. 060 keV aud 7r —1.," Fitting a emission model yielded a hot, optically thin corona with unconstrained $T_{\rm e}$ $\sim$ 60 keV and $\tau$ $\sim$ 1."
+ The disk blackbody: model was rejected. with v /dof = 1067/LL.," The disk blackbody model was rejected, with $\chi^2$ /dof = 1067/445."
+ Again. Nyy and E for the two component model were consistent with the values for the single power law. and NSPEC was unable to estimate the uncertainties for the blackbody parameters.," Again, $N_{\rm H}$ and $\Gamma$ for the two component model were consistent with the values for the single power law, and XSPEC was unable to estimate the uncertainties for the blackbody parameters."
+ We present the spectrum and two component fit in Fie 3.., We present the spectrum and two component fit in Fig \ref{4bhspec}.
+ NBI53 also appears to have been m a non-thermal state., XB153 also appears to have been in a non-thermal state.
+" The observed 210 keV fux for the best fit power law modelwas «10.1 ere ein ? thence. the probability of finding an ACN of this brightnesso within 1"" of αν of tlhe 128 CC's is 1.0410 7. The ↽∙∶≻∖↽"," The observed 2–10 keV flux for the best fit power law model was $\times10^{-13}$ erg cm $^{-2}$ $^{-1}$; hence, the probability of finding an AGN of this brightness within $''$ of any of the 428 GCs is $\times 10^{-5}$ ."
+⊽∙↴∐∖∏⊔↕⋯↴∖↴↕↖⇁⋟∪↥⋅∐∖⊸⊽⋀⋀≓⋀⊽↸∖↖↖↽∪∐I tvfor the NMM-Newt observation was a factor ~2 lower than the peak., The luminosity for the XMM-Newton observation was a factor $\sim$ 2 lower than the peak.
+ XD 163 was observed in 10. ACTS aud 29 TRC observations., XB 163 was observed in 10 ACIS and 29 HRC observations.
+ NB163 exhibited at least 5 outbursts over the ~ lo00 day viewing period. aud was brightest during ACTS observation slsl (2007 February 11) the exposure was 5 ks.," XB163 exhibited at least 5 outbursts over the $\sim$ 4000 day viewing period, and was brightest during ACIS observation 8184 (2007 February 14); the exposure was 5 ks."
+ The 0.3 keV huuinositv of NBIG3 varied by a factor ~650. lence it cannot be an ACN.," The 0.3--10 keV luminosity of XB163 varied by a factor $\sim$ 650, hence it cannot be an AGN."
+ NDBIG63 was, XB163 was
+profile. or it may be because the universalitv of the density profile (Navarroctal.1996:Mooreetal.1998) really fiuds its origin i accretion listorv (Conzález-Casadoetal.,"profile, or it may be because the universality of the density profile \citep{nfw,moore} really finds its origin in accretion history \citep{gsmh}."
+2007).. Let us finally discuss what our newly fouud attractor may do for the mass-velocity anisotropy degeneracy., Let us finally discuss what our newly found attractor may do for the mass-velocity anisotropy degeneracy.
+ When we observe the stellar ttics in a dwarf galaxy we can observe the stellar density aud the stellar dispersion., When we observe the stellar tics in a dwarf galaxy we can observe the stellar density and the stellar dispersion.
+ Then. the jeans equation. eq. (3)).," Then, the jeans equation, eq. \ref{eq:jeans}) ),"
+ tells us that for amy assumed. velocity anisotropy profile for the stars. oC). we can solve for the total eravitating mass.," tells us that for any assumed velocity anisotropy profile for the stars, $\beta(r)$, we can solve for the total gravitating mass."
+ However. if woe had assumed a different (7) then we would rave found a differcut total mass profile (Strigari 2010).," However, if we had assumed a different $\beta(r)$ then we would have found a different total mass profile \citep{strigari}."
+. If we for instance consider a Ieruquist density xofile with a > profile in agreement with uuucrical simulations aud observatious (ITansen&Piffaretti 2010).. hen the reconstructed nass is overestinated bv up to 1054. if the analvsis is made uncer the siuplifug assuniptiou j—0.," If we for instance consider a Hernquist density profile with a $\beta$ profile in agreement with numerical simulations and observations \citep{hansenpiff2007,host2009,wojtak2010}, then the reconstructed mass is overestimated by up to $40\%$, if the analysis is made under the simplifying assumption $\beta=0$."
+ Also the derived inner density slope (from the tota lass) ds systematically found to be more shallow than the true slope is by up to 1056.," Also the derived inner density slope (from the total mass) is systematically found to be more shallow than the true slope is, by up to $10\%$."
+ This means that if the true density slope is —1. then we will mcasure around 0.95. if we assumed ij= Oin the analysis.," This means that if the true density slope is $-1$, then we will measure around $-0.95$, if we assumed $\beta =0$ in the analysis."
+ This max uo longer have to be the case., This may no longer have to be the case.
+ If our attractor solutions also applies to stellar svstenis ii a dwarf ealaxy. or to the dynamics of the ealaxies in a galaxy cluster. then we have a unique connection between the 3 quantities. 5.& and ," If our attractor solutions also applies to stellar systems in a dwarf galaxy, or to the dynamics of the galaxies in a galaxy cluster, then we have a unique connection between the 3 quantities, $\gamma, \kappa$ and $\beta$."
+Therefore. if we have measured (accurately) the stellar density aud dispersion profiles. then we do in principle kuow exactly what 2(r£0) looks like. aud we can then deduce the unique total mass profile.," Therefore, if we have measured (accurately) the stellar density and dispersion profiles, then we do in principle know exactly what $\beta(r)$ looks like, and we can then deduce the unique total mass profile."
+ We have identified an attractor solution for dark matter structures., We have identified an attractor solution for dark matter structures.
+ This implies that any dark matter structure which is repeatedly perturbed (e.g. through violent relaxtion durius mergi) and then allowed to relax (phase imix). will flow towards this l-cdimensional curve iu the 3-diuensional space spanned by the 2 radial derivatives of the deusity aud velocity dispersion. and the velocity anisotropy.," This implies that any dark matter structure which is repeatedly perturbed (e.g. through violent relaxtion during merging) and then allowed to relax (phase mix), will flow towards this 1-dimensional curve in the 3-dimensional space spanned by the 2 radial derivatives of the density and velocity dispersion, and the velocity anisotropy."
+ This finding provides strong support for the idea that the universalities found in cosmological dark matter structures are a property of eravity. and not simply a result of similar accretion and morecr histories of different structures.," This finding provides strong support for the idea that the universalities found in cosmological dark matter structures are a property of gravity, and not simply a result of similar accretion and merger histories of different structures."
+" This attractor solution effectively removes one deeree of freedom frou he Jeans equation. giving hope that we will eveutually be able to solve the Jeans equations analytically, and thereby truely understaud the origin of the universal profiles."," This attractor solution effectively removes one degree of freedom from the Jeans equation, giving hope that we will eventually be able to solve the Jeans equations analytically, and thereby truely understand the origin of the universal profiles."
+ It is a pleasure to thaws Jens ITjorth for discussions., It is a pleasure to thank Jens Hjorth for discussions.
+ The simulations were performed on the facilities provided by the Danish Center for Scientific Computing., The simulations were performed on the facilities provided by the Danish Center for Scientific Computing.
+ The Dark Cosinology Centre is funded by the Danish National Research Foundation., The Dark Cosmology Centre is funded by the Danish National Research Foundation.
+Moreover our results (e.g. temperatures) for C are consistent with those for C3 (?)..,Moreover our results (e.g. temperatures) for $C_2$ are consistent with those for $C_3$ \citep{Adamkovics2003}.
+ The profile of the narrow DIB at apparently depends on the rotational temperature estimated from the dicarbon molecule. being broader for higher temperatures. which is characteristic of both C» and Cs.," The profile of the narrow DIB at apparently depends on the rotational temperature estimated from the dicarbon molecule, being broader for higher temperatures, which is characteristic of both $C_2$ and $C_3$ ."
+ Cs was analysed by ? and ?.., $C_3$ was analysed by \citet{Adamkovics2003} and \citet{Oka2003}.
+ They found that the column density and rotational temperature derived from Cs 15 well correlated with the same parameters of C»., They found that the column density and rotational temperature derived from $C_3$ is well correlated with the same parameters of $C_2$.
+ It is interesting and important to analyse the simplest multicarbon chains (like, It is interesting and important to analyse the simplest multicarbon chains (like
+(10)) of states into the linear combinations (11)) which are still vanishing.,\ref{eq:linearrel}) ) of states into the linear combinations \ref{eq:linearrel2}) ) which are still vanishing.
+" Also. notice that ©, is self-adjoint. again because of the reality of 2? and it commutes with itself since The 2-netfunction W(s.s). defined in (2)). can be written now as More in general. we can celine so that Now. consider the free linear space A formed. by the (formal) linear combinations of spin networks. with complex coefficients There is a natural product defined on A bv ο5'=5Us’. and a natural star operation defined bv s*=s (Ilere we refer to spin networks labeled bv SU(2) representations ancl each representation of 90(2) is conjugate to itself."," Also, notice that $\hat\phi_{s}$ is self-adjoint, again because of the reality of $P$ and it commutes with itself since The 2-netfunction $W(s,s')$, defined in \ref{eq:W2}) ), can be written now as More in general, we can define so that Now, consider the free linear space $\cal A$ formed by the (formal) linear combinations of spin networks, with complex coefficients There is a natural product defined on $\cal A$ by $s\cdot s'=s\cup
+s'$, and a natural star operation defined by $s^{*}=s$ (Here we refer to spin networks labeled by $SU(2)$ representations and each representation of $SU(2)$ is conjugate to itself."
+ When spin networks are labeled by representations of groups which are not sell-conjugate the star operation should replaces representations will dual representations.), When spin networks are labeled by representations of groups which are not self-conjugate the star operation should replaces representations with dual representations.)
+ We define the norm |L1]|= sup;|c;|., We define the norm $||A||=sup_{s}|c_{s}|$ .
+ We obtain in this wav a C algebra structure on A., We obtain in this way a $C^{*}$ algebra structure on $\cal A$.
+ The quantity W(s). defined in (17)). defines a linear functional on A.A straightforward. caleulation shows that the finelional is positive We can thus apply the Gelland-Naimark-Segal construction to the C algebra A and the positive linear functional WW. obtaining a Lilhert space H. a," The quantity $W(s)$, defined in \ref{eq:W}) ), defines a linear functional on $\cal A$.A straightforward calculation shows that the functional is positive We can thus apply the Gelfand-Naimark-Segal construction to the $C^*$ algebra $\cal A$ and the positive linear functional$W$ , obtaining a Hilbert space $\cal H$ , a"
+The second column in Table 1 shows the small effect of this smoothing on the fit parameters.,The second column in Table \ref{tab:1t} shows the small effect of this smoothing on the fit parameters.
+" The fit quality is slightly poorer after accounting for smoothing, suggesting that the emission lines may be emitted on a smaller scale than the broader X-ray emission(Section 3))."," The fit quality is slightly poorer after accounting for smoothing, suggesting that the emission lines may be emitted on a smaller scale than the broader X-ray emission(Section \ref{sect:profiles}) )."
+ The goodness command gave 54.6 per cent of realisations of the best fitting model with a better fit than the data., The goodness command gave 54.6 per cent of realisations of the best fitting model with a better fit than the data.
+" Since the lines are remarkably narrow, and the source extent is small, we can limit the velocity broadening."," Since the lines are remarkably narrow, and the source extent is small, we can limit the velocity broadening."
+ We fitted the spectra in Section 2.1 with the model in assuming that the spectra were broadened by the thermal motion of the ions in the gas (version 12.5.0ah of fixed an error in the line widths) and line-of-sight velocity broadening added in quadrature., We fitted the spectra in Section \ref{sect:1t} with the model in assuming that the spectra were broadened by the thermal motion of the ions in the gas (version 12.5.0ah of fixed an error in the line widths) and line-of-sight velocity broadening added in quadrature.
+ The spectral lines were used in the model., The spectral lines were used in the model.
+ We show the change in fit statistic as a function of velocity in Fig. 3.., We show the change in fit statistic as a function of velocity in Fig. \ref{fig:broadening}.
+ If we conservatively treat the cluster as a point source and do not use the model we obtain a 90 per cent upper limit of 274kms~!., If we conservatively treat the cluster as a point source and do not use the model we obtain a 90 per cent upper limit of $274\kmps$.
+ This limit is improved to 182kms~! with the addition of the broadening model., This limit is improved to $182\kmps$ with the addition of the broadening model.
+" Examining just the Fe-L spectral region between 9.6 and (rest), we obtain a limit of 214kms-!, and for the Ist order spectrum between 13.6 and22.4A, containing the strong O line, a limit of 3830kms~!."," Examining just the Fe-L spectral region between 9.6 and (rest), we obtain a limit of $214\kmps$, and for the 1st order spectrum between 13.6 and, containing the strong O line, a limit of $380\kmps$."
+ We confirmed that such constraints are readily achievable with simulated spectra., We confirmed that such constraints are readily achievable with simulated spectra.
+ We can limit how much gas can be cooling within temperature bins., We can limit how much gas can be cooling within temperature bins.
+" We constructed a model using six cooling flow model components (constructed using models) with temperature ranges of 105.6>2.814070.35—0.0808 keV. We fixed the components to have the same metallicities, allowing individual elements to vary as in Section 2.1.."," We constructed a model using six cooling flow model components (constructed using models) with temperature ranges of $10 \rightarrow 5.6 \rightarrow 2.8 \rightarrow 1.4
+\rightarrow 0.7 \rightarrow 0.35 \rightarrow 0.0808$ keV. We fixed the components to have the same metallicities, allowing individual elements to vary as in Section \ref{sect:1t}."
+" The maximum cooling rates in Moyr~!, assuming isobaric cooling, were free parameters in the fit."," The maximum cooling rates in $\Msunpyr$, assuming isobaric cooling, were free parameters in the fit."
+" The results are shown in Fig. 4,,"," The results are shown in Fig. \ref{fig:coolrate},"
+ with and without the smoothing component., with and without the smoothing component.
+" For comparison, we also fitted a model made up of an isothermal component plus a model cooling from its temperature to zero."," For comparison, we also fitted a model made up of an isothermal component plus a model cooling from its temperature to zero."
+ The best fitting temperature of the isothermal component was 3.85 keV (close to the single temperature fits)., The best fitting temperature of the isothermal component was 3.85 keV (close to the single temperature fits).
+ We obtain cooling rates from this temperature to zero of 70736Moyr!., We obtain cooling rates from this temperature to zero of $70^{+46}_{-56}\Msunpyr$.
+ This rate is shown by the shaded bar in Fig. 4.., This rate is shown by the shaded bar in Fig. \ref{fig:coolrate}.
+ The 90 per cent upper limit is 140Mayr! ., The 90 per cent upper limit is $140\Msunpyr$ .
+ A Markov Chain Monte Carlo analysis produces, A Markov Chain Monte Carlo analysis produces
+are systematically larger than the corresponding distances from the CLLA.,are systematically larger than the corresponding distances from the CLLA.
+ This difference can be explained bv some unaccounted components. the presence of which may lead to an underestimation of photometric distances.," This difference can be explained by some unaccounted components, the presence of which may lead to an underestimation of photometric distances."
+ Another reason lor this discrepancy is the photometric distance calibration adopted in the CLLA., Another reason for this discrepancy is the photometric distance calibration adopted in the CLLA.
+ Most of the stars in our sample were examined for common proper molion components (Allenetal.2000:ZapateroOsorio&Martin2004).," Most of the stars in our sample were examined for common proper motion components \citep{allen,zapatero}."
+". The speckle interferometric observations of 223 sample stars were carried out in 20062007 on the 6 m BTA telescope (Itastegaevetal.2007.2008) which dilfraction-limitec resolution is 0.023"" for A=550 nm and 0.033"" lor A=800 mnm."," The speckle interferometric observations of 223 sample stars were carried out in 2006–2007 on the 6 m BTA telescope \citep{rastegaev_2007,rastegaev_2008} which diffraction-limited resolution is $0.023 ''$ for $\lambda=550$ nm and $0.033 ''$ for $\lambda=800$ nm."
+ Most of the observations were carried out using the svstem (Maksimovetal.2009) based on a 512x512 EAICCD (a CCD [eaturing on-chip multiplication gain) with high «quantum efficiency. ancl linearity.," Most of the observations were carried out using the system \citep{maksimov}
+ based on a $\times$ 512 EMCCD (a CCD featuring on-chip multiplication gain) with high quantum efficiency and linearity."
+" This svstem allowed us to detect objects with magnitude differences between the components ol up to Am=5"".", This system allowed us to detect objects with magnitude differences between the components of up to $\triangle m = 5^{m}$.
+ Takine into account the limiting stellar magnitude of our sample (my« 127).detected secondary component can be as faint as 177.," Taking into account the limiting stellar magnitude of our sample $\mathrm{m_V}<12^{m}$ ),detected secondary component can be as faint as $17^{m}$."
+" The 4.4” field of view of our system allows detection of secondary components at a separation of3"" from the primary star.", The $4.4''$ field of view of our system allows detection of secondary components at a separation of$3''$ from the primary star.
+ The speckle interferograms were recorded using five filters: 545/30. 550/20. /600/40. 800/100/ and 800/110 nm (the first number indicates the central wavelength of the filler. the second the hall-width of the filters bandwidth) with the exposures of 5 to 20 ms.," The speckle interferograms were recorded using five filters: $545/30$, $550/20$, $600/40$, $800/100$ and $800/110$ nm (the first number indicates the central wavelength of the filter, the second — the half-width of the filter's bandwidth) with the exposures of 5 to 20 ms."
+ For each object. we accumulated Irom 500 to 2000 short exposure images depending on weather conditions.," For each object, we accumulated from 500 to $2\ 000$ short exposure images depending on weather conditions."
+" The observations were made with an average seeing of 1.5"".", The observations were made with an average seeing of $1.5''$.
+" The accuracy of our speckle interferogram processing method (Dalegaοἱal.2002) may be as οσους as 0.02"". 0.001"". and 0.1* for the component maenitude dilference. angular separation and position angle respectively."," The accuracy of our speckle interferogram processing method \citep{balega_2002} may be as good as $0.02^m$ $0.001''$ , and $0.1^{\circ}$ for the component magnitude difference, angular separation and position angle respectively."
+baselines.,baselines.
+ Alternatively it may indicate that these largest scales are simply not present in the galaxies under consideration., Alternatively it may indicate that these largest scales are simply not present in the galaxies under consideration.
+ It is clear though that at large scales there is good agreement., It is clear though that at large scales there is good agreement.
+" As the spatial scale probed approaches the size of the natural-weighted beam however, the power in the lintegrated moment maps begins to decrease compared to the mmaps."," As the spatial scale probed approaches the size of the natural-weighted beam however, the power in the integrated moment maps begins to decrease compared to the maps."
+ At about ~1.5 times the natural-weighted beam size we start to see a significant deviation., At about $\sim1.5$ times the natural-weighted beam size we start to see a significant deviation.
+ Also shown in Figure 19 isa power law with slope —3 which is a reasonable description for the power spectrum at intermediate scales., Also shown in Figure \ref{fig:powerspec} is a power law with slope $-3$ which is a reasonable description for the power spectrum at intermediate scales.
+ It is consistent with values found in other galaxies(?)., It is consistent with values found in other galaxies.
+". We will not here to relate the power law to tthe turbulence or the attemptenergy input of the ISM, but we slopedraw attention to the fact that as the small-scale power in the mmaps starts to fall away, the ppower spectrum continues to follow this power-law behavior."," We will not attempt here to relate the power law slope to the turbulence or the energy input of the ISM, but we draw attention to the fact that as the small-scale power in the maps starts to fall away, the power spectrum continues to follow this power-law behavior."
+ This indicates that the probe real small-scale structure more efficiently than the classical mmapsmmaps., This indicates that the maps probe real small-scale structure more efficiently than the classical maps.
+ Note that the power in the mmaps only starts to drop away at scale sizes of half a natural-weighted beam., Note that the power in the maps only starts to drop away at scale sizes of half a natural-weighted beam.
+ Scales probed in classical, Scales probed in classical
+Since the late 1950s. it has been realized that neutrou-star iuterior may consist of a πο of quantum fiuids (see Shapiro Toukolsky 1983 for a review).,"Since the late 1950's, it has been realized that neutron-star interior may consist of a number of quantum fluids (see Shapiro Teukolsky 1983 for a review)."
+ Curreuthe it is thought that both neutron superfiuid and proton superconductor are likely to coexist in the neutron-star cores (see. 6.8.. Link 2007 Or a discussion).," Currently, it is thought that both neutron superfluid and proton superconductor are likely to coexist in the neutron-star cores (see, e.g., Link 2007 for a discussion)."
+ Several researchers have arec that if the proton supercouductivityv were| of the type IL then t1c superconductor fluxtuvos would couple strongv to the notron superfluid vortices.," Several researchers have argued that if the proton superconductivity were of the type II, then the superconductor fluxtubes would couple strongly to the neutron superfluid vortices."
+ This line of rTOasonding is based ou the fact that nuclear forces contain velocity-depeuceut terms. which results i the οιαπλο! of protons in he jeutron supercurrent (Alpar. Langer. Sauls. 1958).," This line of reasoning is based on the fact that nuclear forces contain velocity-dependent terms, which results in the entrainment of protons in the neutron supercurrent (Alpar, Langer, Sauls, 1984)."
+ Therefore. the vortices are sheathe« by charged curreits cutrained in the suοΠα flow. aud are strongv maenetized.," Therefore, the vortices are sheathed by charged currents entrained in the superfluid flow, and are strongly magnetized."
+ Magnetic fiuxtubes interact stroely with t1C Wagjetized vortices. snilar to the wav in wuch ειο fluxtubes interact between each other (Rileruniui. Zhu. Chen 1998. and references therein).," Magnetic fluxtubes interact strongly with the magnetized vortices, similar to the way in which the fluxtubes interact between each other (Ruderman, Zhu, Chen 1998, and references therein)."
+ As a result f this couphue. the vortices ect strongly piune« to +ie proton-clectron plasma in the core.," As a result of this coupling, the vortices get strongly pinned to the proton-electron plasma in the core."
+ Such pininue would have nu»ortant iuplications for he neutron-star plienomenology., Such pinning would have important implications for the neutron-star phenomenology.
+ Ruclerman. Zhu. Chen (1998) have argued. that t1ο vortex-pinning iu the core lay be responsible for he observed elitejes in the pulsar rotation rates.," Ruderman, Zhu, Chen (1998) have argued that the vortex-pinning in the core may be responsible for the observed glitches in the pulsar rotation rates."
+ Link (2003) has considerec the effect of the vortex-fluxtube interaction ou the cdyvuamics of the precessineo pulsar PSR 1828-11 (observed by Stairs. Lyuc. Shelu 2000).," Link (2003) has considered the effect of the vortex-fluxtube interaction on the dynamics of the precessing pulsar PSR 1828-11 (observed by Stairs, Lyne, Shemar 2000)."
+ Building ou the theoretica WNOYs by Shaham (1977) and Sedrakian. Wasscermial.. Cordes (1999). le has concluded that the interaction. if oesenut. would ultimaelv lead to the fast precession.," Building on the theoretical work by Shaham (1977) and Sedrakian, Wasserman, Cordes (1999), he has concluded that the interaction, if present, would ultimately lead to the fast precession."
+ Since PSR JBN28-11 is precessing slowly and persisteuthy. Links (2003) las argued that the core vortex pinning is excluded by tιο observations and hence that either the proton sitperconductor migit be of tvpe L or that both proton aud neun condensates do uot coexist iuside that pulsar.," Since PSR 1828-11 is precessing slowly and persistently, Link (2003) has argued that the core vortex pinning is excluded by the observations and hence that either the proton superconductor might be of type I, or that both proton and neutron condensates do not coexist inside that pulsar."
+ While Links argmeut is suggestive. we believe if is prenature ο rule out strong vortex piunius iu the cores of a] neutron stars.," While Link's argument is suggestive, we believe it is premature to rule out strong vortex pinning in the cores of all neutron stars."
+ Iu this paper we consier hydromaguctic waves in the case when the neutron vortices are stronglv piined to the protou-clectrou plasma iu he core., In this paper we consider hydromagnetic waves in the case when the neutron vortices are strongly pinned to the proton-electron plasma in the core.
+ We have 2 nun astroplivsical motivations or studying this problem., We have 2 main astrophysical motivations for studying this problem.
+ The first one is due to the fairly receut observations, The first one is due to the fairly recent observations
+"Even galaxies with little star formation activity continue to evolve, as evidenced by the substantial increase of their cosmic stellar mass density over the past 7 billion years","Even galaxies with little star formation activity continue to evolve, as evidenced by the substantial increase of their cosmic stellar mass density over the past 7 billion years"
+D aud V stellar maeuitudes.,$B$ and $V$ stellar magnitudes.
+" Using these parameters, we calculate the effective temperature and Imuuinositv via the Zaustra method (I&aloey 1983). ("," Using these parameters, we calculate the effective temperature and luminosity via the Zanstra method (Kaler 1983). ("
+2) Dx locating the central star iu the logTig logL plane. we derive its mass (Mog) from comparison with a set of evolutionary tracks (Stanghellini Renzini 1993). (,"2) By locating the central star in the $\log T_{\rm eff}$ $\log L$ plane, we derive its mass $M_{\rm CS}$ ) from comparison with a set of evolutionary tracks (Stanghellini Renzini 1993). ("
+3) Usine the initial massfinal mass relation. we colpute the progenitor mass. that is the stellar mass on the main sequence (A/S).,"3) Using the initial mass–final mass relation, we compute the progenitor mass, that is the stellar mass on the main sequence $M_{\rm MS}$ )."
+ The stellar properties adopted for the PNe detected iu PCO and the derived. values of the progenitor mass are eiven in Table 3., The stellar properties adopted for the PNe detected in $^{13}$ CO and the derived values of the progenitor mass are given in Table 3.
+ Details ou the individual objects are eiven iu the Appoeudix., Details on the individual objects are given in the Appendix.
+ Let us exanune the uncertainty involved iu the final mass calculations., Let us examine the uncertainty involved in the final mass calculations.
+ Estimates of the stellar temperature and Iuninositv eiven iu Table 3 are affected by errors in magnitudes. fixes. diameters aud extinctions.," Estimates of the stellar temperature and luminosity given in Table 3 are affected by errors in magnitudes, fluxes, diameters and extinctions."
+ However. these quautitics are usually determined with οσους accuraev (~520%). so that the wnecrtainty in the erived mass of the central stars does not exceed ~15 or (0.02AZ.," However, these quantities are usually determined with good accuracy $\sim 5-20\%$ ), so that the uncertainty in the derived mass of the central stars does not exceed $\sim 15$, or $\sim 0.02$."
+.. The values given iu the table do ixt oeiclude errors on the distances to the PN. which can e oeitrinsically hiehl (up to 505€)) but ave difficult to estimate on an individual basis.," The values given in the table do not include errors on the distances to the PN, which can be intrinsically high (up to ) but are difficult to estimate on an individual basis."
+ To inter the main sequence masses. we have used the oeitial massfinal mass relation given by Uervig (1996).," To infer the main sequence masses, we have used the initial mass–final mass relation given by Hervig (1996)."
+ This relation differs from that of Weideiiaun (LOST) vdopted in GSTP., This relation differs from that of Weidemann (1987) adopted in GSTP.
+ We preferred Heorvigs prescription since it is based ou reliable observations of cluster white cdawarts. although the formal errors on the final mass are still substantial. and can amount to about 01 ALL.," We preferred Hervig's prescription since it is based on reliable observations of cluster white dwarfs, although the formal errors on the final mass are still substantial, and can amount to about 0.1 $M_\odot$."
+ Since we derive initial masses from final masses. the errors ou the Ormer quantity can be even larger.," Since we derive initial masses from final masses, the errors on the former quantity can be even larger."
+ Quantitatively. we assign a formal error to the main sequence mass of AAAsτσ1.5 stor low values of the initial mass GM<2 AL.) anda simaller error (AAAs&0.75 i) for higher masses.," Quantitatively, we assign a formal error to the main sequence mass of $\Delta M_{\rm MS}\simeq 1.5$ for low values of the initial mass $M_{\rm MS}<2$ ), and a smaller error $\Delta M_{\rm MS}\simeq
+0.75$ ) for higher masses."
+ This difference is due to the change of the slope of the initial mass - final mass relation at about 2M: sinaller masses are more sensitive to the adopted relation. and the uncertainty is correspondingly larger.," This difference is due to the change of the slope of the initial mass - final mass relation at about 2: smaller masses are more sensitive to the adopted relation, and the uncertainty is correspondingly larger."
+ Tow do we interpret our results on the irafios in the framework of stellar uucleosvuthesis?, How do we interpret our results on the ratios in the framework of stellar nucleosynthesis?
+ To help answering this question. we combine the information provided by the observed lisotopic ratios with the mass estimates of the progenitors of the PNe. and with the predictions of some representative stellar uucleosvuthesis models.," To help answering this question, we combine the information provided by the observed isotopic ratios with the mass estimates of the progenitors of the PNe, and with the predictions of some representative stellar nucleosynthesis models."
+ Since the formation of a PN takes place at the eund of the AGB phase. the significant comparison is between the observed abundancees and those predicted for the stellar ejecta at the AGB tip.," Since the formation of a PN takes place at the end of the AGB phase, the significant comparison is between the observed abundances and those predicted for the stellar ejecta at the AGB tip."
+ Uufortunately. no," Unfortunately, no"
+p and d-sates of Bain.,$p$ and $d$ -sates of Ba.
+ Then. we obtain the transition probabilities in the tensorial irreducible basis.," Then, we obtain the transition probabilities in the tensorial irreducible basis."
+ Afterwards. these propabilities are integrated over impact parameters and Maxwellian distribution of relative velocities to obtain the depolarization and the transfer rates.," Afterwards, these propabilities are integrated over impact parameters and Maxwellian distribution of relative velocities to obtain the depolarization and the transfer rates."
+ We perform calculations varying the temperature to obtain the best analytical fit to the collistonal rates., We perform calculations varying the temperature to obtain the best analytical fit to the collisional rates.
+The spectra of our new program stars have been obtained in (he same fashion as (hose studied by C2003. using the Center for Astrophysics Digital Speedometers (Latham 1935. 1992). primarily with the 1.5-m Wveth reflector at (he Oak Ridge Observatory in Harvarel.Massachusetts’. as well as the 1.5-m Tillinghast reflector and the MMT. instruments atop Alt. Hopkins in Arizona.,"The spectra of our new program stars have been obtained in the same fashion as those studied by C2003, using the Center for Astrophysics Digital Speedometers (Latham 1985, 1992), primarily with the 1.5-m Wyeth reflector at the Oak Ridge Observatory in Harvard, as well as the 1.5-m Tillinghast reflector and the MMT instruments atop Mt. Hopkins in Arizona."
+ The Tillinghast reflector was especially important for the stars south of —20° aand north of +62° iin declination., The Tillinghast reflector was especially important for the stars south of $-20$ and north of $+62$ in declination.
+ Also as before. the wavelength coverage is 45A.. centered near 5187À.. with a resolution of 8.5+.," Also as before, the wavelength coverage is 45, centered near 5187, with a resolution of 8.5."
+. The signal-to-noise ratio varied [rom 10 to 50 per resolution element. will a (vpical value of about 15.," The signal-to-noise ratio varied from 10 to 50 per resolution element, with a typical value of about 15."
+ C2003 described in detail the measurement. of the radial velocities., C2003 described in detail the measurement of the radial velocities.
+ A grid of model almospheres. defined by Zar. log g. and. |Fe/II] values. was computed.," A grid of model atmospheres, defined by $T_{\rm eff}$, log $g$, and [Fe/H] values, was computed."
+ The grid spacing in temperature was 250 Ix. 0.5 dex in log g. and 0.5 dex in metallicity for [Fe/H] <—1.0.," The grid spacing in temperature was 250 K, 0.5 dex in log $g$, and 0.5 dex in metallicity for [Fe/H] $\leq\ -1.0$."
+ functions and elemental abundances in whieh all the a elements (O. Ne. Meg. SiS. Ca. and Ti) were enhanced by 0.4 dex relative to the solar abundances.," functions and elemental abundances in which all the $\alpha$ "" elements (O, Ne, Mg, Si, S, Ca, and Ti) were enhanced by 0.4 dex relative to the solar abundances."
+ More details may be found in Nórrdstrom et ((1994) and C2003., More details may be found in Nörrdstrom et (1994) and C2003.
+ The program SYNTIIE was used to compute svnthetic spectra in Che waveleng(h range 5146-5229A.., The program SYNTHE was used to compute synthetic spectra in the wavelength range 5146-5229.
+ In the case of the Sun. we obtained an excellent line-by-Iine match between the ]vurucz solar model andsynfhetie spectrum compared with theobserved solar [πι spectrum (Ixurucz et 11984).," In the case of the Sun, we obtained an excellent line-by-line match between the Kurucz solar model and spectrum compared with the solar flux spectrum (Kurucz et 1984)."
+ SYNTIIE computes specilic intensity al 17 different emergent angles across the stellar disk. and integration over the disk. including the effects of stellar rotation. vields (he svntlietic [Inx spectrum.," SYNTHE computes specific intensity at 17 different emergent angles across the stellar disk, and integration over the disk, including the effects of stellar rotation, yields the synthetic flux spectrum."
+ 5YNTIIE enables us to include the effects of instrumental resolution. which we chose to be a Gaussian with a FWILM o£ 8.5f... which is appropriate io the instrumentation we enploved.," SYNTHE enables us to include the effects of instrumental resolution, which we chose to be a Gaussian with a FWHM of 8.5, which is appropriate to the instrumentation we employed."
+ The adopted microturbulent. velocities were 2HL. and macroturbulent velocities were 3|.," The adopted microturbulent velocities were 2, and macroturbulent velocities were 3."
+".. At each combination of temperature. eravity. and metallicity in our grid we computed svnthetic spectra employing a wide range of rotational broadening profiles. with Veo, = O0. 1. 2. 4. 6. 5. 10. 12. 16. 20. 25. 30. 35. 40. 50. 60. το SO. 90. LOO. 110. 120. and 1401."," At each combination of temperature, gravity, and metallicity in our grid we computed synthetic spectra employing a wide range of rotational broadening profiles, with $V_{\rm rot}$ = 0, 1, 2, 4, 6, 8, 10, 12, 16, 20, 25, 30, 35, 40, 50, 60, 70, 80, 90, 100, 110, 120, and 140."
+ Once the stellar parameters. Zr. log g. and. [Fe/II] had been estimated. we relied on ihe model atmosphere grid point closest in (hese variables. paving special attention to the primary variable. temperature.," Once the stellar parameters, $T_{\rm eff}$, log $g$, and [Fe/H] had been estimated, we relied on the model atmosphere grid point closest in these variables, paying special attention to the primary variable, temperature."
+ In the case of more than one close match. we emploved the template that gave the hiehest value for the peak correlation. averaged over all the observed," In the case of more than one close match, we employed the template that gave the highest value for the peak correlation, averaged over all the observed"
+matter distribution.,matter distribution.
+ The results of ? apply only to isolated galaxies. while the cross-correlation approach applies to any parent galaxy.," The results of \citet{chen_etal06} apply only to isolated galaxies, while the cross-correlation approach applies to any parent galaxy."
+ This suggests that the environment of parent galaxies is relatively unimportant., This suggests that the environment of parent galaxies is relatively unimportant.
+ My result is also generally consistent with cluster-sized simulations which include cooling and star formation., My result is also generally consistent with cluster-sized simulations which include cooling and star formation.
+ In such simulations. while the distribution of dark matter subhalos is less concentrated than the dark matter. the stellar components of satellite galaxies more closely follow the dark matter distribution (??)..," In such simulations, while the distribution of dark matter subhalos is less concentrated than the dark matter, the stellar components of satellite galaxies more closely follow the dark matter distribution \citep{nagai_kravtsov05,maccio_etal06}."
+ However. the possible discrepancy in values for f£ and fa. suggest that the spatial distribution of a class of satellite objects may depend on the size of the host halo: galaxy-sized halos. groups and clusters.," However, the possible discrepancy in values for $f$ and $f_{\rm cross}$ suggest that the spatial distribution of a class of satellite objects may depend on the size of the host halo: galaxy-sized halos, groups and clusters."
+ ? suggest that the projected radial distribution of satellite galaxies depends upon satellite color. such that redder satellites have a significantly steeper density profile than bluer satellites.," \citet{chen08} suggest that the projected radial distribution of satellite galaxies depends upon satellite color, such that redder satellites have a significantly steeper density profile than bluer satellites."
+ A cross-correlation analysis with a different selection function for satellites then should show such a color dependence., A cross-correlation analysis with a different selection function for satellites then should show such a color dependence.
+ In order to achieve better constraints on the spatial distribution of satellite galaxies and its environmental and color dependency. larger data sets are required.," In order to achieve better constraints on the spatial distribution of satellite galaxies and its environmental and color dependency, larger data sets are required."
+ One possibility for increasing the volume of data is to apply a similar analysis to the photometric redshift catalog. instead of limiting the sample to objects with spectroscopic redshifts.," One possibility for increasing the volume of data is to apply a similar analysis to the photometric redshift catalog, instead of limiting the sample to objects with spectroscopic redshifts."
+ The SDSS photometric redshift catalog goes far deeper than the spectroscopic catalog and without the spectroscopic catalog’s fiber collision problem., The SDSS photometric redshift catalog goes far deeper than the spectroscopic catalog and without the spectroscopic catalog's fiber collision problem.
+ However. photometric redshifts are less accurate and may complicate the statistical error analysis.," However, photometric redshifts are less accurate and may complicate the statistical error analysis."
+ Given that future large surveys such as the Dark Energy Survey and the Panoramic Survey Telescope Rapid Response System (Pan-STARRS) which will produce significantly more photometric redshifts than currently available. it would be interesting to perform a analysis with photometric redshifts.," Given that future large surveys such as the Dark Energy Survey and the Panoramic Survey Telescope Rapid Response System (Pan-STARRS) which will produce significantly more photometric redshifts than currently available, it would be interesting to perform a cross-correlation analysis with photometric redshifts."
+Internal dust might allect seriously the observed. properties of distant objects. ancl observations have indeed. found evidence of dust in high-z objects.,"Internal dust might affect seriously the observed properties of distant objects, and observations have indeed found evidence of dust in $z$ objects."
+ Several active galaxies with z 2 have been detected at submillimeter wavelengths. suggesting the presence of large amounts of cust (Mausx1077 Mo: see Hughes 1996 for a recent review).," Several active galaxies with $z>$ 2 have been detected at submillimeter wavelengths, suggesting the presence of large amounts of dust $_{dust} \sim 10^{8-9}$ $_{\odot}$; see Hughes 1996 for a recent review)."
+ The reddening of the background quamus and their metal abundances suggest the presence of dust in damped Lye absorption systems (Pettini et al., The reddening of the background quasars and their metal abundances suggest the presence of dust in damped $\alpha$ absorption systems (Pettini et al.
+ 1994: Per Fall 1996)., 1994; Pei Fall 1996).
+ Also. the UV polarization of propertieshigh-z radio galaxies can be explained in term of dust scattering. suggesting a significant amount of dust in their ISM (Cimatti 1996 and references therein).," Also, the UV polarization properties of $z$ radio galaxies can be explained in term of dust scattering, suggesting a significant amount of dust in their ISM (Cimatti 1996 and references therein)."
+ A substantial amount of cust. is also expected in evolutionary models of spheroidal galaxies at high-z (Franceschini et al., A substantial amount of dust is also expected in evolutionary models of spheroidal galaxies at $z$ (Franceschini et al.
+ 1994: Mazzei De Zotti 1996 and references therein)., 1994; Mazzei De Zotti 1996 and references therein).
+ Although dust. is likely to be present in most high-z systems. no information is available neither on its spatial distribution. nor about howrealise dust. extinction. can aleet our view of distant galaxies.," Although dust is likely to be present in most $z$ systems, no information is available neither on its spatial distribution, nor about how dust extinction can affect our view of distant galaxies."
+ In Lact. extinction is usually treated in a simplistic way. neelecting the contribution of scattering. and assuming naive spatial distributions (e.g. uniform foreground. screens or infinite slabs).," In fact, extinction is usually treated in a simplistic way, neglecting the contribution of scattering, and assuming naive spatial distributions (e.g. uniform foreground screens or infinite slabs)."
+ Only recentlv. more realistic models have. been developed (IxvlIafis Baheall 1987: Bruzualet al.," Only recently, more realistic models have been developed (Kylafis Bahcall 1987; Bruzual et al."
+ 1988: Witt et al., 1988; Witt et al.
+ 1992: Byun et al., 1992; Byun et al.
+ 1994: Wise Silva 1996: Bianchi. Ferrara Giovanardi 1996. hereafter BEG).," 1994; Wise Silva 1996; Bianchi, Ferrara Giovanardi 1996, hereafter BFG)."
+ Although these models have shown that cust scattering. plavs an important role by reducing the ellects. of dust. absorption. no systematic studies of the ellects of the extinction on the observed. colours of high-z. galaxies have been pertormed.," Although these models have shown that dust scattering plays an important role by reducing the effects of dust absorption, no systematic studies of the effects of the extinction on the observed colours of $z$ galaxies have been performed."
+ Understanding these elfects is crucial in cosmology and galaxy evolution studies., Understanding these effects is crucial in cosmology and galaxy evolution studies.
+ Vor instance. a relevant problem arises in the age estimates of high-z galaxies.," For instance, a relevant problem arises in the age estimates of $z$ galaxies."
+ When deep spectroscopy. is not available. no information on the stellar continuum and absorption features are obtainable. and the age estipiates are based. solely on the fitting of the broad-band photomoetric Spectral Energy. Distributions (SEDs) with svnthetic stellar population spectra.," When deep spectroscopy is not available, no information on the stellar continuum and absorption features are obtainable, and the age estimates are based solely on the fitting of the broad-band photometric Spectral Energy Distributions (SEDs) with synthetic stellar population spectra."
+ In a simplistic scenario where the extinction is entirely. due to absorption (the screen model). itis well known that the colours ofa reddened voung galaxy," In a simplistic scenario where the extinction is entirely due to absorption (the screen model), it is well known that the colours of a reddened young galaxy"
+"As a measure of bar strength. we use Q,=max(hy;/«Pi>). the maximum of tangential force amplitude relative o the mean axisvmmetric radial force. evaluated at the region of the bar.","As a measure of bar strength, we use $Q_g = \max({F_T/<F_R>})$, the maximum of tangential force amplitude relative to the mean axisymmetric radial force, evaluated at the region of the bar."
+ The mass-to-light ratio (AL/£) is assumed to be constant. and the vertical scale height of the disk (ancl bar) is estimated from the exponential scale length. using a Hubble tvpe dependent mean ratio.," The mass-to-light ratio $M/L$ ) is assumed to be constant, and the vertical scale height of the disk (and bar) is estimated from the exponential scale length, using a Hubble type dependent mean ratio."
+ Also. the dillerent 3D density cüstribution of the bulge is corrected. based on bulge models obtained roni decompositions.," Also, the different 3D density distribution of the bulge is corrected, based on bulge models obtained from decompositions."
+" The elfect of including dark halo force field was also investigated. but its influence on (Q, appeared to »e Insignificant at the bar region (63]. 17])."," The effect of including dark halo force field was also investigated, but its influence on $Q_g$ appeared to be insignificant at the bar region ([63], [17])."
+" These calculations also give a proxy for the bar length. rQy. which is the radius where the maximum tangential force (Q, occurs."," These calculations also give a proxy for the bar length, $rQ_g$, which is the radius where the maximum tangential force $Q_g$ occurs."
+" Dar lengths were estimated also from the phases of the ο Fourier amplitude » assuming that it is maintained nearly constant in the bar region (a correlation between this bar length and. rQ,y was shown by 48]).", Bar lengths were estimated also from the phases of the $A_2$ Fourier amplitude by assuming that it is maintained nearly constant in the bar region (a correlation between this bar length and $rQ_g$ was shown by [48]).
+ The maximum of m-—2 Fourier amplitude. clo. was used as an estimate of the relative brightness of the bar.," The maximum of $m$ =2 Fourier amplitude, $A_2$, was used as an estimate of the relative brightness of the bar."
+ These properties are shown in Figure 5 in Laurikainen et al. , These properties are shown in Figure 5 in Laurikainen et al. [
+"51] as à function of Hubble: bars grow in length and in relative xightness (le) towards the earlv-tvpe galaxies. but for €, the trend is opposite.","51] as a function of Hubble: bars grow in length and in relative brightness $A_2$ ) towards the early-type galaxies, but for $Q_g$ the trend is opposite."
+ Notice that although the bar ellipticity (shown in the same figure) correlates with Quy (91]). it has no systematic correlation with the Hubble tvpe.," Notice that although the bar ellipticity (shown in the same figure) correlates with $Q_g$ ([91]), it has no systematic correlation with the Hubble type."
+ In Laurikainen et el. , In Laurikainen et el. [
+"48 ] the tendency. of weakening bar strengths ((,) towards the early-type galaxies was explained by a dilution elfect due to the more massive bulges in the early-tvpe galaxies (c.g. the average Q, parameter may decrease even if the average se amplitude increases. since the bulge contribution to radial force becomes more important toward earlier tvpes).","48 ] the tendency of weakening bar strengths $Q_g$ ) towards the early-type galaxies was explained by a dilution effect due to the more massive bulges in the early-type galaxies (e.g. the average $Q_g$ parameter may decrease even if the average $A_2$ amplitude increases, since the bulge contribution to radial force becomes more important toward earlier types)."
+" There exist à correlation also between Q, and slo. but for the above reason the correlations are dilferent for the early and late-type galaxies (sce Fig."," There exist a correlation also between $Q_g$ and $A_2$, but for the above reason the correlations are different for the early and late-type galaxies (see Fig."
+ S in. 48])., 8 in [48]).
+ The obtained tendeney for bar lengths was originally shown by Elmegreen & Elmoegreen *Ji, The obtained tendency for bar lengths was originally shown by Elmegreen $\&$ Elmegreen [28].
+ For à sub-sample of 26 barred. galaxies in NIRSOS. the radial sto profiles were fitted by single (SC) ancl double (DC) Gaussian functions by Buta et al. ," For a sub-sample of 26 barred galaxies in NIRS0S, the radial $A_2$ profiles were fitted by single (SG) and double (DG) Gaussian functions by Buta et al. ["
+18].,18].
+ Ht appeared that 65% of the bars in SO0-S0/a galaxies have single Gaussian profiles. whereas 35% are best fitted by two Gaussian functions.," It appeared that $\%$ of the bars in S0-S0/a galaxies have single Gaussian profiles, whereas $\%$ are best fitted by two Gaussian functions."
+ Evpical examples of such profiles are shown in Figure 4., Typical examples of such profiles are shown in Figure 4.
+ The galaxies with DG bars tvpically have also significant higher Fourier modes in the bar region (m4. 6. 8). in addition to m=2 (51]).," The galaxies with DG bars typically have also significant higher Fourier modes in the bar region $m$ =4, 6, 8), in addition to $m$ =2 ([51])."
+ Lt was discussed by Buta et al., It was discussed by Buta et al.
+ that the DG-profiles are similar to those predicted by the simulation models (e.g... 5]) in which the bar transfers a large amount of angular momentum to the halo.," that the DG-profiles are similar to those predicted by the simulation models (e.g., [5]) in which the bar transfers a large amount of angular momentum to the halo."
+ An attempt to associate the DC-profiles to specific morphological structures was made by Laurikainen et al. , An attempt to associate the DG-profiles to specific morphological structures was made by Laurikainen et al. [
+51] who suggested that the fat or double-peaked Gaussian amplitude profiles ave due to two bar components. a long and narrow bar. and a shorter component in the inner parts of the bar (or an inner oval).,"51] who suggested that the fat or double-peaked Gaussian amplitude profiles are due to two bar components, a long and narrow bar, and a shorter component in the inner parts of the bar (or an inner oval)."
+ DG bars were found to be more prominent. not only in terms of Qu. but also in ele and bar length 51].," DG bars were found to be more prominent, not only in terms of $Q_g$, but also in $A_2$ and bar length [51]."
+ In Laurikainen et al. , In Laurikainen et al. [
+54] these inner bar components were associated mainly with barlenses (though some of them can be ovals). which are found to appear in 30% of barred. SO-SO/a ealaxies in NIIRSOS. A good example is NGC 4314 (Fig.,"54] these inner bar components were associated mainly with barlenses (though some of them can be ovals), which are found to appear in $\%$ of barred S0-S0/a galaxies in NIRS0S. A good example is NGC 4314 (Fig."
+ 1). in which the barlens is the fat elongated structure inside the bar.," 1), in which the barlens is the fat elongated structure inside the bar."
+ Erwin et al. , Erwin et al. [
+30] have discussed two SOs. in which a superposition of a classical and a pseudo-bulge was suggested.,"30] have discussed two S0s, in which a superposition of a classical and a pseudo-bulge was suggested."
+ These galaxies are NGC 2787 and NGC 3945. which form. part of NIBSOS. In both galaxies the component interpreted as a pseudo-bulge by Erwin et al.," These galaxies are NGC 2787 and NGC 3945, which form part of NIRS0S. In both galaxies the component interpreted as a pseudo-bulge by Erwin et al.,"
+ is called as a fat inner bar component in 51]. and more recently delined as a barlens by us 54].," is called as a fat inner bar component in [51], and more recently defined as a barlens by us [54]."
+ The above question was recently made by Buta et al. , The above question was recently made by Buta et al. [
+S6] with the main emphasis to test the hypothesis by Bournaud. & Combes 14]. in which multiple bar episodes are expected in the Hubble time.,"86] with the main emphasis to test the hypothesis by Bournaud $\&$ Combes [14], in which multiple bar episodes are expected in the Hubble time."
+ In this scenario bars form and evolve in galaxies when they have gas. and the evolution stops when the gas in used in star formation.," In this scenario bars form and evolve in galaxies when they have gas, and the evolution stops when the gas in used in star formation."
+ These stars are then transferred. into he bulge. for example by bars or spiral arms in the central regions of the galaxies.," These stars are then transferred into the bulge, for example by bars or spiral arms in the central regions of the galaxies."
+ When the central mass concentration ormed by star formation becomes very high. the bar will be destroved.," When the central mass concentration formed by star formation becomes very high, the bar will be destroyed."
+" Therefore. if bar strength varies over time the relative requency of galaxies in each (2, bin tells us the relative amount of time a galaxy spends in a certain bar state (strong. weak. non-barred)."," Therefore, if bar strength varies over time the relative frequency of galaxies in each $Q_g$ bin tells us the relative amount of time a galaxy spends in a certain bar state (strong, weak, non-barred)."
+ For spirals this was first tested by Block et al. , For spirals this was first tested by Block et al. [
+SO].,80].
+ They suggested that galaxies might have doubled their mass in 102 vears (see Fig., They suggested that galaxies might have doubled their mass in $^{10}$ years (see Fig.
+ 5. Left panel). evidenced by the extended: tail towards strong bars. and the lack of weak bars. which features are predicted in the strong gas accretion models by Bournaud & Combes 14].," 5, left panel), evidenced by the extended tail towards strong bars, and the lack of weak bars, which features are predicted in the strong gas accretion models by Bournaud $\&$ Combes [14]."
+" This test was later repeated by Buta. Laurikainen and Salo 87] for the same galaxy sample. but using the refined. bar orque method described in the previous section (we also discussed why the obtained (, distribution was different [rom that w Block et al. "," This test was later repeated by Buta, Laurikainen and Salo [87] for the same galaxy sample, but using the refined bar torque method described in the previous section (we also discussed why the obtained $Q_g$ distribution was different from that by Block et al. ["
+SO.,80]).
+ In Buta et al. , In Buta et al. [
+17] the bar and spiral llüxes were additionally separated from each other.,17] the bar and spiral fluxes were additionally separated from each other.
+" Although the correction allected €, in à few individual cases having very strong spiral armis. it barely alfected the Q, distribution."," Although the correction affected $Q_g$ in a few individual cases having very strong spiral arms, it barely affected the $Q_g$ distribution."
+" The refined Q, distribution (SY]. 17]) has a large number of weak bars lacking from that obtained by Block et al."," The refined $Q_g$ distribution ([87], [17]) has a large number of weak bars lacking from that obtained by Block et al."
+ Likewise. it yas a slightly smaller number of very strong bars.," Likewise, it has a slightly smaller number of very strong bars."
+" In fact. the obtained Q, distribution (sce Fig."," In fact, the obtained $Q_g$ distribution (see Fig."
+ Sa in. 17]) largely resembles he non-accretion model by Bournaud & Combes shown in Figure 5 (left panel). thus supporting the view that bars in spirals are fairly robust.," 8a in [17]) largely resembles the non-accretion model by Bournaud $\&$ Combes shown in Figure 5 (left panel), thus supporting the view that bars in spirals are fairly robust."
+ In Buta et al. , In Buta et al. [
+"86] the (, distribution for NIIGSOS was calculated.",86] the $Q_g$ distribution for NIRS0S was calculated.
+ Most importantly. a clear dillerence was ound between SOs and early-tvpe spirals (see Fig.," Most importantly, a clear difference was found between S0s and early-type spirals (see Fig."
+ 5. right panel).," 5, right panel)."
+ This was suggested to support the view according to which SOs have not acercted gas for a long time. evidenced by the lack of the extended tail. and the existence of a large number of weak bars.," This was suggested to support the view according to which S0s have not accreted gas for a long time, evidenced by the lack of the extended tail, and the existence of a large number of weak bars."
+ As discussed above (see Section :3.2) the smaller number of strong bars among the SOs can be due to a dilution elect caused by the more massive bulges and thicker disks in SOs., As discussed above (see Section 3.2) the smaller number of strong bars among the S0s can be due to a dilution effect caused by the more massive bulges and thicker disks in S0s.
+ However. it was also discussed by Buta et al. ," However, it was also discussed by Buta et al. ["
+"SG] that this cannot produce all of the dillerence in Q,y between the SOs and earlv-type spirals: although spirals have a larger number of",86] that this cannot produce all of the difference in $Q_g$ between the S0s and early-type spirals: although spirals have a larger number of
+A key question in the study of the formation and evolution of galaxies concerns the relationship between their observational properties and the large-scale cosmological environment.,A key question in the study of the formation and evolution of galaxies concerns the relationship between their observational properties and the large–scale cosmological environment.
+ In recent years. a flourishing of observational campaign has provided a detailed description of the evolution of the galaxy population in clusters.," In recent years, a flourishing of observational campaign has provided a detailed description of the evolution of the galaxy population in clusters."
+ Indeed. galaxy clusters play a key role in the characterization of the galaxy evolution.," Indeed, galaxy clusters play a key role in the characterization of the galaxy evolution."
+ Each cluster provides a large sample of galaxies. all placed at the same redshift.," Each cluster provides a large sample of galaxies, all placed at the same redshift."
+ Furthermore. clusters offer the possibility of sampling a variety of environments. from their dense core regions. to the outskirts where the properties of the cluster galaxy population tends to approach that of the field.," Furthermore, clusters offer the possibility of sampling a variety of environments, from their dense core regions, to the outskirts where the properties of the cluster galaxy population tends to approach that of the field."
+ A diversity of the galaxy population in nearby clusters. with respect to that in the field. was noticed already by ? and by ?..," A diversity of the galaxy population in nearby clusters, with respect to that in the field, was noticed already by \cite{1974ApJ...194....1O} and by \cite{1980ApJ...236..351D}."
+ Rich clusters were shown to contain a higher fraction of bulge—dominated (early type and SO) galaxies. and a correspondingly lower fraction of star forming galaxies. than poor systems.," Rich clusters were shown to contain a higher fraction of bulge--dominated (early type and S0) galaxies, and a correspondingly lower fraction of star forming galaxies, than poor systems."
+ ? noticed that moderately distant clusters ἐς~ 0.3) have a galaxy," \cite{1978ApJ...226..559B}
+ noticed that moderately distant clusters $z\sim 0.3$ ) have a galaxy"
+c3x1015erg.,$\simeq 3 \times 10^{48} \erg$.
+" With no mass estimate to calculate the kinetic energy, we can only speculate that the total energy is in the order of 103:50erg, larger than that of other ILOTs with the same time scale."," With no mass estimate to calculate the kinetic energy, we can only speculate that the total energy is in the order of $\sim 10^{49-50} \erg$, larger than that of other ILOTs with the same time scale."
+" As SN 2002bu is far out of the OTS, we conclude that it is not an ILOT, and the most likely explanation is that it is a peculiar class of SN."," As SN 2002bu is far out of the OTS, we conclude that it is not an ILOT, and the most likely explanation is that it is a peculiar class of SN."
+ This demonstrates that the ETD can also be used to identify transients with different physical properties than the ones considered as accretion powered ILOTs., This demonstrates that the ETD can also be used to identify transients with different physical properties than the ones considered as accretion powered ILOTs.
+" The transient M85 OT2006 was discovered in the lenticular galaxy M85 in Jan 7, 2006 (Kulkarni et al."," The transient M85 OT2006 was discovered in the lenticular galaxy M85 in Jan 7, 2006 (Kulkarni et al."
+ 2007a)., 2007a).
+ It had a peak luminosity of ~2x10*°ergs! and total radiated energy of Eyaa(M85 OT2006) ~6x1045erg over a duration of ~180days (Kulkarni et al., It had a peak luminosity of $\sim 2 \times 10^{40} \erg \s^{-1}$ and total radiated energy of $E_{\rm{rad}}$ (M85 OT2006) $\simeq 6 \times 10^{46} \erg$ over a duration of $\sim 180 \days$ (Kulkarni et al.
+ 2007a)., 2007a).
+ The progenitor’s mass was estimated to be <7Me (Ofek et al., The progenitor's mass was estimated to be $< 7~\rm{M_{\odot}}$ (Ofek et al.
+ 2008) and it was suggested that the origin of M85 OT2006 is a stellar merger (Kulkarni et al., 2008) and it was suggested that the origin of M85 OT2006 is a stellar merger (Kulkarni et al.
+ 2007a)., 2007a).
+ Rau et al. (, Rau et al. (
+2007) estimated that the effective temperature and the stellar radius at the time of peak luminosity were Teec4600K and R~3600Ro respectively.,2007) estimated that the effective temperature and the stellar radius at the time of peak luminosity were $T_{\rm{eff}} \simeq 4600\K$ and $R \simeq 3600~\rm{R_{\odot}}$ respectively.
+ At a later time in the eruption the star cooled down and expanded to have Tig~950K and R~20000Ro (see their table 2).," At a later time in the eruption the star cooled down and expanded to have $T_{\rm{eff}} \simeq 950\K$ and $R \simeq 20\,000~\rm{R_{\odot}}$ (see their table 2)."
+ The work of Shara et al. (, The work of Shara et al. (
+2010a) extends the previous nova models of Yaron et al. (,2010a) extends the previous nova models of Yaron et al. (
+2005).,2005).
+ The data of six of their models is given in their tables 2-4., The data of six of their models is given in their tables 2–4.
+ In the new extreme nova models of Shara et al. (, In the new extreme nova models of Shara et al. (
+2010a) the mass of the white dwarf is 0.4 — 0.65Μο and the accretion rate is as low as 10? — 107?Moyr!.,2010a) the mass of the white dwarf is $0.4$ – $0.65~\rm{M_{\odot}}$ and the accretion rate is as low as $10^{-12}$ – $10^{-10}~\rm{M_{\odot} \yr^{-1}}$.
+ The ejected mass in the novae ranges between 5.3x10* — 2.2x1073Mo at velocities of 150 — 480kms!., The ejected mass in the novae ranges between $5.3 \times 10^{-4}$ – $ 2.2 \times 10^{-3}~\rm{M_{\odot}}$ at velocities of $150$ – $480 \km \s^{-1}$.
+ Shara et al. (, Shara et al. (
+"2010a) claim that the eruption of M85 OT2006 was a nova, and suggest that it can be explained in the frame of the new extreme nova models.","2010a) claim that the eruption of M85 OT2006 was a nova, and suggest that it can be explained in the frame of the new extreme nova models."
+ Their conclusion is mainly based on the following. (, Their conclusion is mainly based on the following. (
+1) Their new results showing that novae can reach comparable luminosities of ILOTs such as M85 OT2006 and M31 RV (few x10’ Lo). (,1) Their new results showing that novae can reach comparable luminosities of ILOTs such as M85 OT2006 and M31 RV (few $\times 10^7~\rm{L_{\odot}}$ ). (
+"2) The models produce red eruptions, as the spectra of ILOTs.","2) The models produce red eruptions, as the spectra of ILOTs."
+ We hereby show that the mass ejected in the eruption of M85 OT2006 is much larger than the nova models can produce., We hereby show that the mass ejected in the eruption of M85 OT2006 is much larger than the nova models can produce.
+ Let us consider the M85 OT2006 transient according to data given by Rau et al. (, Let us consider the M85 OT2006 transient according to data given by Rau et al. (
+"2007), discussed above.","2007), discussed above."
+" The column density required above the photosphere is given by where p is the average density above the photosphere, Ar is the thickness of the shell above the photosphere, and K is its average opacity."," The column density required above the photosphere is given by where $\rho$ is the average density above the photosphere, $\Delta r$ is the thickness of the shell above the photosphere, and $\kappa$ is its average opacity."
+ We parameterize the thickness of the shell with Ar=GR with 8~ 0.1., We parameterize the thickness of the shell with $\Delta r = \beta R$ with $\beta \sim 0.1$ .
+ The final mass we obtain for the shell above the photosphere changes by a factor <2 for 0.01«60.1., The final mass we obtain for the shell above the photosphere changes by a factor $<2$ for $0.01 \leqslant \beta \leqslant 0.1$.
+ The reason is because the value of 6 determines the density that weakly influences the opacity., The reason is because the value of $\beta$ determines the density that weakly influences the opacity.
+ We use opacities from Ferguson et al. (, We use opacities from Ferguson et al. (
+"2005) (slightly extrapolated), using compositions from Asplund et al. (","2005) (slightly extrapolated), using compositions from Asplund et al. ("
+2004) with hydrogen abundance X=0.7 and metallicity Z=0.1; other composition from Lodders (2003) gives very close opacity values.,2004) with hydrogen abundance $\rm{X}=0.7$ and metallicity $\rm{Z}=0.1$; other composition from Lodders (2003) gives very close opacity values.
+ The total mass above the photosphere is Using the temperature at peak luminosity we find that the opacity at the peak luminosity is &~1.3x 10-?., The total mass above the photosphere is Using the temperature at peak luminosity we find that the opacity at the peak luminosity is $\kappa \simeq 1.3 \times 10^{-3}$ .
+ Substituting theopacity and the radius at peak luminosity in equation (7)) we find that the mass above the photosphere at the time of peak luminosity is Mpn0.2 Mo., Substituting theopacity and the radius at peak luminosity in equation \ref{eq:M_ph}) ) we find that the mass above the photosphere at the time of peak luminosity is $M_{\rm{ph}} \sim 0.2~\rm{M_{\odot}}$ .
+ Our results are summarized in Table 2.., Our results are summarized in Table \ref{tab:M_ph}.
+" The real amount of mass is even larger, as some mass is well above the photosphere and we expect a large amount of mass to be below the photosphere as well."," The real amount of mass is even larger, as some mass is well above the photosphere and we expect a large amount of mass to be below the photosphere as well."
+ The largest value of ejected mass in the nova models of Shara et al. (, The largest value of ejected mass in the nova models of Shara et al. (
+20102) is 2.2x10?Mo or two orders of magnitude below the expected ejected mass we calculate.,2010a) is $ 2.2 \times 10^{-3}~\rm{M_{\odot}}$ or two orders of magnitude below the expected ejected mass we calculate.
+ We therefore conclude that M85 OT2006 is not a nova., We therefore conclude that M85 OT2006 is not a nova.
+" Previously in KFS10, we used the assumption of Ofek et al. ("," Previously in KFS10, we used the assumption of Ofek et al. ("
+"2008) that the total ejected mass of M85 OT2006 is 0.1 Mo, and consequently obtained a total energy of ~1.4x1047 erg.","2008) that the total ejected mass of M85 OT2006 is $0.1~\rm{M_{\odot}}$ , and consequently obtained a total energy of $\sim 1.4 \times 10^{47} \erg$ ."
+ Our new estimate of ejected mass is much higher and we update our estimate of the total energy in M85 OT2006., Our new estimate of ejected mass is much higher and we update our estimate of the total energy in M85 OT2006.
+" Our new estimate, taking ejected mass velocity of ~870kms! (Rau et al."," Our new estimate, taking ejected mass velocity of $\sim 870 \km \s^{-1}$ (Rau et al."
+" 2007) is 1.6x1048 — 4.6x107?erg, corresponding to 0.2Mo and 0.6Mo, respectively."," 2007) is $1.6 \times 10^{48}$ -- $4.6 \times 10^{48}\erg$, corresponding to $0.2~\rm{M_{\odot}}$ and $0.6~\rm{M_{\odot}}$, respectively."
+ These values are summarized in Table 1 and the updated location of M85 OT2006 is shown in the ETD (Figure 1))., These values are summarized in Table \ref{tab:data} and the updated location of M85 OT2006 is shown in the ETD (Figure \ref{fig:totEvst}) ).
+ We calculate the total (radiated and kinetic) energy of the new nova modelsof Shara el al. (, We calculate the total (radiated and kinetic) energy of the new nova modelsof Shara el al. (
+2010a) andplot them in Figure 1..,2010a) andplot them in Figure\ref{fig:totEvst}. .
+ We include all models from Yaron et al. (, We include all models from Yaron et al. (
+2005),2005)
+revisit the same image twice for a total of 3 snapshots per field of view.,revisit the same image twice for a total of 3 snapshots per field of view.
+" Observing in 4 bands, for a total of 28 days integration time over 2 years, such a survey expects to see ~10 red supergiant PISNe and ~300 CCSNe."," Observing in 4 bands, for a total of 28 days integration time over 2 years, such a survey expects to see $\sim$ 10 red supergiant PISNe and $\sim$ 300 CCSNe."
+" To the extent that PISN spectra can be represented as a distribution of blackbodies at different temperatures, since the temperature and redshift would be degenerate, it will be impossible to acquire photometric redshifts without further information about the SN epoch."," To the extent that PISN spectra can be represented as a distribution of blackbodies at different temperatures, since the temperature and redshift would be degenerate, it will be impossible to acquire photometric redshifts without further information about the SN epoch."
+" However, our simulated spectra show significant deviations from a blackbody in the UV (I< 3500A)) due to metal-line blanketing in the SN photosphere, providing spectral and photometric signatures that could be used as redshift indicators, depending on their strength."," However, our simulated spectra show significant deviations from a blackbody in the UV $l < 3500$ ) due to metal-line blanketing in the SN photosphere, providing spectral and photometric signatures that could be used as redshift indicators, depending on their strength."
+" Although the UV flux of PISNe is relatively short lived, the more massive PISNe stay bright in its rest frame visible band for over a year."," Although the UV flux of PISNe is relatively short lived, the more massive PISNe stay bright in its rest frame visible band for over a year."
+" Given this brightness and long intrinsic duration, coupled with the (1+z) time dilation at high redshifts, it is conceivable that PISNe could contribute to the luminosity function of all objects at high redshifts whengalaxies were dim."," Given this brightness and long intrinsic duration, coupled with the $(1+z)$ time dilation at high redshifts, it is conceivable that PISNe could contribute to the luminosity function of all objects at high redshifts whengalaxies were dim."
+" Figure 9 illustrates the luminosity function of PISNe at ~4000A,, calculated using the helium core progenitor models for PISN luminosity, and the Pop III Flat or Pop III Salpeter models for the star formation rate."," Figure \ref{luminosity_function} illustrates the luminosity function of PISNe at $\sim$, calculated using the helium core progenitor models for PISN luminosity, and the Pop III Flat or Pop III Salpeter models for the star formation rate."
+" Shown for comparison are the projected galaxy luminosity functions at high redshifts, using the ? best fit Schechter parameterization for the UV luminosity function, and shifting to the visible band using U—Vz0.4, 0.3 for z—7, 8 respectively, measured using the Spitzer Infrared Array Camera (??).."," Shown for comparison are the projected galaxy luminosity functions at high redshifts, using the \citet{Bouwens2011b} best fit Schechter parameterization for the UV luminosity function, and shifting to the visible band using $U-V \approx 0.4$, $0.3$ for $z=7,$ $8$ respectively, measured using the Spitzer Infrared Array Camera \citep{Labb'e2010, Labb'e2010a}."
+" Applying this U-V shift is a crude approximation, as luminous and faint galaxies have different rest frame UV-to-optical color; however, we are most interested in the bright end of the luminosity function, where this current U-V measurement is applicable."," Applying this U-V shift is a crude approximation, as luminous and faint galaxies have different rest frame UV-to-optical color; however, we are most interested in the bright end of the luminosity function, where this current U-V measurement is applicable."
+ The luminosity function for PISNe implied by our Pop III IMF models overlaps with the galaxy luminosity function at the brightest magnitudes., The luminosity function for PISNe implied by our Pop III IMF models overlaps with the galaxy luminosity function at the brightest magnitudes.
+" If a top-heavy Pop III IMF was solely responsible for reionization, PISNe will contaminate the brightest end of the galaxy luminosity function, unless great care is taken to remove these supernovae."," If a top-heavy Pop III IMF was solely responsible for reionization, PISNe will contaminate the brightest end of the galaxy luminosity function, unless great care is taken to remove these supernovae."
+" Since the volumetric count of the brightest galaxies and PISNe is very low, it will take a wide infrared survey to observe this effect."," Since the volumetric count of the brightest galaxies and PISNe is very low, it will take a wide infrared survey to observe this effect."
+" In our discussion we ignored complicating factors such as metallicity and rotation, and calculated the PISN and CCSN event rate using only the SN progenitor mass range along with the star formation rate."," In our discussion we ignored complicating factors such as metallicity and rotation, and calculated the PISN and CCSN event rate using only the SN progenitor mass range along with the star formation rate."
+" However, at low redshifts z«1, the measured CCSN rate is a factor of ~2 smaller than that predicted by the analogous calculation using the measured cosmic star formation rate."," However, at low redshifts $z<1$, the measured CCSN rate is a factor of $\sim 2$ smaller than that predicted by the analogous calculation using the measured cosmic star formation rate."
+ The discrepancy is likely due to many intrinsically low-luminosity or obscured SNe being missed in surveys (?).., The discrepancy is likely due to many intrinsically low-luminosity or obscured SNe being missed in surveys \citep{Horiuchi2011}.
+" As this discrepancy is lower than the uncertainty in our SFR. model parameters, and we already account for lower intrinsic luminosities for the lower progenitor mass PISNe, we do not take obscuration into account for our predictions of the SN rate as seen by JWST."," As this discrepancy is lower than the uncertainty in our SFR model parameters, and we already account for lower intrinsic luminosities for the lower progenitor mass PISNe, we do not take obscuration into account for our predictions of the SN rate as seen by JWST."
+ The IMF of early stellar populations responsible for reionization should also leave an imprint on the metal enrichment pattern via their SN products., The IMF of early stellar populations responsible for reionization should also leave an imprint on the metal enrichment pattern via their SN products.
+" So far, the abundance patterns observed to date in extremely metal-deficient stars in the Galactic halo (?) are more consistent with an IMF that produced much more CCSNe instead of PISNe (?).."," So far, the abundance patterns observed to date in extremely metal-deficient stars in the Galactic halo \citep{Beers2005} are more consistent with an IMF that produced much more CCSNe instead of PISNe \citep{Joggerst2010}."
+" However, in previous surveys, subtle selection effects might have disfavored finding PISN-enriched stars; the metal yields of PISNe are so high that the metal abundances of stars formed out of PISN ejecta (7) are already higher than the metallicity range targeted by metal-deficient star surveys (?).."," However, in previous surveys, subtle selection effects might have disfavored finding PISN-enriched stars; the metal yields of PISNe are so high that the metal abundances of stars formed out of PISN ejecta \citep{Greif2008} are already higher than the metallicity range targeted by metal-deficient star surveys \citep{Karlsson2008}."
+" Large carbon enhancements observed in metal-poor stars, when interpreted as the outcome of pollution by winds from binary companions that have gone through the AGB phase, suggest the existence of a large number of intermediate-mass stars (~1— 8Mc) at high redshifts (??).."," Large carbon enhancements observed in metal-poor stars, when interpreted as the outcome of pollution by winds from binary companions that have gone through the AGB phase, suggest the existence of a large number of intermediate-mass stars $\sim 1-8 M_{\odot}$ ) at high redshifts \citep{Tumlinson2007,Tumlinson2007a}."
+" Alternatively, nucleosynthesis in faint CCSNe from higher mass stars could also explain the observed carbon enhancement in metal-poor stars (?).."," Alternatively, nucleosynthesis in faint CCSNe from higher mass stars could also explain the observed carbon enhancement in metal-poor stars \citep{Iwamoto2005}."
+" Observing the Type Ia SN rate during the epoch of reionization will be an complementary way to test these models, and constrain the number of intermediate-mass stars at high redshifts."," Observing the Type Ia SN rate during the epoch of reionization will be an complementary way to test these models, and constrain the number of intermediate-mass stars at high redshifts."
+ The predicted initial mass range of ~140 to 260 Mo for PISN progenitors assumed the stars to be non-rotating (?).., The predicted initial mass range of $\sim$ 140 to 260 $M_{\odot}$ for PISN progenitors assumed the stars to be non-rotating \citep{Heger2002}. .
+" However, observations find that at very low metallicities, stars rotate faster (?).."," However, observations find that at very low metallicities, stars rotate faster \citep{Martayan2007}."
+" The fast rotation of the first stars is supported by the latest hydrodynamic simulations of their formation (?),, and also by observations of anomalously high abundances of Ba and La with respect to Fe in ancient stars (?),, which could originate in metal-poor fast- massive stars."," The fast rotation of the first stars is supported by the latest hydrodynamic simulations of their formation \citep{Stacy2011}, and also by observations of anomalously high abundances of Ba and La with respect to Fe in ancient low-mass stars \citep{Chiappini2011}, , which could originate in metal-poor fast-rotating massive stars."
+" Generally, rotationshould increase the required PISN progenitor mass by increasing mass loss."," Generally, rotationshould increase the required PISN progenitor mass by increasing mass loss."
+The general nucleosvnthesis associated with rapid acliabatic expansion and freeze-oul from Nuclear Statistical Equilibrinm (NSE) at high entropy-per-baryon. s. was first considered in the landmark paper bv Wagoner.Fowler.&Iovle(1967) (hereafter WEIL).,"The general nucleosynthesis associated with rapid adiabatic expansion and freeze-out from Nuclear Statistical Equilibrium (NSE) at high entropy-per-baryon, $s$, was first considered in the landmark paper by \cite{wag} (hereafter WFH)."
+" Those authors concentrated on the environments associated wilh exploding supermassive objects. where s/h,~10. and on Dig Dang Nucleosvnthesis (BBN). where s/h,zzLOM."," Those authors concentrated on the environments associated with exploding supermassive objects, where $s/k_b\sim 10^3$, and on Big Bang Nucleosynthesis (BBN), where $s/k_b\approx10^{10}$."
+ In both supermassive objects aud BBN the conditions are expected to be proton-rich (preponderance of protons over neutrons) and the characteristic expansion timescale large (e.g.. 7444~1008 for BBN).," In both supermassive objects and BBN the conditions are expected to be proton-rich (preponderance of protons over neutrons) and the characteristic expansion timescale large (e.g., $\tau_{\rm dyn}\sim 100{\rm s}$ for BBN)."
+ The physics of relativistic outflows aud potential nucleosvinthesis in (hese sites has also been the subject of some recent. attention (Pruet.FullerBurrows.&Mever2001:Otsukiοἱal.," The physics of relativistic outflows and potential nucleosynthesis in these sites has also been the subject of some recent attention \citep{ultra,thomps,otsuki}."
+ 2000).. llere we extend the WFII study., Here we extend the WFH study.
+ We consider [reeze-out. from NSE over a wide range of entropv-per-baryon spanning that in WEIL ancl a range of neutron-to-proton ratios. all the way [rom proton-rich. (o neutron-rich.," We consider freeze-out from NSE over a wide range of entropy-per-baryon spanning that in WFH, and a range of neutron-to-proton ratios, all the way from proton-rich to neutron-rich."
+" NSE freezeout in these scenarios is calculated for expansion limescales ranging [rom those appropriate for relativistic flows from compact objects to those associated with BBN (10""s<zq4,1005).", NSE freezeout in these scenarios is calculated for expansion timescales ranging from those appropriate for relativistic flows from compact objects to those associated with BBN $10^{-6}{\rm s}<\tau_{\rm dyn}<100{\rm s})$.
+ We find that the general nuclear physics features of BBN are a recurring theme throughout all of these parameter ranges., We find that the general nuclear physics features of BBN are a recurring theme throughout all of these parameter ranges.
+ In particular. 71 svnthesis can be significant for relativistic flows.," In particular, ${^2}{\rm H}$ synthesis can be significant for relativistic flows."
+ The expected possibly neutron-rich conditions of a GRB fireball are. in a sense. the isospin mirror of those for BBN.," The expected possibly neutron-rich conditions of a GRB fireball are, in a sense, the isospin mirror of those for BBN."
+ However. we also show that it is not sufficient to simply consider only (he NSE freeze-out nucleosvnthesis in some parameter regimes appropriate for. e.g.. GRB fireballs.," However, we also show that it is not sufficient to simply consider only the NSE freeze-out nucleosynthesis in some parameter regimes appropriate for, e.g., GRB fireballs."
+ In fact. the dvnaimics of these fireballs ean ciffer dramatically on a microscopic scale from (he conditions treated by WEIL.," In fact, the dynamics of these fireballs can differ dramatically on a microscopic scale from the conditions treated by WFH."
+ In particular. for initially neutron-rich material (he few protons accelerate with the relativistic photon and. e= pair fluid to a very high Lorentz [actor (5~10?— 107).," In particular, for initially neutron-rich material, the few protons accelerate with the relativistic photon and $e^{\pm}$ pair fluid to a very high Lorentz factor $\gamma\sim10^2-10^3$ )."
+ In this scenario. there will be high energy (~ GeV) collisions of protons on left-behind neutron stragglers (Fuller.Pruet.&Abazajian2000).," In this scenario, there will be high energy $\sim {\rm GeV}$ ) collisions of protons on “left-behind” neutron stragglers \citep{fpa}."
+. It has been shown that high energy collisions of (his kind can result in significant destruction and/or production of light nuclei 1983).., It has been shown that high energy collisions of this kind can result in significant destruction and/or production of light nuclei \citep{dim}.
+ We show below that these non-thermal nuclear reactions can modilv significantly the simple BBN-like NSE freeze-ont abundances of 7H and other species., We show below that these non-thermal nuclear reactions can modify significantly the simple BBN-like NSE freeze-out abundances of ${^2}{\rm H}$ and other species.
+ In the right conditions. IL number fraction vields can approach YyL~1056.," In the right conditions, ${^2}{\rm H}$ number fraction yields can approach $Y_{\rm D}\sim 10\%$."
+ This is a stageerinelySSSt. high5S vield given the known. and small. primordial ?II abundance.," This is a staggeringly high yield given the known, and small, primordial ${^2}\rm H$ abundance."
+ The deuteron is a particularly interesting nucleus because it has a binding energv οἱ only 222.2 MeV. and is notoriously Iragile and difficult to synthesize., The deuteron is a particularly interesting nucleus because it has a binding energy of only $\approx$ 2.2 MeV and is notoriously fragile and difficult to synthesize.
+ At present it is thought that essentially all of the deuterium in the universe is primordial., At present it is thought that essentially all of the deuterium in the universe is primordial.
+ Estimates of the primordial “11/11 come from measurements of the 82 km/s isotope shift in the Lyman a line in a handful ol high redshift’ Lyman limit clouds., Estimates of the primordial $\hh$ /H come from measurements of the 82 km/s isotope shift in the Lyman $\alpha$ line in a handful of high redshift Lyman limit clouds.
+ These are clouds that interseet our lines of sight to, These are clouds that intersect our lines of sight to
+aud the stellaz component is caleulated using an N-body echuique.,and the stellar component is calculated using an N-body technique.
+ The initial conditiois wore built according to Spriugel et al. (, The initial conditions were built according to Springel et al. (
+2005). based ou the analytic work of Mo et al. (,"2005), based on the analytic work of Mo et al. ("
+1998).,1998).
+ The two model galaxies were chosen sucht hat trev represent a Milkv Way type and a simall spiral ooOalaxy. with the mass ratio of the two ealaxics Όσιο 8:1.," The two model galaxies were chosen such that they represent a Milky Way type and a small spiral galaxy, with the mass ratio of the two galaxies being 8:1."
+ There‘Ore. the total mass of the model galaxies A and [2L) Is « 1072U ph1 M. aud L671Iss 4107 53 M;. respectivev.," Therefore, the total mass of the model galaxies A and B is $\times$ $^{12}$ $h^{-1}$ $M_{\sun}$ and $\times$ $^{11}$ $h^{-1}$ $M_{\sun}$, respectively."
+ The combined N-bodv/SPII simulation then calculates 5 Cir of evoluion., The combined N-body/SPH simulation then calculates 5 Gyr of evolution.
+ For every time step. we know he veocity of cach particle aud can hence extract realistic 2D veocitv fields of the eus.," For every time step, we know the velocity of each particle and can hence extract realistic 2D velocity fields of the gas."
+ons Concerning the osoval. aliennneut aud impact )araucters. we follow he notation iutroduced by Duc et al. (," Concerning the spatial alignment and impact parameters, we follow the notation introduced by Duc et al. ("
+2000) o describe the interaction geometry.,2000) to describe the interaction geometry.
+ The )araiueer corresponds to the minim separation of he galaxies rajectories. as if they were ponit masses Ou Weplerian orbits.," The parameter corresponds to the minimum separation of the galaxies' trajectories, as if they were point masses on Keplerian orbits."
+ Additionally. two angles. O aud ᾧ dchine the spalal orientation of the disc.," Additionally, two angles, $\Theta$ and $\Phi$ define the spatial orientation of the disc."
+ For he complete sample. we selected. the aliguimeuts iu such a way as to cover as many ecoletrics as possible. including nuxx and major mergers aud fiv-bys (achieved lyVv increasing the minium separation).," For the complete sample, we selected the alignments in such a way as to cover as many geometries as possible, including minor and major mergers and fly-bys (achieved by increasing the minimum separation)."
+ For this analvsis NXο just use three differeut mteraction eeonmetzies. which NXο list iu Table 1..," For this analysis we just use three different interaction geometries, which we list in Table \ref{sims}."
+ These three simulations allow us to study clifferent classes of ποιαισα]. cistortiois. detailed in Sect. 3..," These three simulations allow us to study different classes of kinematical distortions, detailed in Sect. \ref{extract}."
+" Additional important quantities of the simulations. as e.g. the particle αμαος, are smnunarized iu Table 2.."," Additional important quantities of the simulations, as e.g. the particle numbers, are summarized in Table \ref{galaxyproperties_resolution}."
+ Throughout the paper we adopt the standard ACDAM cosinoloey wi hOQy=.7. 0enti=0.3. aud h=0.7.," Throughout the paper we adopt the standard $\Lambda$ CDM cosmology with $\Omega_\Lambda=0.7$, $\Omega_m=0.3$, and h=0.7."
+ Il OCer to construct realistic 2D velocity fields. we projec all gas particles of the N-bodyv/SPII simulations [9]uto a Cartesian. equally spaced exid.," In order to construct realistic 2D velocity fields, we project all gas particles of the N-body/SPH simulations onto a Cartesian, equally spaced grid."
+ The spacing is Crosea such. that the spatial resolution at the assumed redshit of the model galaxy correspouds to the angular resolujon of current state-ofthe-art observations.," The spacing is chosen such, that the spatial resolution at the assumed redshift of the model galaxy corresponds to the angular resolution of current state-of-the-art observations."
+ For t lewrole investigated redshift range from z- to z=1.0 iMο adopt an aneular resolution typical for IFU or EPI [9]servations. namely 0.0 (ee. Chemin et al.," For the whole investigated redshift range from z=0 to z=1.0 we adopt an angular resolution typical for IFU or FPI observations, namely 0.4"" (e.g. Chemin et al."
+ 2006)., 2006).
+ The angular resolution o|: SAURON. for example. would be IY or 0.97.," The angular resolution of SAURON, for example, would be 0.3"" or 0.9""."
+ For iuterinediate redshüfts we additiona hose 0.257. which is the pixel scale along the FORS2 s sed bv Ziegler et al. (," For intermediate redshifts we additionally chose 0.25"", which is the pixel scale along the FORS2 slit as used by Ziegler et al. ("
+2006) and 0.527. which is f et size of FLAMES/GIRAFFE at VET (c.g. Flores al.,"2006) and 0.52"", which is the lenslet size of FLAMES/GIRAFFE at VLT (e.g. Flores et al."
+ 2006)., 2006).
+ We calculate for each redshift he physic 10360)utiion according o the eiven aieular resolution usi 16 concordance cosmological iiocel (see Table 3)., We calculate for each redshift the physical resolution according to the given angular resolution using the concordance cosmological model (see Table \ref{redshift}) ).
+ T vclocity field of the galaxy is binned using this spatial 10360)mlon., The velocity field of the galaxy is binned using this spatial resolution.
+ Frei the kuowledeoe of the full 3D velociv fields ancl the iucraction historv of the galaxies in our siuulations iMο define three kinematical classes. in principle analogue tc) classificatious in Flores et al. (," From the knowledge of the full 3D velocity fields and the interaction history of the galaxies in our simulations we define three kinematical classes, in principle analogue to classifications in Flores et al. ("
+2006) or Krajnovic et al. (,2006) or Krajnović et al. (
+2006): The main question tlat we investigate in this paper is how such classifications depend ou the redshift of re observed salaxy. ie. the actual spatial resolution of ie galaxy.,"2006): The main question that we investigate in this paper is how such classifications depend on the redshift of the observed galaxy, i.e. the actual spatial resolution of the galaxy."
+ Iu this contest the κος plays a crucial 1¢de. as it typically exceeds the augular resolution of the oeistriuueit.," In this context the seeing plays a crucial role, as it typically exceeds the angular resolution of the instrument."
+ To simulate secius effects on our velocity-field 1ieasureunieits. a c0wolution with a Gaussian poiut nixead function was appied.," To simulate seeing effects on our velocity-field measurements, a convolution with a Gaussian point spread function was applied."
+ We adopted a value of 087 or the FWA of the Gaussian seeiug. which is a typical Vidue for eround. based observatious (see c.g. Jagecr et al.," We adopted a value of 0.8"" for the FWHM of the Gaussian seeing, which is a typical value for ground based observations (see e.g. Jägger et al."
+ 2)D., 2004).
+ Tje appearance of the velocity fields is therefore clenunmated hy the secing., The appearance of the velocity fields is therefore dominated by the seeing.
+ Note that we do not calculate a1 evolution of properics of the galaxy with redshift but Edy how a eiven kinematical state of a galaxy is observed a different recshifts., Note that we do not calculate an evolution of properties of the galaxy with redshift but study how a given kinematical state of a galaxy is observed at different redshifts.
+ All studies imnoenutioned above are inainly ποσο cleuinated., All studies mentioned above are mainly seeing dominated.
+ Wo addijonally study the possibility to XCentity mergers at high redshift (22) with an adaptive Qfies Iustrument such as SINFONT at the VET., We additionally study the possibility to identify mergers at high redshift $\sim$ 2) with an adaptive optics instrument such as SINFONI at the VLT.
+" For this investigation we adop an angular resolution of 0.15"" as aclüeved by Cieuzel et al. ("," For this investigation we adopt an angular resolution of 0.15"" as achieved by Genzel et al. ("
+2006).,2006).
+ Iu the subsequent sub-sectious we will svstemiaticallv investigate velocity fields from cach of the kiueimatical classes defined im Sect. 3.., In the subsequent sub-sections we will systematically investigate velocity fields from each of the kinematical classes defined in Sect. \ref{extract}.
+ The focus lies ou the visibility of kinematical distortions as a function of redshift., The focus lies on the visibility of kinematical distortions as a function of redshift.
+ Iu order to «nantifv the distortions and to interpret the partly complex structures iu the velocity fields we use the kincimnetry package o ‘Wrajnovie et al. (, In order to quantify the distortions and to interpret the partly complex structures in the velocity fields we use the kinemetry package of Krajnović et al. (
+2006).,2006).
+ This analysis Is based on the assumption. that the mean velocity along best fitting cllipses can be reproduced by a cosine law. 1.6. Ww.here « is the lenet1 of senianajor axis of the cllipse. V Is for discs the azimuthal anele measured from the major axis iu the plane o| the Ooealaxv.," This analysis is based on the assumption, that the mean velocity along best fitting ellipses can be reproduced by a cosine law, i.e. where $a$ is the length of semi-major axis of the ellipse, $\Psi$ is for discs the azimuthal angle measured from the major axis in the plane of the galaxy."
+ Note that in Oogeneral a siuple siue correction for the inclination of the dise is a»plied iu observations too., Note that in general a simple sine correction for the inclination of the disc is applied in observations too.
+ The velocity fields presented, The velocity fields presented
+The external galactic fiekl due to the bulge is represented by the potential of the Schwarzschild model (Sehwarzschild1979).,The external galactic field due to the bulge is represented by the potential of the Schwarzschild model \citep{scw}.
+. The Schwarzschild model is a non-rotating. consistent triaxial ellipsoid with axis ratios 2: 1.25: 1. typical of many galaxies ," The Schwarzschild model is a non-rotating, self-consistent triaxial ellipsoid with axis ratios 2: 1.25: 1, typical of many galaxies \citep{ber91}."
+"Defining adimensional units as rj being the bulge core radius. the potential (x.i.z) is expressed as the sum of a spherically simmetric term. ὃν (07= r/rj). which corresponds to the potential given by a density distribution following the modified IIubble's law pr)=p[1+(η. (po=My rj). plus two spherical armonics. 94(z.7) and Pole.ii): where and Ad, is the bulge"," Defining adimensional units as $r_b$ being the bulge core radius, the potential $\Phi(x^{\prime},y^{\prime},z^{\prime})$ is expressed as the sum of a spherically simmetric term, $\Phi_{r^{\prime}}(r^{\prime})$, $r^{\prime}=r/r_b$ ), which corresponds to the potential given by a density distribution following the modified Hubble's law $\rho(r^{\prime})=\rho_0 \left[1+\left(r^{\prime}\right)^2\right]^{-3/2}$, $\rho_0\equiv M_b/r_b^3 $ ), plus two spherical armonics, $\Phi_1(z^{\prime},r^{\prime})$ and $\Phi_2(x^{\prime},y^{\prime},r^{\prime})$ where and $M_b$ is the bulge"
+There remain many pieces needed to assemble the circiumnuclear starburst ring puzzle.,There remain many pieces needed to assemble the circumnuclear starburst ring puzzle.
+ A large-scale bar drives dust ancl gas towards the center of the galaxy., A large-scale bar drives dust and gas towards the center of the galaxy.
+ At the boundary of the bulge and the disk. a starlormine ring develops.," At the boundary of the bulge and the disk, a starforming ring develops."
+ Within (his ring. there is a fairly coherent loosely wound dust spiral.," Within this ring, there is a fairly coherent loosely wound dust spiral."
+" [tis unclear. however. whether this nuclear dust spiral is composed of dust aud gas originating Irom the large-scale bar and/or the ring itself. or if it is native dust that is ""stiimred-up by the starburst ring."," It is unclear, however, whether this nuclear dust spiral is composed of dust and gas originating from the large-scale bar and/or the ring itself, or if it is native dust that is “stirred-up” by the starburst ring."
+ Simulations by Regan&Teuben and Maciejewski(2004) imply that bars in low velocitv-dispersion svstemis are only capable ol driving gas down to the radius at which the ring formsbut not furthersuggesting that native dust comprises the nuclear spiral., Simulations by \citet{regan04} and \citet{maciejewski04b} imply that bars in low velocity-dispersion systems are only capable of driving gas down to the radius at which the ring forms—but not further—suggesting that native dust comprises the nuclear spiral.
+ A galaxy. with a high central velocitv-dispersion. however. could potentially have mass transported into the ring (Maciejewski2004)..," A galaxy with a high central velocity-dispersion, however, could potentially have mass transported into the ring \citep{maciejewski04b}. ."
+Following similar. transformations with the remaining Eqs. (,Following similar transformations with the remaining Eqs. (
+6)-(9). we arrived at the linearized equations needed for stability analysis.,"6)-(9), we arrived at the linearized equations needed for stability analysis."
+" We consider perturbations with the space- dependence «export—fk-r). where k=(k,.O.Kk.) Is the wavevector."," We consider perturbations with the space-time dependence $\propto \exp ( \sigma t - i \vec{k} \cdot \vec{r})$, where $\vec{k}= (k_{s}, 0, k_{z})$ is the wavevector."
+" Then. the linearized MHD-equations read with accuracy in the lowest order in /l/s: where w,=gk."," Then, the linearized MHD-equations read with accuracy in the lowest order in $\lambda/s$: where $\omega_{\eta}= \eta k^2$."
+ The dispersion equation. corresponding to Eqs. (," The dispersion equation, corresponding to Eqs. ("
+18)-(22) is rather complex in the general case.,18)-(22) is rather complex in the general case.
+ Therefore. we consider only a particular case when the wavevector is perpendicular to the magnetic field. k-B=0.," Therefore, we consider only a particular case when the wavevector is perpendicular to the magnetic field, $\vec{k} \cdot \vec{B} = 
+0$."
+ After some algebra. Eqs. (," After some algebra, Eqs. ("
+18)-(22) can be combined into a sixth-order dispersion relation. The coefficients of this equation are expressed in terms of characteristic frequencies. where jt=- . The characteristic frequences are given by modes d,"18)-(22) can be combined into a sixth-order dispersion relation, The coefficients of this equation are expressed in terms of characteristic frequencies, where $\mu = k^{2}_{z}/k^{2}$ The characteristic frequences are given by where $c_{s} = \sqrt{\gamma p/ \rho}$ is the sound speed."
+, Eq. (
+escribes that can generally exist in a compressible: rotating. magnetized gas.,"23) describes six modes that can generally exist in a compressible, rotating, magnetized gas."
+" In the non-dissipative. incompressible limit when the sound speed c, is very large. Eq. ("," In the non-dissipative, incompressible, limit when the sound speed $c_{s}$ is very large, Eq. ("
+2) reduces to the dispersion relation for the inertial waves. This equation allows for unstable solutions if the Rayleigh criterion is fulfilled. «7<0.,"23) reduces to the dispersion relation for the inertial waves, This equation allows for unstable solutions if the Rayleigh criterion is fulfilled, $\kappa^{2} < 0$."
+ Note that Eq. (, Note that Eq. (
+23) does not describethe magnetorotational instability since it does not occur for perturbations with K-B=0.,23) does not describethe magnetorotational instability since it does not occur for perturbations with $\vec{k} \cdot \vec{B} = 0$.
+ In non-dissipative limit. 7-0. Eq. (," In non-dissipative limit, $\eta \rightarrow 0$, Eq. ("
+23) recovers the dispersion relation by Bonanno Urpin (2006) for MD Wwith Mk-B=0. The authors argued that Eq. (,"23) recovers the dispersion relation derived by Bonanno Urpin (2006) for perturbations with $\vec{k} \cdot \vec{B} = 0$, The authors argued that Eq. ("
+25) has unstable solutions even when the criteria of the magnetorotational and Rayleigh instability are not satisfied.,25) has unstable solutions even when the criteria of the magnetorotational and Rayleigh instability are not satisfied.
+" The condition of instability ts wpo*O and. hence. the shear-driven instability arises if B,#0."," The condition of instability is $\omega_{B \Omega} \neq 0$ and, hence, the shear-driven instability arises if $B_{s} \neq 0$."
+ The paper by Bonanno Urpin (2007) considers the dispersion relation for any wavevectors. but does so in a conductivity limit 77—0.," The paper by Bonanno Urpin (2007) considers the dispersion relation for any wavevectors, but does so in a high-conductivity limit $\eta \rightarrow 0$ ."
+ Eq. (, Eq. (
+"23) can be simplified in many cases of interest if we take into account that. likely. the s-component of the magnetic field in disks 1s greater than the z-component. B,>>B-.","23) can be simplified in many cases of interest if we take into account that, likely, the $s$ -component of the magnetic field in disks is greater than the $z$ -component, $B_s \gg B_z$."
+" For perturbations with k-B=0. we have k,=-k-B-/B,."," For perturbations with $\vec{k} \cdot 
+\vec{B} = 0$, we have $k_s = - k_z B_z/ B_s$."
+ Then. Substituting HR| into Eq. a(," Then, Substituting $\mu \approx 1$ into Eq. ("
+2ONG can transform it into Two roots of this equation deseribe the inertial waves which can be unstable only if the Rayleigh criterion is satisfied.,"23), we can transform it into Two roots of this equation describe the inertial waves which can be unstable only if the Rayleigh criterion is satisfied."
+ Other four modes. are deseribed by the dispersion relation This equation deseribes fast and slow magnetoacoustic waves. and we consider the stability of these modes.," Other four modes are described by the dispersion relation This equation describes fast and slow magnetoacoustic waves, and we consider the stability of these modes."
+ Note that simplification (26) is made for mathematical convenience rather than for physical relevance., Note that simplification (26) is made for mathematical convenience rather than for physical relevance.
+ In. fact. the dispersion relation (28) for fast and slow magnetosonie waves will change litle if wo#1.," In fact, the dispersion relation (28) for fast and slow magnetosonic waves will change little if $\mu \neq 1$."
+ This particularly concerns the modes with jo]>Q because the dispersion equation for them will differ from Eq. (, This particularly concerns the modes with $|\sigma| > \Omega$ because the dispersion equation for them will differ from Eq. (
+28) only by terms of the order of (ji—Lafor in coefficients.,28) only by terms of the order of $(\mu -1) \kappa^2/ \sigma^2$ in coefficients.
+ Therefore. our results can be applied with a sufficient accuracy also for the perturbations with 4i#I.," Therefore, our results can be applied with a sufficient accuracy also for the perturbations with $\mu \neq 1$."
+ The Hurwitz theorem states that an equation of the fourth order. has at least one root with a positive real part (unstable mode) if one of the following inequalities is fulfilled (see Aleksandrov et al.," The Hurwitz theorem states that an equation of the fourth order, has at least one root with a positive real part (unstable mode) if one of the following inequalities is fulfilled (see Aleksandrov et al."
+ 1985)., 1985).
+ Conditions (30) are never satisfied 1n. protostellar disks., Conditions (30) are never satisfied in protostellar disks.
+ Therefore. the instability arises rf one of the inequalities (31) or (32) are satisfied.," Therefore, the instability arises if one of the inequalities (31) or (32) are satisfied."
+ These inequalities can be rewritten for Eq. (, These inequalities can be rewritten for Eq. (
+"28) as Generally. the HOOexpressionsOj on the2, Lh.s.","28) as Generally, the expressions on the l.h.s."
+ of both these inequalities can have a positive or negative sign. depending on the value of «go.," of both these inequalities can have a positive or negative sign, depending on the value of $\omega_{B \Omega}$."
+ Therefore. Eq. (," Therefore, Eq. ("
+33) and (34) impose restrictions on the rate of differential rotation which should be greater than some critical value.,33) and (34) impose restrictions on the rate of differential rotation which should be greater than some critical value.
+ Note that. if the azimuthal field is generated winding up the radial field. then διο>0 and. hence. o;byBC;> Oas well.," Note that, if the azimuthal field is generated by winding up the radial field, then $B_s B_{\varphi} \Omega' >0$ and, hence, $\omega_{B \Omega}^3 > 0$ as well."
+ Therefore. we consider only this case.," Therefore, we consider only this case."
+Denoting 6=Casca/ 03. We have from Eq. (,"Denoting $\delta = c_{As} 
+c_{A \varphi}/c_{A}^2$ , we have from Eq. ("
+33 Thiscondition can ie satisfied only if the differential rotation is sufficiently strong.,33) Thiscondition can be satisfied only if the differential rotation is sufficiently strong.
+ Since Q’<0 in astrophysical disks. we," Since $\Omega' < 0$ in astrophysical disks, we"
+"For finite quark mass mg. F(r.T) remains finite lor r—x. since the ""string. between the two color charges ‘breaks’ when the corresponding potential enerey becomes equal to the mass AM, of the lowest hadron: bevond this point. it becomes energetically more favourable to produce an additional hadron.","For finite quark mass $m_q$ , $F(r,T)$ remains finite for $r \to \infty$, since the `string' between the two color charges `breaks' when the corresponding potential energy becomes equal to the mass $M_h$ of the lowest hadron; beyond this point, it becomes energetically more favourable to produce an additional hadron."
+" Hence now £ no longer vanishes in the confined phase. but only becomes exponentially small there. L(T)——Al,/Th:: here M, is a tvpical hadron mass. of the order of 0.5 to 1.0 GeV. so that al 7.c170 MeV. L~107. rather than zero."," Hence now $L$ no longer vanishes in the confined phase, but only becomes exponentially small there, L(T); here $M_h$ is a typical hadron mass, of the order of 0.5 to 1.0 GeV, so that at $T_c \simeq 170$ MeV, $L \sim 10^{-2}$, rather than zero."
+ Deconfinement is thus indeed much like the transition. for which the order parameter. the conductivity (7). also does not really vanish for T'>0. but with o(7)~exp{-AL/T} is only exponentially small. since thermal ionisation (with ionisation energv AZ) produces a small number of unbound electrons even in (he insulator phase.," Deconfinement is thus indeed much like the insulator-conductor transition, for which the order parameter, the conductivity $\sigma(T)$, also does not really vanish for $T>0$, but with $\sigma(T) \sim \exp\{-\Delta E/T\}$ is only exponentially small, since thermal ionisation (with ionisation energy $\Delta E$ ) produces a small number of unbound electrons even in the insulator phase."
+ rel?ja illustrates (heschematically (he behavior of LCD) and of the corresponding susceptibility VLUD)e(L2)=(L7. as obtained in finite temperature lattice studies [2022].. for the case of two flavors of lisht quarks.," \\ref{2_4}$ $~\!$ a illustrates theschematically the behavior of $L(T)$ and of the corresponding susceptibility $\x_L(T) \sim \langle L^2 \rangle - \langle L \rangle^2$, as obtained in finite temperature lattice studies \cite{K&L,cheng75,
+cheng77}, for the case of two flavors of light quarks."
+ We note that L(7) undergoes the expected sudden increase from a small confinement to a much larger deconfinement value., We note that $L(T)$ undergoes the expected sudden increase from a small confinement to a much larger deconfinement value.
+ The sharp peak of νε} defines «quite well a transition temperature 7). which we shall shortly specify in physical units.," The sharp peak of $\chi_L(T)$ defines quite well a transition temperature $T_L$, which we shall shortly specify in physical units."
+ The next quantity to consider is (he effective quark mass: it is measured by (he expectation value of the corresponding term in the Lagrangian. (c0)(T).," The next quantity to consider is the effective quark mass; it is measured by the expectation value of the corresponding term in the Lagrangian, $\langle {\bar \psi} \psi \rangle(T)$."
+ In the limit of vanishing current quark mass. (he Lagraneian becomes chirally νοτας and (σὺ(1) the corresponding order parameter.," In the limit of vanishing current quark mass, the Lagrangian becomes chirally symmetric and $\langle {\bar \psi} \psi \rangle(T)$ the corresponding order parameter."
+" In the confined phase. with effective constituent. quark masses AL,c0.3 GeV. this chiral svuumetry is spontaneously broken. while in the deconfinecl phase. athieh enough temperature. we expect its restoration."," In the confined phase, with effective constituent quark masses $M_q \simeq 0.3$ GeV, this chiral symmetry is spontaneously broken, while in the deconfined phase, athigh enough temperature, we expect its restoration."
+" Hence now (cc)(T) constitutes a genuine order parameter. [mite for 7<7), ancl vanishing for 7> T),. as shown in relchi.."," Hence now $\langle {\bar \psi} \psi \rangle(T)$ constitutes a genuine order parameter, finite for $T< T_m$ and vanishing for $T\geq T_m$ , as shown in \\ref{chi}. ."
+Millisecond (recycled) radio pulsars are distinguished from ordinary pulsars by their very short and stable periods. P< ms. P—102!107? s sl,"Millisecond (recycled) radio pulsars are distinguished from ordinary pulsars by their very short and stable periods, $P\la 10$ ms, $\dot{P}\sim 10^{-21}-10^{-19}$ s $^{-1}$."
+ [t is generally accepted that they are very old objects. with spin-down ages P/2P~10°—10? yr and low surface magnetic fields B(PP)!v40—10'? G (eg. Taylor. Manchester. Lyne 1993).," It is generally accepted that they are very old objects, with spin-down ages $\tau=P/2\dot{P}\sim 10^9-10^{10}$ yr and low surface magnetic fields $B\propto (P\dot{P})^{1/2}\sim 10^8-10^{10}$ G (e.g., Taylor, Manchester, Lyne 1993)."
+ Similar to ordinary pulsars. a millisecond pulsar can emit nonthermal X-rays from its magnetosphere. with a hard power-law spectrum and sharp pulsations.," Similar to ordinary pulsars, a millisecond pulsar can emit nonthermal X-rays from its magnetosphere, with a hard power-law spectrum and sharp pulsations."
+ In addition to this nonthermal radiation. thermal X-rays can be emitted from the neutron star (NS) surface. provided the surface is hot enough.," In addition to this nonthermal radiation, thermal X-rays can be emitted from the neutron star (NS) surface, provided the surface is hot enough."
+ According to the models of NS thermal evolution (albeit rather uncertain at these old ages). recycled pulsars are too cold (surface temperature Τ<0.1 MK — see. e. g.. Tsuruta 1998) to be detectable in X-rays.," According to the models of NS thermal evolution (albeit rather uncertain at these old ages), recycled pulsars are too cold (surface temperature $T\la 0.1$ MK — see, e. g., Tsuruta 1998) to be detectable in X-rays."
+ However. their polar caps can be heated up to X-ray temperatures by relativistic particles impinging onto the magnetic poles from the acceleration zones in the magnetosphere.," However, their polar caps can be heated up to X-ray temperatures by relativistic particles impinging onto the magnetic poles from the acceleration zones in the magnetosphere."
+ The radio pulsar models (e.g.. Cheng Ruderman 1980: Arons 1981: Michel 1991; Beskin. Gurevich. Istomin 1993) predict polar cap radii Ay.~QzRPol? (where R~10 km ts the NS radius). i.e... Ry;— 1-5 km for millisecond pulsars. although different models predict quite different polar cap temperatures. in the range of 1-10 MK.," The radio pulsar models (e.g., Cheng Ruderman 1980; Arons 1981; Michel 1991; Beskin, Gurevich, Istomin 1993) predict polar cap radii $\rpc\sim (2\pi R^3/Pc)^{1/2}$ (where $R\approx 10$ km is the NS radius), i.e., $\rpc\sim 1$ –5 km for millisecond pulsars, although different models predict quite different polar cap temperatures, in the range of 1–10 MK."
+ Detection of the polar cap thermal radiation would allow one to discriminate between various models of radio pulsars. study the properties of NS surface layers. and constrain the NS mass-to-radius ratio (Pavlov Zavlin 1997; Zavlin Pavlov 1998 [ZP98]).," Detection of the polar cap thermal radiation would allow one to discriminate between various models of radio pulsars, study the properties of NS surface layers, and constrain the NS mass-to-radius ratio (Pavlov Zavlin 1997; Zavlin Pavlov 1998 [ZP98])."
+ However. just as in the case of ordinary pulsars. this radiation is detectable only if it is not buried under stronger nonthermal radiation.," However, just as in the case of ordinary pulsars, this radiation is detectable only if it is not buried under stronger nonthermal radiation."
+ The current theoretical models are not elaborate enough to predict in which (if any) of millisecond pulsars the thermal component can be brighter than the nonthermal one (in particular. both the thermal and nonthermal luminosities are expected to increase with energy loss E. perhaps with different rates).," The current theoretical models are not elaborate enough to predict in which (if any) of millisecond pulsars the thermal component can be brighter than the nonthermal one (in particular, both the thermal and nonthermal luminosities are expected to increase with energy loss $\dot{E}$, perhaps with different rates)."
+ Therefore. we have to rely upon the analysis of X-ray observations to distinguish the thermal and nonthermal components.," Therefore, we have to rely upon the analysis of X-ray observations to distinguish the thermal and nonthermal components."
+ The X-ray observatoriesROSAT..ASCA.. and have detected 11 millisecond pulsars (nearly 1/3 of all X-ray-detected rotation-powered pulsars — see Becker Pavlov 2001 for a recent review).," The X-ray observatories, and have detected 11 millisecond pulsars (nearly 1/3 of all X-ray-detected rotation-powered pulsars — see Becker Pavlov 2001 for a recent review)."
+ Five of these pulsars are identified in X-rays only by positional comeidence with the radio pulsars and. due to the low number of recorded counts. provide only crude flux estimates.," Five of these pulsars are identified in X-rays only by positional coincidence with the radio pulsars and, due to the low number of recorded counts, provide only crude flux estimates."
+ The radiation from 3 pulsars — B1821-24 (Saito et al., The radiation from 3 pulsars — B1821--24 (Saito et al.
+ 1997). B1937+21 (Takahashi et al.," 1997), B1937+21 (Takahashi et al."
+ 2001). and JO218+4232 (Mineo et al.," 2001), and J0218+4232 (Mineo et al."
+ 2000) — is clearly nonthermal: their power-law spectra. detected with and up to energies of 5-10 keV. are very hard. with photon indices ~ |. and their pulse profiles show sharp. peaks.," 2000) — is clearly nonthermal: their power-law spectra, detected with and up to energies of 5–10 keV, are very hard, with photon indices $\gamma\sim 1$ , and their pulse profiles show sharp peaks."
+ Interestingly. these 3 pulsars are characterized by particularly large E values. E=(2-20)«10? erg s!. and their magnetic fields at the light cylinder. By=ΒΚΔΙΣ~10° G. are close to that of the Crab pulsar.," Interestingly, these 3 pulsars are characterized by particularly large $\dot{E}$ values, $\dot{E}=(2-20)\times 10^{35}$ erg $^{-1}$, and their magnetic fields at the light cylinder, $B_{\rm lc}=B(R/R_{\rm lc})^3\sim 10^6$ G, are close to that of the Crab pulsar."
+ The case for the other 3 pulsars — J0437-4715 (Becker Trümmper 1993. 1999 [BT93. BT99]|: ZP98). J2124—3358 (BT99). and JO030+0451 (Becker et al.," The case for the other 3 pulsars — J0437–4715 (Becker Trümmper 1993, 1999 [BT93, BT99]; ZP98), J2124–3358 (BT99), and J0030+0451 (Becker et al."
+ 2000) — is less certain., 2000) — is less certain.
+ These pulsars show broad peaks of X-ray pulsations. but it does not necessarily mean that their radiation is thermal because broad peaks can be produced by nonthermal emission at some viewing angles.," These pulsars show broad peaks of X-ray pulsations, but it does not necessarily mean that their radiation is thermal because broad peaks can be produced by nonthermal emission at some viewing angles."
+ High-quality spectra have been recorded for the brightest of these pulsars. J0437—4715. but their interpretation has been controversial — e.g.. ZP98 suggest that the radiation detected with and can be interpreted as thermal radiation from hot polar caps. whereas BT99 argue that the radiation is nonthermal (see $22).," High-quality spectra have been recorded for the brightest of these pulsars, J0437–4715, but their interpretation has been controversial — e.g., ZP98 suggest that the radiation detected with and can be interpreted as thermal radiation from hot polar caps, whereas BT99 argue that the radiation is nonthermal (see 2)."
+ To resolve this controversy. the pulsar needed to be observed at energies abovethe soft and bands (EmS>2 keV). and with high spatial resolution to avoid," To resolve this controversy, the pulsar needed to be observed at energies abovethe soft and bands $E\ga 2$ keV), and with high spatial resolution to avoid"
+a phase transition (22222)..,"a phase transition \citep{Baym:1995fk,Martin:1995su,Hindmarsh:1997tj,Boyanovsky:2002wa,Kahniashvili:2009qi}."
+ Fields can also be produced by the production of non-linear vorticity from linear density perturbations(???????).. but the impact these have on the CMB ts complicated by their evolving. non-trivial nature.," Fields can also be produced by the production of non-linear vorticity from linear density \citep{Gopal:2004ut,Matarrese:2004kq,Takahashi:2005nd,Ichiki:2006cd,Siegel:2006px,Kobayashi:2007wd,Maeda:2008dv}, but the impact these have on the CMB is complicated by their evolving, non-trivial nature."
+ The impact primordial magnetic fields have on the CMB and its anisotropies have been well-studied. with (22222222222922222?) being some instructive examples.," The impact primordial magnetic fields have on the CMB and its anisotropies have been well-studied, with \citep{Barrow:1997mj,Subramanian:1998fn,Durrer:1998ya,Koh:2000qw,Kahniashvili:2000vm,Mack:2001gc,Clarkson:2002dd,Lewis:2004ef,Kahniashvili:2006hy,Kahniashvili:2008hx,Yamazaki:2008gr,Finelli:2008xh,Paoletti:2008ck,Bonvin:2010nr,Giovannini:2009fu,Yamazaki:2010nf,Paoletti:2010rx,Kahniashvili:2010wm} being some instructive examples."
+ While older literature tended to assume a homogeneous background component with an inhomogeneous perturbation. more recent work has typically focused on tangled configurations without a background component and | assume this throughout.," While older literature tended to assume a homogeneous background component with an inhomogeneous perturbation, more recent work has typically focused on tangled configurations without a background component and I assume this throughout."
+ These studies fairly consistently suggest that the field is constrained to be of at most nano-Gauss in magnitude., These studies fairly consistently suggest that the field is constrained to be of at most nano-Gauss in magnitude.
+ The spectral index is restricted to be approximately scale-invariant (??).. with limits growing extremely tight for a primordial magnetic field with index far from scale-invarianee (2)..," The spectral index is restricted to be approximately scale-invariant \citep{Yamazaki:2010nf,Paoletti:2010rx}, with limits growing extremely tight for a primordial magnetic field with index far from scale-invariance \citep{Caprini:2001nb}."
+ A large-scale homogeneous field also introduces characteristic correlations between multipole moments with A/€[-2.0.2] and Am€(0.x1.£2] which vanish in the standard scenario (?)..," A large-scale homogeneous field also introduces characteristic correlations between multipole moments with $\Delta l\in\{-2,0,2\}$ and $\Delta m\in\{0,\pm 1,\pm2\}$ which vanish in the standard scenario \citep{Kahniashvili:2008sh}."
+ However. the magnetic 2-point signal is overwhelmed on large-scales by the standard perturbations. with the B-mode polarisation being perhaps the most realistic option if we are to detect it directly.," However, the magnetic 2-point signal is overwhelmed on large-scales by the standard perturbations, with the $B$ -mode polarisation being perhaps the most realistic option if we are to detect it directly."
+ The increasing accuracy of measurements of the CMB non-Gaussianity provides an alternative., The increasing accuracy of measurements of the CMB non-Gaussianity provides an alternative.
+ The stress tensor of a magnetic field is implying that the statistics induced on matter perturbations are intrinsically non-Gaussian. regardless of the nature of the underlying magnetic field.," The stress tensor of a magnetic field is non-linear, implying that the statistics induced on matter perturbations are intrinsically non-Gaussian, regardless of the nature of the underlying magnetic field."
+ Since the standard scenario contains relatively few sources of primordial it is possible that a magnetic signal is dominant.," Since the standard scenario contains relatively few sources of primordial non-Gaussianity, it is possible that a magnetic signal is dominant."
+ Viewed another way. predicted signals from a magnetic field are likely to be of a characteristic nature. and must be found in and cleaned from the CMB data before any conclusions on early-universe physics can be made.," Viewed another way, predicted signals from a magnetic field are likely to be of a characteristic nature, and must be found in and cleaned from the CMB data before any conclusions on early-universe physics can be made."
+ Aspects of the three-point moments have been studied in a series of papers in the last few years (22222222) ," Aspects of the three-point moments have been studied in a series of papers in the last few years \citep{Brown:2005kr,Brown:2006wv,Seshadri:2009sy,Caprini:2009vk,Trivedi:2010gi,Cai:2010uw,Shiraishi:2010yk,Kahniashvili:2010us}. ."
+A bispectrum is set by three wavevectors. which we denote with k. p and q.," A bispectrum is set by three wavevectors, which we denote with $\mathbf{k}$, $\mathbf{p}$ and $\mathbf{q}$."
+ Since to retain statistical isotropy these must form a closed triangle. this geometry can equivalently be expressed with the scalars &.7.o. where p=rk and ὁ is the angle between p and q.," Since to retain statistical isotropy these must form a closed triangle, this geometry can equivalently be expressed with the scalars $k,r,\phi$, where $p=rk$ and $\phi$ is the angle between $\mathbf{p}$ and $\mathbf{q}$ ."
+ Employing these variables the bispectrum geometry can be written as a foliation of planes of constant rand for each constant angle @ we then have a one-dimensional line through the bispectrum which in broad terms is expected to act ina similar manner to the power spectra., Employing these variables the bispectrum geometry can be written as a foliation of planes of constant $r$ and for each constant angle $\phi$ we then have a one-dimensional line through the bispectrum which in broad terms is expected to act in a similar manner to the power spectra.
+" The magnetic bispectra studied thus far have typically been along only three such lines. all in the 7=1 plane — the ""colinear case where k=p4/2 and so 6=0 (???.hereafterBCOS.BO6andCFPRO9).. the “equilateral” case wherek=p qand so 6=27/3 (???.hereafterSS09.CFPRO9andTSS10).. and the ""local"" or degenerate case where kοὐ=p. q=0 and so @=x (SS09. CFPRO9. TSS10)."," The magnetic bispectra studied thus far have typically been along only three such lines, all in the $r=1$ plane – the “colinear” case where $k=p=q/2$ and so $\phi=0$ \citep[hereafter BC05, B06 and CFPR09]{Brown:2005kr,Brown:2006wv,Caprini:2009vk}, the “equilateral” case where $k=p=q$ and so $\phi=2\pi/3$ \citep[hereafter SS09, CFPR09 and TSS10]{Seshadri:2009sy,Caprini:2009vk,Trivedi:2010gi}, and the “local” or degenerate case where $k\approx p$, $q\approx 0$ and so $\phi\approx\pi$ (SS09, CFPR09, TSS10)."
+ TSS10 also considered configurations where @=0 but kzp., TSS10 also considered configurations where $\phi=0$ but $k\neq p$.
+ The recent studies have expanded the previous results considerably., The recent studies have expanded the previous results considerably.
+ SS09 considered the equilateral and degenerate lines of the bispectrum of the magnetic energy density for nearly scale-invariant magnetic fields. concluding that the degenerate line provides the greatest contribution to the integral and employing an approximation to this dominant term to estimate the CMB signal.," SS09 considered the equilateral and degenerate lines of the bispectrum of the magnetic energy density for nearly scale-invariant magnetic fields, concluding that the degenerate line provides the greatest contribution to the integral and employing an approximation to this dominant term to estimate the CMB signal."
+ Likewise. CFPRO9 considered the bispectrum of the energy density and considered the colinear. equilateral and degenerate lines.," Likewise, CFPR09 considered the bispectrum of the energy density and considered the colinear, equilateral and degenerate lines."
+ The authors generally relied on approximations that neglect angular terms in the integrations. or apply only on large scales.," The authors generally relied on approximations that neglect angular terms in the integrations, or apply only on large scales."
+ Doing so recovers the scaling behaviour of the bispectrum at the expense of an accurate calculation of the relative amplitudes between lines., Doing so recovers the scaling behaviour of the bispectrum at the expense of an accurate calculation of the relative amplitudes between lines.
+" Since the degenerate line was found to diverge as g-""~ as g—O this term ts likely to dominate.", Since the degenerate line was found to diverge as $q^{2n+3}$ as $q\rightarrow 0$ this term is likely to dominate.
+ CFPRO9 also present exact solutions for the colinear case for both a causal field and a field relatively close to scale-invariance. which enable them to test their approximations.," CFPR09 also present exact solutions for the colinear case for both a causal field and a field relatively close to scale-invariance, which enable them to test their approximations."
+ The approximations are certainly reasonable. but not ideal.," The approximations are certainly reasonable, but not ideal."
+ In particular. since the bispectra are not positive-definite it is unclear whether there are strong cancellations to the degenerate line arising from other parts of the bispectrum.," In particular, since the bispectra are not positive-definite it is unclear whether there are strong cancellations to the degenerate line arising from other parts of the bispectrum."
+ ? employed the approximations of SS09 and CFPRO9 and extended the treatment to full transfer functions., \citet{Cai:2010uw} employed the approximations of SS09 and CFPR09 and extended the treatment to full transfer functions.
+ Morerecently. TSS10 considered the bispectrum of theanisotropic pressure of a near scale-invariant magnetic field.," Morerecently, TSS10 considered the bispectrum of theanisotropic pressure of a near scale-invariant magnetic field."
+ Unlike the previous papers they evaluated the bispectrum along the degenerate line in full. without neglecting any angular," Unlike the previous papers they evaluated the bispectrum along the degenerate line in full, without neglecting any angular"
+Finally we look at the RV data themselves.,"Finally, we look at the RV data themselves."
+ The calculation of the formal uncertainties assumes that the RV errors scale linearly with the inverse of the spectrum SNR (Bouchyetal.2001)., The calculation of the formal uncertainties assumes that the RV errors scale linearly with the inverse of the spectrum SNR \citep{bou01}.
+". On the other hand, we suspect that the dependency of the errors on the SNR is stronger than linear."," On the other hand, we suspect that the dependency of the errors on the SNR is stronger than linear."
+" To test this, we divided the RV residuals shown in Fig."," To test this, we divided the RV residuals shown in Fig."
+ 3 by the SNR of the corresponding spectra., \ref{fig3} by the SNR of the corresponding spectra.
+" The scatter of the lower-SNR data points is significantly larger than that of the high-SNR points, by a much larger amount than indicated by the formal error bars."," The scatter of the lower-SNR data points is significantly larger than that of the high-SNR points, by a much larger amount than indicated by the formal error bars."
+ The same holds for the residuals relative to the mean., The same holds for the residuals relative to the mean.
+ This observation clearly indicates that the dependency of the uncertainties on SNR is underestimated., This observation clearly indicates that the dependency of the uncertainties on SNR is underestimated.
+" It can be modelled by assuming the presence of additional instrumental uncertainties with a steep dependence on SNR, at the ~5 mmss! level near the median SNR, increasing to about ~10 ! at the low end of the SNR range."," It can be modelled by assuming the presence of additional instrumental uncertainties with a steep dependence on SNR, at the $\sim 5$ $^{-1}$ level near the median SNR, increasing to about $\sim 10$ $^{-1}$ at the low end of the SNR range."
+ The standard way of searching for the signature of planetary orbits in RV data is to use a Lomb-Scargle or generalized periodogram (Horne&Baliunas1986;PressRybicki1989;Zechmeister&Kürster 2009).," The standard way of searching for the signature of planetary orbits in RV data is to use a Lomb-Scargle or generalized periodogram \citep{Hor86,Pre86,Zec09}."
+. The periodogram of the HARPS RV data for CoRoT-7 is highly complex., The periodogram of the HARPS RV data for CoRoT-7 is highly complex.
+" In the course of their pre-whitening analysis, Q09 identify no less than eleven peaks, all of them highly ‘significant’ in the sense that they correspond to low formal false alarm probabilities."," In the course of their pre-whitening analysis, Q09 identify no less than eleven peaks, all of them highly `significant' in the sense that they correspond to low formal false alarm probabilities."
+" However, on should bear in mind that the false alarm probability expresses the probability that a given peak is due to Gaussian white noise, but neither activity nor instrumental noise are expected to be white or Gaussian."," However, on should bear in mind that the false alarm probability expresses the probability that a given peak is due to Gaussian white noise, but neither activity nor instrumental noise are expected to be white or Gaussian."
+" The very irregular sampling of the data also implies that one or more of the peaks may arise from signal at a single, apparently unrelated frequency."," The very irregular sampling of the data also implies that one or more of the peaks may arise from signal at a single, apparently unrelated frequency."
+" Nonetheless, the main peaks in the RV periodogram are clearly related to the signal from activity, being near the rotation period, its harmonics, and their one-day aliases."," Nonetheless, the main peaks in the RV periodogram are clearly related to the signal from activity, being near the rotation period, its harmonics, and their one-day aliases."
+" There is also a peak corresponding to the period of the transit signal detected in the CoRoT photometry, P;=0.854 dd. Fig."," There is also a peak corresponding to the period of the transit signal detected in the CoRoT photometry, $P_b=0.854$ d. Fig."
+" 4 shows the semi-amplitude of the best-fit sinusoid at the period and phase of the photometric transit, as a function of SNR threshold, as measurements derived from low-SNR spectra are progressively discarded, starting with the lowest SNR."," \ref{fig6} shows the semi-amplitude of the best-fit sinusoid at the period and phase of the photometric transit, as a function of SNR threshold, as measurements derived from low-SNR spectra are progressively discarded, starting with the lowest SNR."
+" For the most stringent threshold (SNR>1.15SNRmea), about one third of the measurements remain."," For the most stringent threshold ${\rm SNR}>1.15\,{\rm SNR}_{\rm
+med}$ ), about one third of the measurements remain."
+" We performed this calculation with both the raw RV, and the residuals from our activity models."," We performed this calculation with both the raw RV, and the residuals from our activity models."
+" The stellar rotation and planetary orbital frequencies are widely separated, but the latter is close to the one-day alias of the third harmonic of the former."," The stellar rotation and planetary orbital frequencies are widely separated, but the latter is close to the one-day alias of the third harmonic of the former."
+" As a result, it is not clear a priori whether correcting for the activity signal improves the semi-amplitude measurements, or on the contrary adds noise to them."," As a result, it is not clear a priori whether correcting for the activity signal improves the semi-amplitude measurements, or on the contrary adds noise to them."
+" Fortunately, the two methods give very similar results."," Fortunately, the two methods give very similar results."
+" In both cases, the measured orbital semi-amplitude depends strongly on the SNR threshold: including lower-SNR measurements favours a higher value."," In both cases, the measured orbital semi-amplitude depends strongly on the SNR threshold: including lower-SNR measurements favours a higher value."
+" Low-SNR measurements are more likely to be outliers (as their formal uncertainties are underestimated), and would favour a higher amplitude for all fitted features: the higher number of measurements is offset by their poorer quality."," Low-SNR measurements are more likely to be outliers (as their formal uncertainties are underestimated), and would favour a higher amplitude for all fitted features: the higher number of measurements is offset by their poorer quality."
+" It is therefore not clear whether the most reliable value of Ky is the one derived from all the RV measurements, or from only the best third or half."," It is therefore not clear whether the most reliable value of $K_b$ is the one derived from all the RV measurements, or from only the best third or half."
+ The monotonic trend in Ky versus SNR threshold suggests that SNR-dependent effects play a large role in the detected amplitude., The monotonic trend in $K_b$ versus SNR threshold suggests that SNR-dependent effects play a large role in the detected amplitude.
+" Figure 5 shows the radial-velocity data corrected by one of our variability models and phased to the transit signal, together with the best-fit Keplerian orbit with and without an SNR cut."," Figure \ref{phasedvr}
+ shows the radial-velocity data corrected by one of our variability models and phased to the transit signal, together with the best-fit Keplerian orbit with and without an SNR cut."
+" 'To evaluate the effect of the steeper rise in the total uncertainty at the lower end of the SNR range, we re-calculate the constraints on K in the following way: we add a new term to the radial-velocity uncertainties, with a quadratic rather than linear dependence on the inverse of the signal-to-noise ratio."," To evaluate the effect of the steeper rise in the total uncertainty at the lower end of the SNR range, we re-calculate the constraints on $K$ in the following way: we add a new term to the radial-velocity uncertainties, with a quadratic rather than linear dependence on the inverse of the signal-to-noise ratio."
+ We set the magnitude of this term so that the reduced x? of the residuals of the, We set the magnitude of this term so that the reduced $\chi^2$ of the residuals of the
+among different observables (MHDO3: McHardy et al.,among different observables (MHD03; McHardy et al.
+ 2006). supporting the notion that AGN are indeed scaled-up galactic black holes.," 2006), supporting the notion that AGN are indeed scaled-up galactic black holes."
+ However. physical models for the disc-jet coupling in BHXRB based on the observed correlations between radio and X-ray luminosity (Fender.Gallo&Jonker2003:Kórding.Fender&Migliari2006) all face a large uncertainty due to the lack of reliable measurements of the jet kinetic power.," However, physical models for the disc-jet coupling in BHXRB based on the observed correlations between radio and X-ray luminosity \cite{fender:03,koerding:06b} all face a large uncertainty due to the lack of reliable measurements of the jet kinetic power."
+ In this context. it is interesting to include in Figure 2. the only GBH for which a measurement of the kinetic output has been made. Cyg X-I (Galloetal.2005)... which turns out to be consistent with the relationship derived from the AGN sample.," In this context, it is interesting to include in Figure \ref{fig:lkinledd} the only GBH for which a measurement of the kinetic output has been made, Cyg X-1 \cite{gallo:05}, which turns out to be consistent with the relationship derived from the AGN sample."
+ Clearly. systematic efforts to estimate kinetic power of BHXRB in the low/hard state are needed in order to assess their similarity with radio mode AGN.," Clearly, systematic efforts to estimate kinetic power of BHXRB in the low/hard state are needed in order to assess their similarity with radio mode AGN."
+ If the above correlation (3)) directly reveals fundamental physical properties of jet-producing AGN of lowpower. it still shows a non- intrinsic scatter.," If the above correlation \ref{eq:lklambda}) ) directly reveals fundamental physical properties of jet-producing AGN of lowpower, it still shows a non-negligible intrinsic scatter."
+ On the other hand. one should expect a more direet relationship between the nuclear radio core emission and the larger scale Kinetic power. as both originate from the jet.," On the other hand, one should expect a more direct relationship between the nuclear radio core emission and the larger scale kinetic power, as both originate from the jet."
+ All theoretical models for AGN flat-spectrum compact jet cores (Blandford&Kónigl1979:FalekeBiermann1996:HeinzSunyaev2003) predict a dependence of the radio luminosity on the jet power in the form Ly-x ," All theoretical models for AGN flat-spectrum compact jet cores \cite{blandford:79,falcke:96,heinz:03} a dependence of the radio luminosity on the jet power in the form $L_{\rm R} \propto
+L_{\rm kin}^{17/12}$."
+The current sample provides by far the best opportunity to test Lus.these predictions., The current sample provides by far the best opportunity to test these predictions.
+ A Kendalls tau correlation test reveals that the kinetic power is correlated with theobserved radio core luminosity Lion... with Pun=9210 (see the empty circles in Figure +)).," A Kendall's tau correlation test reveals that the kinetic power is correlated with the radio core luminosity $L_{\rm
+R,obs}$, with $P_{\rm null}= 9.2\times 10^{-5}$ (see the empty circles in Figure \ref{fig:pboth}) )."
+ We have fitted the data with a linear relationship:Ly. once again making use of a symmetric regression algorithm that takes into account errors on both variables.," We have fitted the data with a linear relationship:, once again making use of a symmetric regression algorithm that takes into account errors on both variables."
+ We obtain clan.=(22.1$3.5). Ba=(0.54+0.09). with a large intrinsic scatter of y=0.41.," We obtain $A_{\rm obs}= (22.1 \pm 3.5)$, $B_{\rm obs}=(0.54 \pm 0.09)$, with a large intrinsic scatter of $\sigma=0.47$."
+ Such a correlation. however. must be at some level biased by relativistic Doppler boosting of the radio emission in the relativistic jets.," Such a correlation, however, must be at some level biased by relativistic Doppler boosting of the radio emission in the relativistic jets."
+" An alternative way to proceed would be to use estimators of the nuclear radio core luminosity which are less affected by relativistic beaming (Heinz&Grimm2005).. as. for example. the multivariate relation between BH mass. radio core and hard X-ray luminosity. the so-called ""fundamental plane’ (FP) of active black holes (MHDO3)."," An alternative way to proceed would be to use estimators of the nuclear radio core luminosity which are less affected by relativistic beaming \cite{heinz:05}, as, for example, the multivariate relation between BH mass, radio core and hard X-ray luminosity, the so-called `fundamental plane' (FP) of active black holes (MHD03)."
+ Recent analysis of this correlation (Heinz Merloni 2004: Kórrding. Faleke Corbel 2006 [KFCO6]: lerloni et al.," Recent analysis of this correlation (Heinz Merloni 2004; Körrding, Falcke Corbel 2006 [KFC06]; Merloni et al."
+ 2006) have shown that both Doppler boosting and sample selection play a crucial role in the exact determination of he intrinsic correlation coefficients of the FP. which also need to be accounted for.," 2006) have shown that both Doppler boosting and sample selection play a crucial role in the exact determination of the intrinsic correlation coefficients of the FP, which also need to be accounted for."
+ In the Appendix. we diseuss in detail a possible way © overcome such a bias with the aid of a Monte Carlo simulation of the samples used to derive the FP relation.," In the Appendix, we discuss in detail a possible way to overcome such a bias with the aid of a Monte Carlo simulation of the samples used to derive the FP relation."
+" That study allows us to estimate statistically the intrinsic (un-boosted) radio core uminosity of the AGN jets as a function of their (mean) Lorentz actor. Li, in a way that ean be approximated by the following expression: where Lip ds the intrinsic (un-boosted) radio core luminosity of the jet at 5 GHz. Lx the nuclear 2-10 keV intrinsic (un-absorbed) luminosity and Ly, the mean Lorentz factor of the jets."," That study allows us to estimate statistically the intrinsic (un-boosted) radio core luminosity of the AGN jets as a function of their (mean) Lorentz factor, $\Gamma_{\rm m}$ in a way that can be approximated by the following expression: where $ L_{\rm R,FP}$ is the intrinsic (un-boosted) radio core luminosity of the jet at 5 GHz, $L_{\rm X}$ the nuclear 2-10 keV intrinsic (un-absorbed) luminosity and $\Gamma_{\rm m}$ the mean Lorentz factor of the jets."
+ In what follows. we will adopt the specitic version of the FP relation derived from a sample of low luminosity AGN only. i.e. free from the bias introduced by the inclusion of bright. radiatively efficient AGN or QSOs (see discussion in KFCO6).," In what follows, we will adopt the specific version of the FP relation derived from a sample of low luminosity AGN only, i.e. free from the bias introduced by the inclusion of bright, radiatively efficient AGN or QSOs (see discussion in KFC06)."
+ For that. the correlation coefficients are ax=0.71. £n;=0.62. slightly different (but only at the 1-7 level) from those found in MHDO3.," For that, the correlation coefficients are $\xi_{\rm RX}=0.71$, $\xi_{\rm RM}=0.62$, slightly different (but only at the $\sigma$ level) from those found in MHD03."
+ Given Eq. (4)).," Given Eq. \ref{eq:fp_corr}) ),"
+" assuming a distribution of Lorentz factors for the AGN jets tor its mean. provided that the distribution is not too broad). we determine the ""true"" relationship between the Kinetic luminosity and radio core luminosityby fitting the 15 data points in our sample with the linear relationship The fitted values for the intrinsic slope. Di. as a function of Puy. are shown as a dot-dashed line in the bottom panel of Figure 3.."," assuming a distribution of Lorentz factors for the AGN jets (or its mean, provided that the distribution is not too broad), we determine the “true” relationship between the Kinetic luminosity and radio core luminosityby fitting the 15 data points in our sample with the linear relationship The fitted values for the intrinsic slope, $B_{\rm int}$, as a function of $\Gamma_{\rm m}$, are shown as a dot-dashed line in the bottom panel of Figure \ref{fig:bobs}."
+ From this we can see that the higher the mean Lorentz ‘actor of the jets. the steeper must the intrinsic correlation between Kinetic power and jet core luminosity be. and the larger the discrepancy with the measured slope. 4...=0.54+0.09 (solid ines in Figure 39). obtained using simply the observed radio core uminosity. without any attempt to correct for relativistic beaming.," From this we can see that the higher the mean Lorentz factor of the jets, the steeper must the intrinsic correlation between kinetic power and jet core luminosity be, and the larger the discrepancy with the measured slope, $B_{\rm
+obs}=0.54 \pm 0.09$ (solid lines in Figure \ref{fig:bobs}) ), obtained using simply the observed radio core luminosity, without any attempt to correct for relativistic beaming."
+ In fact. such a discrepancy between the intrinsic and the observed slopes of the Ly - Li; relation is indeed expected if the 15 sources of our sample harbor relativistic jet randomly oriented with respect o the line of sight?.," In fact, such a discrepancy between the intrinsic and the observed slopes of the $L_{\rm R}$ - $L_{\rm kin}$ relation is indeed expected if the 15 sources of our sample harbor relativistic jet randomly oriented with respect to the line of ."
+. In order to show this quantitatively. we have simulated (107 imes) the observed sample. assuming an underlying relationship," In order to show this quantitatively, we have simulated $10^4$ times) the sample, assuming an underlying relationship"
+"CH, mixing ratio with altitude with a scale height of ~20 km. due to photolysis.","$_4$ mixing ratio with altitude with a scale height of $\sim$ 20 km, due to photolysis."
+ We obtained the same column density as above. indicating a partial pressure of methane of 9.843.7 nbar. 1.8. a surface density of (1.940.7) x 107 οι”.," We obtained the same column density as above, indicating a partial pressure of methane of $\pm$ 3.7 nbar, i.e. a surface density of $\pm$ 0.7) x $^{12}$ $^{-3}$."
+" This appears to be 43. times larger than inferred from Voyager in 1989. adopting the CH, number densities of Herbert and Sandel (1991) and Strobel and Summers (1995) (4.7x10!! em"". within a factor 1.7. averaging ingress and egress)."," This appears to be $^{+5}_{-2.5}$ times larger than inferred from Voyager in 1989, adopting the $_4$ number densities of Herbert and Sandel (1991) and Strobel and Summers (1995) $\times$ $^{11}$ $^{-2}$, within a factor 1.7, averaging ingress and egress)."
+" An even larger enhancement factor (57°) is indicated 1f the Krasnopolsky and Cruikshank (1995) reanalysis of the Voyager UV data. giving CH; = 3.1z0.8x 10"" em? at the surface. is used."," An even larger enhancement factor $^{+6}_{-2}$ ) is indicated if the Krasnopolsky and Cruikshank (1995) reanalysis of the Voyager UV data, giving $_4$ = $\pm$ $\times$ $^{11}$ $^{-3}$ at the surface, is used."
+ Results are independent on the surface pressure. as collisional broadening is negligible.," Results are independent on the surface pressure, as collisional broadening is negligible."
+" They clearly demonstrate that the CH, partial pressure has increased in the last 20 years.", They clearly demonstrate that the $_4$ partial pressure has increased in the last 20 years.
+ The 2335-2365 nm part of the Triton spectrum (see excerpts in Fig., The 2335-2365 nm part of the Triton spectrum (see excerpts in Fig.
+ 3) shows the detection of 8 lines due to the CO(2-0) band (R2-R5. P2. P3. P5 and P5). providing the first detection of CO in its atmosphere.," 3) shows the detection of 8 lines due to the CO(2-0) band (R2-R5, P2, P3, P5 and P8), providing the first detection of CO in its atmosphere."
+ An accurate determination of the CO abundance is particularly difficult. as at infinite spectral resolution. these features are very narrow. saturated Doppler-shaped lines.," An accurate determination of the CO abundance is particularly difficult, as at infinite spectral resolution, these features are very narrow, saturated Doppler-shaped lines."
+ Nonetheless. assuming a vertically uniform CO distribution. and utilizing the whole set of CO lines (see Fig.," Nonetheless, assuming a vertically uniform CO distribution, and utilizing the whole set of CO lines (see Fig."
+ + on-line). we determine a CO column of 0.30 em-am. re. a CO partial pressure of 24 nbar. within a factor of 3.," 4 on-line), we determine a CO column of 0.30 cm-am, i.e. a CO partial pressure of 24 nbar, within a factor of 3."
+" The column density CO/CH, ratio is nominally ~3.75 (surface partial pressure ratio CO/CH, -2.5). with a factor of 4 uncertainty."," The column density $_4$ ratio is nominally $\sim$ 3.75 (surface partial pressure ratio $_4$ $\sim$ 2.5), with a factor of 4 uncertainty."
+ Deriving the CO/N> and CHj/N» mixing ratio is complicated by the fact that the surface pressure in 2009 is unknown., Deriving the $_2$ and $_4$ $_2$ mixing ratio is complicated by the fact that the surface pressure in 2009 is unknown.
+ Stellar occultation results (Olkin et al., Stellar occultation results (Olkin et al.
+" 1997, Sicardy et al."," 1997, Sicardy et al."
+ 1998. Elliotet al.," 1998, Elliot et al."
+ 1998. 20002) indicate that the pressure has been doubling in ~ 10 years from the 14 jibar value determined by Voyager in 1989 (Gurrola 1995).," 1998, 2000a) indicate that the pressure has been doubling in $\sim$ 10 years from the 14 $\mu$ bar value determined by Voyager in 1989 (Gurrola 1995)."
+ A reasonable assumption for 2009 is 40 jibar. providing CO/Ns ~ 6x 1077 and CH4/N> -2.4x 1077 at the surface. within factors of3 and 1.4 respectively.," A reasonable assumption for 2009 is 40 $\mu$ bar, providing $_2$ $\sim$ $\times$ $^{-4}$ and $_4$ $_2$ $\sim$ $\times$ $^{-4}$ at the surface, within factors of 3 and 1.4 respectively."
+ The CO abundance we determine is many times less than previous upper limits (Broadfoot et al., The CO abundance we determine is many times less than previous upper limits (Broadfoot et al.
+ 1989. Young et al.," 1989, Young et al."
+ 2001)., 2001).
+" Near-infrared observations indicate that CO and CH, are present on Triton’s surface with mixing ratios of 0.05 and 0.1 relative to N». and at least for CH. mostly in solid solution in N? (Cruikshank et al."," Near-infrared observations indicate that CO and $_4$ are present on Triton's surface with mixing ratios of 0.05 and 0.1 relative to $_2$, and at least for $_4$, mostly in solid solution in $_2$ (Cruikshank et al."
+ 1993. Quirico et al.," 1993, Quirico et al."
+" 1999, Grundy et al."," 1999, Grundy et al."
+ 2010)., 2010).
+ In this situation. the expected partial pressure of each species is the product of its solid mole fraction and its pure vapor pressure (Raoult's law for an ideal mixture).," In this situation, the expected partial pressure of each species is the product of its solid mole fraction and its pure vapor pressure (Raoult's law for an ideal mixture)."
+" This scenario leads to atmospheric abundances of CO and CH, that are about | and 3 orders of magnitude lower than observed. respectively (Fig."," This scenario leads to atmospheric abundances of CO and $_4$ that are about 1 and 3 orders of magnitude lower than observed, respectively (Fig."
+ 5)., 5).
+" Although Henry's law may be more applicable than Raoult's in the case of the N:-CH, system."," Although Henry's law may be more applicable than Raoult's in the case of the $_2$ $_4$ system,"
+noise and resolution properties and the presence of damped systems for the flux power.,noise and resolution properties and the presence of damped systems for the flux power.
+ We summarize here the constraints found., We summarize here the constraints found.
+" For the pdf: B=0.08£0.05 and 6<0.19 (2e C.L.) by using the flux pdf alone in the range F=[0.1— 0.8], with a reduced x/v=1.09 (35 d.o.£);"," For the pdf: $\beta=0.08\pm
+0.05$ and $\beta<0.19$ $\sigma$ C.L.) by using the flux pdf alone in the range $F=[0.1-0.8]$ , with a reduced $\chi^2/\nu=1.09$ (35 d.o.f.);"
+" 8=0.04+ and 8«0.1 (2c C.L.) by using the flux pdf alone in the whole range F=[0—1], with a reduced y?/v=1.21 (53 d.o.f.)."," $\beta=0.04\pm 0.04$ and $\beta<0.1$ $\sigma$ C.L.) by using the flux pdf alone in the whole range $F=[0-1]$, with a reduced $\chi^2/\nu=1.21$ (53 d.o.f.)."
+" The ranges at low and high transmissivity are those that are most difficult to model due to the presence of strong systems and continuum fitting errors, respectively."," The ranges at low and high transmissivity are those that are most difficult to model due to the presence of strong systems and continuum fitting errors, respectively."
+" Thus, we regard the first result presented as more conservative even though we do model continuum fitting errors and correct for numerical resolution (?).."," Thus, we regard the first result presented as more conservative even though we do model continuum fitting errors and correct for numerical resolution \citep{vbh09}."
+" For the flux power only we obtain: 6=0.07+0.04 and 6<0.14 (2e C.L.) using all the 132 data points (x?/v=1.16, for 120 d.o.f.)."," For the flux power only we obtain: $\beta=0.07\pm 0.04$ and $\beta<0.14$ $\sigma$ C.L.) using all the 132 data points $\chi^2/\nu=1.16$, for 120 d.o.f.)."
+ All these numbers are reasonable and demonstrate that the regions of high transmissivity have a constraining power which is stronger than the power spectrum alone., All these numbers are reasonable and demonstrate that the regions of high transmissivity have a constraining power which is stronger than the power spectrum alone.
+" If we combine the two measurements we find the same trends as in ?:: there is not a very good fit to the data (x?=200 for 164 d.o.f.),"," If we combine the two measurements we find the same trends as in \cite{vbh09}: there is not a very good fit to the data $\chi^2=200$ for 164 d.o.f.),"
+ and a reasonable x? is obtained only when neglecting the three highest redshift bins of the SDSS flux power., and a reasonable $\chi^2$ is obtained only when neglecting the three highest redshift bins of the SDSS flux power.
+" In this case, we obtain 8=0.05-£0.03 and 8«0.1 (2e C.L.) with a reduced x?/v=1.09 (146 d.o.£.)."," In this case, we obtain $\beta=0.05\pm 0.03$ and $\beta<0.1$ $\sigma$ C.L.) with a reduced $\chi^2/\nu=1.09$ (146 d.o.f.)."
+ All the other parameters are not affected significantly by the new parameter introduced and there are not strong degeneracies for B., All the other parameters are not affected significantly by the new parameter introduced and there are not strong degeneracies for $\beta$.
+ From the analysis performed we can conclude that robust 2c upper limits on the coupling constant are in the range 6«0.1—0.2 (depending on the subset of data chosen)., From the analysis performed we can conclude that robust $2\sigma$ upper limits on the coupling constant are in the range $\beta<0.1-0.2$ (depending on the subset of data chosen).
+ These bounds are exclusively derived by the analysis of the observed properties of the IGM and represent a completely new and independent test of cDE cosmologies w.r.t., These bounds are exclusively derived by the analysis of the observed properties of the IGM and represent a completely new and independent test of cDE cosmologies w.r.t.
+ previous constraints (ase.g.???)..," previous constraints \citep[as \eg
+][]{Bean_etal_2008,LaVacca_etal_2009,xia09}."
+ We regard a 2c limit of 8<0.15 as a conservative overall bound once the statistical limitations of the different samples are taken into account., We regard a $2\sigma $ limit of $\beta \lesssim 0.15$ as a conservative overall bound once the statistical limitations of the different samples are taken into account.
+ In this work we have explored the possibility of constraining the coupling 9 between CDM and DE through the statistical properties of the transmitted flux in forest QSO spectra at z—24.2., In this work we have explored the possibility of constraining the coupling $\beta $ between CDM and DE through the statistical properties of the transmitted flux in forest QSO spectra at $z=2-4.2$.
+" For this purpose, we have performed the first high-resolution hydrodynamical simulations with gas cooling and star formation in the context of cDE models and quantitatively exploited the capabilities of flux 1-pt and 2-pt functions to constrain the strength of the coupling 6 between DE and CDM."," For this purpose, we have performed the first high-resolution hydrodynamical simulations with gas cooling and star formation in the context of cDE models and quantitatively exploited the capabilities of flux 1-pt and 2-pt functions to constrain the strength of the coupling $\beta$ between DE and CDM."
+" 'The main results can be summarized as 'This work quantitatively shows that the range of scales and redshifts, where the growth of structures can be radically different from that measured from a naive extrapolation of either local or very high redshift probes, is promising for constraining coupled dark energy cosmologies."," The main results can be summarized as This work quantitatively shows that the range of scales and redshifts, where the growth of structures can be radically different from that measured from a naive extrapolation of either local or very high redshift probes, is promising for constraining coupled dark energy cosmologies."
+" The increasing number of QSO spectra that are being collected (e.g.BOSS?]. offers the exciting prospect of further improvingX-Shooteif]) the numbers and of understanding in a more refined way, by performing simulations and by addressing systematic errors, the impact that coupled dark energy cosmologies can have on the diffuse gas at high redshift."," The increasing number of QSO spectra that are being collected (e.g., ) offers the exciting prospect of further improving the numbers and of understanding in a more refined way, by performing simulations and by addressing systematic errors, the impact that coupled dark energy cosmologies can have on the diffuse gas at high redshift."
+" MB is supported by the DFG Cluster of Excellence ""Origin and Structure of the Universe"" and partly supported by the TRR Transregio Collaborative Research Network on the “Dark Universe”."," MB is supported by the DFG Cluster of Excellence “Origin and Structure of the Universe"" and partly supported by the TRR Transregio Collaborative Research Network on the “Dark Universe""."
+" MV is partly supported by ASI/AAE, INFN-PD51 and PRIN/MIUR."," MV is partly supported by ASI/AAE, INFN-PD51 and PRIN/MIUR."
+ Numerical simulations have been performed atRZG Computing Centre in Garching., Numerical simulations have been performed atRZG Computing Centre in Garching.
+" Post processing and data analysis have been carried out at COSMOS and HPCS (Cambridge), and CINECA thanks to a CINECA/INAF grant."," Post processing and data analysis have been carried out at COSMOS and HPCS (Cambridge), and CINECA thanks to a CINECA/INAF grant."
+Fieure G compares the joint distribution of colors and absolute magnitudes for ALLOL. AJSI. and the MIOL clusters found by Chandaretal.(2004).,"Figure \ref{cmd} compares the joint distribution of colors and absolute magnitudes for M101, M81, and the M101 clusters found by \citet{cwl04}."
+.. Chandaretal.(2004) were specifically attempting to select old. elobular clusters in M1OL. while were searching onlv [or ‘compact’ star clusters in M81. with no selection on age.," \citet{cwl04} were specifically attempting to select old, globular clusters in M101, while \citet{cft1} were searching only for `compact' star clusters in M81 with no selection on age."
+ The ellect of our {ραπ magnitude cut is particularly clear in the right-hand panel of this figure: our cluster sample is missing faint. blue clusters.," The effect of our $I$ -band magnitude cut is particularly clear in the right-hand panel of this figure: our cluster sample is missing faint, blue clusters."
+ It is clear that the excess of red clusters in M81 compared to ATLOL is mostly at brighter cluster Iuminositües: brighter than Af;—3.4 (the expected peak of the globular cluster luminosity function). the M81 clusters are mostly red. while the M1OI clusters are mostly blue.," It is clear that the excess of red clusters in M81 compared to M101 is mostly at brighter cluster luminosities: brighter than $M_V=-7.4$ (the expected peak of the globular cluster luminosity function), the M81 clusters are mostly red, while the M101 clusters are mostly blue."
+ We also confirm the detection by of a number of faint. red clusters in MIOI.," We also confirm the detection by \citet{cwl04} of a number of faint, red clusters in M101."
+ Some of these could be faint. backeround ealaxies. especially ellipticals or the bulges of spirals whose disks are (oo [aint to observe.," Some of these could be faint background galaxies, especially ellipticals or the bulges of spirals whose disks are too faint to observe."
+ Such contamination is unlikelv (o account for all of the faint red. clusters. whose nature is discussed further in Section 3.2..," Such contamination is unlikely to account for all of the faint red clusters, whose nature is discussed further in Section \ref{sec:spatdist}. ."
+ The cluster candidate luminosity distribution is shown in Figure 7.. for the full sample and (he ved and blue subsamples.," The cluster candidate luminosity distribution is shown in Figure \ref{lum-dist}, for the full sample and the red and blue subsamples."
+ The strong fall-olff at V.2»23 (Mj.=—6.1) is due to our imposed magnitude limit (see above)., The strong fall-off at $V>23$ $M_V>-6.1$ ) is due to our imposed magnitude limit (see above).
+ Brighter (han this limit. we find that the blue ]usters are about. 0.25 mae brighter in the median than the red clusters.," Brighter than this limit, we find that the blue clusters are about 0.25 mag brighter in the median than the red clusters."
+ From population svnthesis models. such a difference is consistent with the blue clusters being vounger. by a few Gyr if both populations are Z8 Gyr old. or less if the clusters are younger.," From population synthesis models, such a difference is consistent with the blue clusters being younger, by a few Gyr if both populations are $\gtrsim 8$ Gyr old, or less if the clusters are younger."
+ Inferred ages are of course stronglv dependent on additional factors such as metallicity. recldenine. and initial mass function.," Inferred ages are of course strongly dependent on additional factors such as metallicity, reddening, and initial mass function."
+" As ciseussed below. there are many more τος clusters than the expected number of globular clusters for a galaxy of MIOIs huninositw: (he smooth curve in the left. panel of the figure shows a ‘standard’ elobular cluster luminosity function (a Gaussian with mean A4,=—T.4 and standard deviation ¢= 1.3) scaled to the number of clusters with Ay«—1.4 (also see Section 3.2))."," As discussed below, there are many more red clusters than the expected number of globular clusters for a galaxy of M101's luminosity: the smooth curve in the left panel of the figure shows a `standard' globular cluster luminosity function (a Gaussian with mean $M_V=-7.4$ and standard deviation $\sigma=1.3$ ) scaled to the number of clusters with $M_V<-7.4$ (also see Section \ref{sec:spatdist}) )."
+" The bright end of the LF for the full sample is consistent with the power-law distribution of luminosity d:N(L)/dLxL? CN(L)xL 4), similar to the distributions seen for voung clusters in mergers (Whitmoreetal.1999:οἱal.1997). and Z/57-lentiliel (not necessarily voung) clusters in other spirals 2002).. as well as the bright end of the GCLF for the Milkv Way ancl M31. 1994)."," The bright end of the LF for the full sample is consistent with the power-law distribution of luminosity $dN(L)/dL \propto L^{-2}$ $N(L) \propto L^{-1}$ ), similar to the distributions seen for young clusters in mergers \citep{whi99,mwsf97}
+ and -identified (not necessarily young) clusters in other spirals \citep{lar02b}, as well as the bright end of the GCLF for the Milky Way and M31 \citep{hp94,mcl94}."
+. Do the blue and red subsamples of M1OI correspond to voung! aud ‘lobular cluster eroups?, Do the blue and red subsamples of M101 correspond to `young' and `globular' cluster groups?
+ One way to find out is to compare (he number of clusters per unit galaxy. Iuminosity in MIOLto values for other galaxies., One way to find out is to compare the number of clusters per unit galaxy luminosity in M101to values for other galaxies.
+ The total magnitudes of MIOLI as given by, The total magnitudes of M101 as given by
+ (Skrutskieetal., \citep{Skrutskie90}.
+1990).. ~30 µια. (Brownal., $\sim$ $\mu$ \citep{Brown07}.
+2007).. ~10? (Alexanderetal.2006).. (Liu&Papaloizou1979).. (e.8Jensen&Mathieu1997:Dutrev2005:White&Wil," $\sim10^5$ \citep{Alexander06}, \citep{Lin79}. \citep[e.g][]{Jensen97,Beust05,White05}."
+lebrand2005).. (Espaillatetal.2007) (Careutheretal., \citep{Espaillat07} \citep{Guenther07}.
+2007).. (e.g.Furlanetal.2007) in the Taurus star-forming region (l~2 MM. —115 ppc).," \cite[e.g.][]{Furlan07} in the Taurus star-forming region $\sim$ Myr, $\sim$ pc)."
+ οσα απο was discovered to have a large ~20- JS0;0u excess and no excess at wavelengths <& san throughTelescope Tutrared Spectrograph observatious (Forrestetal.2001)., CoKu Tau/4 was discovered to have a large $\sim$ $\mu$ m excess and no excess at wavelengths $<8$ $\mu$ m through Infrared Spectrograph observations \citep{Forrest04}.
+. The disk has been modeled as having an inner hole of radius —10 AAU (D'Alessioetal.2005).. with suggestions that this hole is due to a giant planet (Quillenctal.2001).," The disk has been modeled as having an inner hole of radius $\sim$ AU \citep{DAlessio05}, with suggestions that this hole is due to a giant planet \citep{Quillen04}."
+. Tn this paper. we describe nem-iufrared aperture-nasking interferometry and niaegnmeg observations that demonstrate that Colu Tau/l1 is a nuear-equal mass ünuarv star system.," In this paper, we describe near-infrared aperture-masking interferometry and imaging observations that demonstrate that CoKu Tau/4 is a near-equal mass binary star system."
+ We show that the predicted mner role size from dynamical models is comparable to. mt larecr than. that predicted from radiative transfer nodels.," We show that the predicted inner hole size from dynamical models is comparable to, but larger than, that predicted from radiative transfer models."
+ Finally. we discuss whether other so-called “transitional” disks could be circiuunubiuarv disks.," Finally, we discuss whether other so-called “transitional” disks could be circumbinary disks."
+ We conclude that for candidates in Taurus. much of the lass ratio-separation space for stellar companions cau by ruled out bv existiug observations. but that defuitively ruling out the possibility of binaritv is generally difficult or iudividual so-called “transitional” disks.," We conclude that for candidates in Taurus, much of the mass ratio-separation space for stellar companions can by ruled out by existing observations, but that definitively ruling out the possibility of binarity is generally difficult for individual so-called “transitional” disks."
+ Colku αι was observed with the NIRC2 camera behind Adaptive Optics (AO) at the heck IE telescope ou 2007 Nov 23 as part of an ongoing aperture-masking survey of nearby vouug star-forming associations., CoKu Tau/4 was observed with the NIRC2 camera behind Adaptive Optics (AO) at the Keck II telescope on 2007 Nov 23 as part of an ongoing aperture-masking survey of nearby young star-forming associations.
+ Apertureauaskiug interferometry (e.e.Tuthilletal.2000)x is a well established technique for achieving the full diffraction limit of a single telescope. receutly applied to observations belind adaptive optics svstenis (6.8.Lloydal.2006:Krauset 2008)..," Aperture-masking interferometry \citep[e.g.][]{Tuthill00} is a well established technique for achieving the full diffraction limit of a single telescope, recently applied to observations behind adaptive optics systems \citep[e.g.][]{Lloyd06,Kraus08a}."
+. A 9-hole mask was placed in a filter wheel near a pupil plane in the NIRC2 camera. euablius interfercuce fringes ou 236 baseline to be simultaneously recorded ou the cameras Waging array.," A 9-hole mask was placed in a filter wheel near a pupil plane in the NIRC2 camera, enabling interference fringes on 36 baselines to be simultaneously recorded on the camera's imaging array."
+ The observations of Colxu Tau/1 consisted of two iniage setstakenthroughaIC filter. each with eight 20 secondexposures.calibratedbytwointerleaved image sets of CN Tan.," The observations of CoKu Tau/4 consisted of two image setstakenthroughaK' filter, each with eight 20 secondexposures,calibratedbytwointerleaved image sets of CX Tau."
+ The airmass of Col&u Tau/1 observations varied, The airmass of CoKu Tau/4 observations varied
+demonstrates that the periodicity is appareut. but is nof detected with sufficicut significance to allow us to obtain a useful dprovement in the precision of the orbital period.,"demonstrates that the periodicity is apparent, but is not detected with sufficient significance to allow us to obtain a useful improvement in the precision of the orbital period."
+ Seven separate peaks in the inteusitv of this source are apparent in the latter part of the ASM light curve., Seven separate peaks in the intensity of this source are apparent in the latter part of the ASM light curve.
+ Their average separation is close to 111 davs and is probably not consistent with the period determined hrough pulse timing., Their average separation is close to 114 days and is probably not consistent with the period determined through pulse timing.
+ The outbursts of some other Be starfneutron star binaries do not occur at precisely the orbital periods determined by pulse timing. so there is 10 reason in this case to doubt the accuracy of the pulse uius value.," The outbursts of some other Be star/neutron star binaries do not occur at precisely the orbital periods determined by pulse timing, so there is no reason in this case to doubt the accuracy of the pulse timing value."
+ is à huumnous IIMXND in the Large Magellanic Cloud which highly likely comprises astcllar-nass black hole and its normal OB-type companion., is a luminous HMXB in the Large Magellanic Cloud which highly likely comprises astellar-mass black hole and its normal OB-type companion.
+ The properties of the svstem are described in detail by Oroszetal.(2009)., The properties of the system are described in detail by \citet{orosz09}.
+. Oroszetal.(2009). also present the ASM results and how they relate to other observations iucludiug other period deterünations., \citet{orosz09} also present the ASM results and how they relate to other observations including other period determinations.
+ Oroszotal.(2009) adopt the value of 3.909175E0.00005 days as their best estimate of the period based on optical photoimoetry and spectroscopy.," \citet{orosz09} adopt the value of $3.90917
+\pm 0.00005$ days as their best estimate of the period based on optical photometry and spectroscopy."
+ An updated ASAD power spectrma is shown in Fieure 13) aud revised estimates. based solely ou the ASM power spectrum. of the orbital frequency and period are eiven in Table 3..," An updated ASM power spectrum is shown in Figure \ref{fig:pdslmcx1} and revised estimates, based solely on the ASM power spectrum, of the orbital frequency and period are given in Table \ref{tbl:detect}."
+ We note that the ASAT detection provides the onlv reported evidence to date of the orbital period i N-aravs aud that our N-arav-based period estimate of P=3.90898+0.00021[0.00153] days is fully consistent with the opticallv-based value of Oroszetal. (2009).," We note that the ASM detection provides the only reported evidence to date of the orbital period in X-rays and that our X-ray-based period estimate of $P = 3.90898 \pm
+0.00021 [\pm 0.00153]$ days is fully consistent with the optically-based value of \citet{orosz09}."
+ is a trausicut X-ray pulsar in a Be/X-vay binary discovered with the BATSE experimceut ontheO, is a transient X-ray pulsar in a Be/X-ray binary discovered with the BATSE experiment on the.
+bservatory2007).. Shraderct used ~2 vears of ASM data to obtain au carly estimate of the outburst period of P~135 d. Levine&Corbet(2006) reported a relatively precise determination of the outburst period. {κε=218.940.5 d. that was based on the time intervals between widely spaced outbursts.," \citet{shrad99} used $\sim2$ years of ASM data to obtain an early estimate of the outburst period of $P \sim 135$ d. \citet{lcatel06}
+ reported a relatively precise determination of the outburst period, $P_{outburst} = 248.9 \pm 0.5$ d, that was based on the time intervals between widely spaced outbursts."
+ We believe this is still the best current estimate of the outburst cycle time., We believe this is still the best current estimate of the outburst cycle time.
+ The ASM power spectrum is shown in the top panel of Figure 1L., The ASM power spectrum is shown in the top panel of Figure \ref{fig:pdsj1008}.
+ Though the peaks due to the outburst periodicity are clear. they are not sufficiently well-defined to vield a superior period estimate.," Though the peaks due to the outburst periodicity are clear, they are not sufficiently well-defined to yield a superior period estimate."
+ Coeetal.(2007) have estimated the orbital period through timing of the— 93-s pulsations secu in the BATSE data., \citet{coe07} have estimated the orbital period through timing of the 93-s pulsations seen in the BATSE data.
+ They obtained the result P—217.5XU. d and noted that is it is iu good agreement with the result of Levine&Corbet(2006).," They obtained the result $P =
+247.8 \pm 0.4$ d and noted that is it is in good agreement with the result of \citet{lcatel06}."
+. is also a transicut pulsar in a BesX-ray binary system., is also a transient pulsar in a Be/X-ray binary system.
+ A brief description of the early history and of a pulse-timine analysis usine BATSE data may be found in Fingeretal.(1996)., A brief description of the early history and of a pulse-timing analysis using BATSE data may be found in \citet{fwc96}.
+. The BATSE timine analysis vielded an estimate of the orbital period of P=12.12£0.03 d as well as the projected seminajor axis. eccentricity. epoch of periastron passage. aud other orbital elements.," The BATSE timing analysis yielded an estimate of the orbital period of $P =
+42.12 \pm 0.03$ d as well as the projected semimajor axis, eccentricity, epoch of periastron passage, and other orbital elements."
+ A slight revision to the orbital period and time of periastrou passage have been eiven by Iunii L).., A slight revision to the orbital period and time of periastron passage have been given by \citet{inam04}.
+ Evidence of outbursts recurring at intervals that are close iu duration to the orbital period (or nmltiples thereof) is secu in the ASAD power spectrum (see Fie. 15))., Evidence of outbursts recurring at intervals that are close in duration to the orbital period (or multiples thereof) is seen in the ASM power spectrum (see Fig. \ref{fig:pds1417}) ).
+ Frequency aud period estimates obtaimed from the spectrum shown in Fie., Frequency and period estimates obtained from the spectrum shown in Fig.
+ 15 are listed iu Table but are not as precise as the pulse-timing values., \ref{fig:pds1417} are listed in Table \ref{tbl:detect} but are not as precise as the pulse-timing values.
+ is au X-ray pulsar with the rather long pulse period of ~1300 s (Lutovinovetal. 2005).., is an X-ray pulsar with the rather long pulse period of $\sim 1300$ s \citep{lutov05}. .
+ (e.gDuquennoy&Aavor&Marcy1992) (c.getal.2002).. (Goodwin," \citep[e.g][]{Duquennoy1991,Fischer1992} \citep[e.g][]{Patience2002}, \citep{Goodwin2009}."
+" AQ Ado. Af,>Af. AA, Mo Ξἀλλ. 4 ou the primary mass", $M_{1}$ $M_{2}$ $M_{1}>M_{2}$ $M_{1}$ $M_{2}$ $q=M_{2}/M_{1}$ $q$ on the primary mass.
+ —CTraditionallv these classes of nodels have been divided iuto capture aud fragmnenutatiou scenarios., Traditionally these classes of models have been divided into capture and fragmentation scenarios.
+ Capture refers to the tidal capture of two nabound objects on a timescale that is long compares o the collapse time of cach component (c.g.MeDon-ald&Clarke 1993)., Capture refers to the tidal capture of two unbound objects on a timescale that is long compared to the collapse time of each component \citep[e.g.][]{McDonald1993}.
+. For cach primary star the mass of he secondary is chosen randomly from the single star nass function aud the secondary-lmass distribution woul reflect the IMIF., For each primary star the mass of the secondary is chosen randomly from the single star mass function and the secondary-mass distribution would reflect the IMF.
+ While tidal capture appears to be too ineticicut in reproducing high binary fractious. it has con noticed that. particularly in small eroups of stars. stir-disk encounters may form binaries (MeDonald&Clarke 1995).," While tidal capture appears to be too inefficient in reproducing high binary fractions, it has been noticed that, particularly in small groups of stars, star-disk encounters may form binaries \citep{McDonald1995}."
+. Tn iu case even this disk assisted capture. whereby a star passing through the disk of another which dissipates enough kinetic cucrey to form a bound svsten. is unlikely to be the most relevant binary formation nechamsin (Doffinctal.1998).," In any case even this disk assisted capture, whereby a star passing through the disk of another which dissipates enough kinetic energy to form a bound system, is unlikely to be the most relevant binary formation mechanism \citep{Boffin1998}."
+. Fragiucntation scenarios are the preferred mechanism or the formation of multiple systems., Fragmentation scenarios are the preferred mechanism for the formation of multiple systems.
+ The so-called rasnientation models are usually classified as pronuipt ragiueutation (60.8.Boss1986:Bounell&Basticu1992) and disk fraeioeutation (e.g.Bonnell1991:Stamatellos&Wlitworth 2009).," The so-called fragmentation models are usually classified as prompt fragmentation \citep[e.g.][]{Boss1986,Bonnell1992} and disk fragmentation \citep[e.g.][]{Bonnell1994,Stamatellos2009}."
+. In the prowpt fragmentation scenario voth primary aud secondarystarsform byfragmentationofthesame collapsing molecular cloud core., In the prompt fragmentation scenario both primary and secondarystarsform byfragmentationofthesame collapsing molecular cloud core.
+ Disk ragiueutation takes place in a uewly formed star-disk, Disk fragmentation takes place in a newly formed star-disk
+times. the high-mass mass spectrum converges to a slope between = Bands=2.5 consistent with the accretion during the stellar dominated phase. and to the Salpeter slope of=235.,"times, the high-mass mass spectrum converges to a slope between $\gamma=-2$ and $\gamma=-2.5$ consistent with the accretion during the stellar dominated phase, and to the Salpeter slope of $\gamma=-2.35$."
+ 1n order ascertain whether we are correct in our interpretation of a two-power law LAL resulting from accretion in. gas-clominatecl ancl stellardominated: regimes. we evaluated. how much. of the eventual stellar. mass. was added: in each regime.," In order ascertain whether we are correct in our interpretation of a two-power law IMF resulting from accretion in gas-dominated and stellar-dominated regimes, we evaluated how much of the eventual stellar mass was added in each regime."
+ Figure 4 plots the amount. of mass accumulated by cach star during the stellar-dominate yhase against the final stellar mass.," Figure \ref{accdyn}
+ plots the amount of mass accumulated by each star during the stellar-dominated phase against the final stellar mass."
+ Llieh-mass stars accumulate the majority of their mass during this phase whereas most of the low-mass stars only acerete a smal raction of their mass during this phase., High-mass stars accumulate the majority of their mass during this phase whereas most of the low-mass stars only accrete a small fraction of their mass during this phase.
+ The break between he two regimes occurs at approximately the same fina mass (2 LM.) where the mass function displavs a break tween the two slopes., The break between the two regimes occurs at approximately the same final mass $\approx 1 \solm$ ) where the mass function displays a break between the two slopes.
+ This supports the assertion that the wo different power-laws in the LAL derive. from. dilferen ohvsical regimes which allect how the stars accrete., This supports the assertion that the two different power-laws in the IMF derive from different physical regimes which affect how the stars accrete.
+ The ow-niass stars derive thei mass from tical-lobe accretion during the gas-domünated regime whereas the high-mass stars derive the majority of their mass from a Boncdi-LHovle ype accretion that occurs in the inner parts of the cluster where the potential is dominated by the stars themselves., The low-mass stars derive their mass from tidal-lobe accretion during the gas-dominated regime whereas the high-mass stars derive the majority of their mass from a Bondi-Hoyle type accretion that occurs in the inner parts of the cluster where the potential is dominated by the stars themselves.
+ In addition to producing the two power-law mass spectrum. competitive accretion naturally results in a certain degree of mass segregation.," In addition to producing the two power-law mass spectrum, competitive accretion naturally results in a certain degree of mass segregation."
+ This arises due to the accretion in the gas-dominated phase where there is a strong correlation between aceretion rate. and thus the final mass. and position in the cluster. (sce equations(16)) ancl (20)).," This arises due to the accretion in the gas-dominated phase where there is a strong correlation between accretion rate, and thus the final mass, and position in the cluster (see \ref{tidmsvsrad}) ) and \ref{massvsrad}) )."
+ This. direct correlation between the final mass and position in the eluster neelects variations in the initial masses and the relative movements of the stars due to their interactions., This direct correlation between the final mass and position in the cluster neglects variations in the initial masses and the relative movements of the stars due to their interactions.
+" If the cluster is mass segregated entering the stellar. dominated ohase. then the Ag,xAZ? implies that the mass segregation =μαill persist."," If the cluster is mass segregated entering the stellar dominated phase, then the $\macc \propto \ms^2$ implies that the mass segregation will persist."
+ Simulations of accretion in clusters show that 10 Mass segregation does result. but that there is not a one to one correlation between mass and radius (Bonnell 2000)., Simulations of accretion in clusters show that the mass segregation does result but that there is not a one to one correlation between mass and radius (Bonnell 2000).
+ In fact. low-mass stars are located throughout 1e cluster. including in the core. but the high-mass stars are edominantly located in the central regions as is found in voung stellar clusters such as the ONC (Llillenbranc 1997).," In fact, low-mass stars are located throughout the cluster, including in the core, but the high-mass stars are predominantly located in the central regions as is found in young stellar clusters such as the ONC (Hillenbrand 1997)."
+ lt ds also worth noting that although the moclels oesented. here are meant to. consider accretion onto voung stars. they are equally appropriate for the erowth of clumps in a molecular cloud.," It is also worth noting that although the models presented here are meant to consider accretion onto young stars, they are equally appropriate for the growth of clumps in a molecular cloud."
+ As the clumps evolve owards eravitational instability. they will accrete from the surrounding σας and this accretion will be governed by the ohvsies described here.," As the clumps evolve towards gravitational instability, they will accrete from the surrounding gas and this accretion will be governed by the physics described here."
+ Thus. for example. the clump mass-tunction found by Motte. André Neri (1998) for the p Oph molecular cloud. could be due to the accretion by the pre-stellar clumps as the whole system collapses down to form a cluster.," Thus, for example, the clump mass-function found by Motte, André Neri (1998) for the $\rho$ Oph molecular cloud could be due to the accretion by the pre-stellar clumps as the whole system collapses down to form a cluster."
+ The observed y=3/2 slope would imply that the whole svstem is in a px27 density configuration and is subvirial (dominated by the diffuse gas not in the elumps)., The observed $\gamma=-3/2$ slope would imply that the whole system is in a $\rho \propto R^{-2}$ density configuration and is subvirial (dominated by the diffuse gas not in the clumps).
+ The steeper slope found by. Motte (1998) at the hieh-mass end of the mass spectrum can be interpreted as arising from a regionὃν of near-uniform oσας density., The steeper slope found by Motte (1998) at the high-mass end of the mass spectrum can be interpreted as arising from a region of near-uniform gas density.
+ A test of such a possibility is to estimate the degree of mass segregation of the clumps in this pre-stellar cluster system., A test of such a possibility is to estimate the degree of mass segregation of the clumps in this pre-stellar cluster system.
+" Finally. it is possible that the mass spectrum for massive stars. Al,<LOAL.. is significantly different than that for lower mass stars if they do not form in a similar fashion."," Finally, it is possible that the mass spectrum for massive stars, $\ms
+\simgreat 10 \solm$, is significantly different than that for lower mass stars if they do not form in a similar fashion."
+ Lf massive stars cannot acercte above 1041. due to the cllect of radiation pressure on the infalling dust (Yorke Ixrüsgecl 1977: Yorke 1993) but form through a merger process in a dense core (Bonnell 1998) then the expected mass spectrum could be significantly dilferent from that presented here., If massive stars cannot accrete above $10 \solm$ due to the effect of radiation pressure on the infalling dust (Yorke Krüggel 1977; Yorke 1993) but form through a merger process in a dense core (Bonnell 1998) then the expected mass spectrum could be significantly different from that presented here.
+ Competitive accretion in voung stellar clusters results in a two power-law mass spectrum with a slope of 5=3Mle for low-mass stars and a steeper slope of 252.5 for high-mass stars., Competitive accretion in young stellar clusters results in a two power-law mass spectrum with a slope of $\gamma\approx - 3/2$ for low-mass stars and a steeper slope of $-2 \ge \gamma \simgreat -2.5$ for high-mass stars.
+ The dillerent slopes are due to whether the eas or the stars dominate the cluster. potential., The different slopes are due to whether the gas or the stars dominate the cluster potential.
+ When eas dominates the cluster potential. the accretion rates are given by a tidal-DIobe accretion radius and results in a," When gas dominates the cluster potential, the accretion rates are given by a tidal-lobe accretion radius and results in a"
+The detailed variability study on the AGB. stars in NGC 1846 undertaken in Lebzelter Wood (2007)) allows us to investigate the observed star-to-star differences in the C/O ratio in light of the pulsational properties of these stars.,The detailed variability study on the AGB stars in NGC 1846 undertaken in Lebzelter Wood \cite{LW07}) ) allows us to investigate the observed star-to-star differences in the C/O ratio in light of the pulsational properties of these stars.
+ In refpl we show the location of our sample stars in the P-L-diagram of the cluster., In \\ref{pl} we show the location of our sample stars in the P-L-diagram of the cluster.
+ Data are taken from Lebzelter Wood., Data are taken from Lebzelter Wood.
+ LE9 could not be plotted due to the lack of any detected periodicity in its small-amplitude light. variability., LE9 could not be plotted due to the lack of any detected periodicity in its small-amplitude light variability.
+ The other stars can be easily attributed to a pulsation mode: LEI3 and LE16 are first overtone pulsators. while LES and H39 pulsate in the second overtone.," The other stars can be easily attributed to a pulsation mode: LE13 and LE16 are first overtone pulsators, while LE8 and H39 pulsate in the second overtone."
+ Long secondary periods beyond the fundamental mode period have been found in the light curves of several of these stars. but no shorter secondary periods.," Long secondary periods beyond the fundamental mode period have been found in the light curves of several of these stars, but no shorter secondary periods."
+ As already pointed out in Lebzelter Wood (2007)). LE17 cannot be attributed unequivocally to any pulsation sequence.," As already pointed out in Lebzelter Wood \cite{LW07}) ), LE17 cannot be attributed unequivocally to any pulsation sequence."
+ Stars located on the first overtone sequence show a higher C/O ratio than objects on the second overtone sequence refpl)., Stars located on the first overtone sequence show a higher C/O ratio than objects on the second overtone sequence \\ref{pl}) ).
+ As can be seen this cannot be attributed to a simple difference in luminosity as the two groups clearly overlap., As can be seen this cannot be attributed to a simple difference in luminosity as the two groups clearly overlap.
+ This is further illustrated in refemd2 where we marked the sample stars in the color- diagram according to their pulsation mode., This is further illustrated in \\ref{cmd2} where we marked the sample stars in the color-magnitude diagram according to their pulsation mode.
+ For LE9 (C/O=0.2). no periodic change was found by Lebzelter Wood (2007)).," For LE9 (C/O=0.2), no periodic change was found by Lebzelter Wood \cite{LW07}) )."
+ Assuming that it would have been nore likely to find a first overtone pulsation than a shorter period higher overtone pulsation. we attributed this star to the group of higher overtone pulsators.," Assuming that it would have been more likely to find a first overtone pulsation than a shorter period higher overtone pulsation, we attributed this star to the group of higher overtone pulsators."
+ These findings suggest that in addition to a relation between the C/O ratio and the luminosity. seems to exist some dependency of this ratio on the pulsation mode.," These findings suggest that in addition to a relation between the C/O ratio and the luminosity, seems to exist some dependency of this ratio on the pulsation mode."
+ At this point it ts not clear how to explain the existence of such a relation., At this point it is not clear how to explain the existence of such a relation.
+ As mentioned above. all sample stars are on the AGB. we can exclude an RGB contamination.," As mentioned above, all sample stars are on the AGB, we can exclude an RGB contamination."
+ From the difference in the C/O ratio. we can assume that the sample stars are in at least three different thermal pulse cycles.," From the difference in the C/O ratio, we can assume that the sample stars are in at least three different thermal pulse cycles."
+ Both from our models and from the literature (e.g. Vassiliadis Wood 1993)). we would expect that the stars become cooler from cycle to cycle.," Both from our models and from the literature (e.g. Vassiliadis Wood \cite{VW93}) ), we would expect that the stars become cooler from cycle to cycle."
+ As the two C enhanced stars LEI6 and LEI3 would form a sequence at higher temperatures than the other stars. it can be excluded that a difference in the thermal pulse cyele is responsible for the observed differences in pulsation mode and temperature.," As the two C enhanced stars LE16 and LE13 would form a sequence at higher temperatures than the other stars, it can be excluded that a difference in the thermal pulse cycle is responsible for the observed differences in pulsation mode and temperature."
+ A difference in pulsation mode between two stars of the same luminosity requires a difference in some other global characteristic of the star., A difference in pulsation mode between two stars of the same luminosity requires a difference in some other global characteristic of the star.
+ This could be a difference in mass., This could be a difference in mass.
+ Mackey Broby Nielsen (2007)) report indications for two stellar populations in NGC 1846 separated in age by about 300 Myr., Mackey Broby Nielsen \cite{MB07}) ) report indications for two stellar populations in NGC 1846 separated in age by about 300 Myr.
+ At a given time we may then find AGB stars of different masses at the same time on the AGB., At a given time we may then find AGB stars of different masses at the same time on the AGB.
+ Indeed such an age or mass difference could easily account for the observed color difference refemd2))., Indeed such an age or mass difference could easily account for the observed color difference \\ref{cmd2}) ).
+ This would also be compatible with the observed C/O ratios as the more massive stars may show a more efficient dredge up., This would also be compatible with the observed C/O ratios as the more massive stars may show a more efficient dredge up.
+ Furthermore. the starting abundances of the two populations may be different. which could enhance the observed difference in C/O and  C/UC. It has to be stressed that. even in the case of a single stellar population. a scatter in the abundances of C and O on the main sequence and after the first dredge up may be present.," Furthermore, the starting abundances of the two populations may be different, which could enhance the observed difference in C/O and $^{12}$ $^{13}$ C. It has to be stressed that, even in the case of a single stellar population, a scatter in the abundances of C and O on the main sequence and after the first dredge up may be present."
+ Such an abundance scatter has been found in various clusters., Such an abundance scatter has been found in various clusters.
+ An interesting case is LEI3., An interesting case is LE13.
+ This star has been classified as spectral type 83/3 by Lloyd Evans (1983))., This star has been classified as spectral type S3/3 by Lloyd Evans \cite{LE83}) ).
+ Indeed the ZrO band head at is clearly visible in the spectrum presented by him (see his 1H)., Indeed the ZrO band head at is clearly visible in the spectrum presented by him (see his 1).
+ According to its spectral classification. TiO-bands and ZrO bands should be of similar strength (Keenan Boeshaar 1980)).," According to its spectral classification, TiO-bands and ZrO bands should be of similar strength (Keenan Boeshaar \cite{KB80}) )."
+ YO should occur. but is not covered by the spectra shown by Lloyd Evans.," YO should occur, but is not covered by the spectra shown by Lloyd Evans."
+ As expected, As expected
+"factor Wer is the fraction of the eveuts produced above a given threshold £i, which are able to reach a cistance equal to the distance from the source to us. d. with an cucrey still above that sae threshold.","factor $W_{GZK}$ is the fraction of the events produced above a given threshold $E_{th}$ which are able to reach a distance equal to the distance from the source to us, $d$, with an energy still above that same threshold."
+" Assuming a power law spectrum at the sources dN ανE* this factor is just where ζειο) is the iitial cucrey that a CR misthave in order to survive with £=Ey, after traveling a distance d."," Assuming a power law spectrum at the sources $N$ $E\propto
+E^{-s}$ this factor is just where $E_i(E_{th},d)$ is the initial energy that a CR musthave in order to survive with $E=E_{th}$ after traveling a distance $d$."
+" The resulting factors for Ly,=60 aud SO EeV. (keep in imiud that there are still significant systematic macertaintics in the energy recoustructed in CR experiments) are shown in fie. ὃν,"," The resulting factors for $E_{th}=60$ and 80 EeV (keep in mind that there are still significant systematic uncertainties in the energy reconstructed in CR experiments) are shown in fig. \ref{fracp.fig},"
+ computed following reference |0].., computed following reference \cite{horizon}.
+ We asstmed a source spectral index of 2.2. but the results are not much scnusitive to the particular value adopted. aud considered a proton composition.," We assumed a source spectral index of 2.2, but the results are not much sensitive to the particular value adopted, and considered a proton composition."
+ The suppressious are qualitatively simular for Fe uuclei but are stronger for intermediate mass nuclei:, The suppressions are qualitatively similar for Fe nuclei but are stronger for intermediate mass nuclei.
+ Iucludius these factors iu the previous comparisons changes the results with respect to those obtained using the simple cutoff at LOO Mpe. but not iu a drastic way. although with larger unuuber of eveuts the proper iuclusion of the GZIN factors could become more nuportaut.," Including these factors in the previous comparisons changes the results with respect to those obtained using the simple cutoff at 100 Mpc, but not in a drastic way, although with larger number of events the proper inclusion of the GZK factors could become more important."
+ Another aspect that iav affect the source models is the fact that the catalogs considered here ouly contain sources brighter than a certain lint. aud although they may be quite complete above those liuiting briehtuesses. the fraction of sources which are faint increases significantly with increasing distance. and those wuebservecl sources may actually contribute to the CR fluxes. »ossdblv eiviug a more diffuse background.," Another aspect that may affect the source models is the fact that the catalogs considered here only contain sources brighter than a certain limit, and although they may be quite complete above those limiting brightnesses, the fraction of sources which are faint increases significantly with increasing distance, and those unobserved sources may actually contribute to the CR fluxes, possibly giving a more diffuse background."
+ In models where the sources are weighted by their fluxes he expected contribution from the unobserved sources is however reduced., In models where the sources are weighted by their fluxes the expected contribution from the unobserved sources is however reduced.
+ In the past some works 11.0] using the IRAS galaxy catalog. for which the selection effects as a function of distance are shown. corrected for the incompleteness of the catalog by dividiug the observed deusity of ealaxics at a eiven distance by the corresponding selection. fiction in order to obtain a more complete represcutation of the ealaxy distribution.," In the past some works \cite{wax,cuoco}
+ using the IRAS galaxy catalog, for which the selection effects as a function of distance are known, corrected for the incompleteness of the catalog by dividing the observed density of galaxies at a given distance by the corresponding selection function in order to obtain a more complete representation of the galaxy distribution."
+ A possible drawback of this approach is that one assigus he unobserved galaxies to the same locations where bright ealaxies are observed. aud this way not ο Very precise when the galaxies are sparscly siuupled. as happens at large distances.," A possible drawback of this approach is that one assigns the unobserved galaxies to the same locations where bright galaxies are observed, and this may not be very precise when the galaxies are sparsely sampled, as happens at large distances."
+ For sinaller catalogs. such as the SWIFT one. even if a selection function were kuown it would not be very realistic to asstune that the unobserved faint AGN are in the same locations as the bright ones.," For smaller catalogs, such as the SWIFT one, even if a selection function were known it would not be very realistic to assume that the unobserved faint AGN are in the same locations as the bright ones."
+ Auvway. if one restricts the sources to the ucarby ones (ising a distance cutoff or weighting them by the GZIs attenuation factors). iu the models where the sources are weighted bv their fluxes the contribution from the faint unobserved sources is not large. so that the model expectations obtained should still be reasonably accurate.," Anyway, if one restricts the sources to the nearby ones (using a distance cutoff or weighting them by the GZK attenuation factors), in the models where the sources are weighted by their fluxes the contribution from the faint unobserved sources is not large, so that the model expectations obtained should still be reasonably accurate."
+ For instance. if onemodels the Ipuuinositv distribution of the sources according to a Schechter fuuctiou.. «N/AxLSexptEf E). where a typical value," For instance, if onemodels the luminosity distribution of the sources according to a Schechter , $N/$ $L\propto L^a\exp(-L/L_*)$ , where a typical value"
+"where ff.) is the Dubble parameter at redshift τν e is fF(0.0) is the as a thefunction speedof ofpositionlight, ou the sky. aud O is samplingthe solid fractionangle.","where $H(z)$ is the Hubble parameter at redshift $z$, $c$ is the speed of light, $f(\theta,\phi)$ is the sampling fraction as a function of position on the sky, and $\Omega$ is the solid angle."
+ We take the function. f(0.0) to be a coustant over the eutire 8032 square deerce areal coverage of the SDSS DR? Leeacy survey.," We take the function $f(\theta, \phi)$ to be a constant over the entire 8032 square degree areal coverage of the SDSS DR7 Legacy survey."
+ The augular part of Equation 7 thus becomes (1/17)(8032/112253)(13) = 0.19£7., The angular part of Equation \ref{vmax-integral} thus becomes $(1/4\pi) (8032/41253) (4\pi)$ = 0.1947.
+" V, values are given in the data table presented in Section ??..", $V_{max}$ values are given in the data table presented in Section \ref{bigtables}.
+ We based the quality assessment of our measured structural paraicters on three science-notivated metres., We based the quality assessment of our measured structural parameters on three science-motivated metrics.
+ Figure 7? shows the size-Iununositv of disks iu the g-baud., Figure \ref{disk-lumsize} shows the size-luminosity of disks in the $g$ -band.
+ Disks are expected to follow a well-defined, Disks are expected to follow a well-defined
+dvnamically in the merger remnant.,dynamically in the merger remnant.
+ Dynamical friction might be the cause of the slower velocity. of rotation of the GC's of GI compared to the gaseous component., Dynamical friction might be the cause of the slower velocity of rotation of the GCs of G1 compared to the gaseous component.
+ The peak at |Fe/II] — - 1 in all three groups may be an indication of a significant contribution of 10? to LOMM. haloes during the hierarchical galaxy formation. assuming that these haloes produced GCs with roughly the same metallicity ancl (hat thev follow the mass-metallicitv relation al z  1 (Lamareille et al.," The peak at [Fe/H] = - 1 in all three groups may be an indication of a significant contribution of $10^{9}$ to $10^{10} M_{\odot}$ haloes during the hierarchical galaxy formation, assuming that these haloes produced GCs with roughly the same metallicity and that they follow the mass-metallicity relation at z $\simeq$ 1 (Lamareille et al."
+ 2009)., 2009).
+ For comparison. we mention (hat the metallicity distribution of M87. another giant. elliptical galaxy with an active nucleus. also peaks around - 1 (Cohen et al.1993). but. unlike NGC 5128. does not have a very. low metallicity tail.," For comparison, we mention that the metallicity distribution of M87, another giant elliptical galaxy with an active nucleus, also peaks around - 1 (Cohen et al.1998), but, unlike NGC 5128, does not have a very low metallicity tail."
+ There are in [act other analogies between the metallicity distributions of the three groups and those of GC's in nearby galaxies: for example between the GCs of AIL 31 (Darmby et al., There are in fact other analogies between the metallicity distributions of the three groups and those of GCs in nearby galaxies: for example between the GCs of M 31 (Barmby et al.
+ 2000) and GI. or the GCs of the LAICG (Beasley et al.," 2000) and G1, or the GCs of the LMC (Beasley et al."
+ 2002) and the metal-rich subgroup of G3., 2002) and the metal-rich subgroup of G3.
+ However. (hese comparisons can only give order of magnitude estimates. since the LMC may have been metal-enrichecl by interactions with our Galaxy. and M 31 might have a very different evolutionary. history [rom the progenitor galaxy of Gl. which merged at least 200 Myrs ago with NGC 5123.," However, these comparisons can only give order of magnitude estimates, since the LMC may have been metal-enriched by interactions with our Galaxy, and M 31 might have a very different evolutionary history from the progenitor galaxy of G1, which merged at least 200 Myrs ago with NGC 5128."
+ In summary. the above scenario rests mainlv on the assuniption of accretion events which shaped the evolution of NGC 5128 and its GCs.," In summary, the above scenario rests mainly on the assumption of accretion events which shaped the evolution of NGC 5128 and its GCs."
+ The latter have several possible origins: the GCs of G2 were produced in a major merger. while the GC's of the (wo other eroups were pre-existing in smaller galaxies Chat were subsequently accreted and disrupted.," The latter have several possible origins: the GCs of G2 were produced in a major merger, while the GCs of the two other groups were pre-existing in smaller galaxies that were subsequently accreted and disrupted."
+ This favors the categories ii) and v) listed in the introduction lor the lormation of the galaxy itself. namely a major merger and several accretions and in-situ merging.," This favors the categories ii) and v) listed in the introduction for the formation of the galaxy itself, namely a major merger and several accretions and in-situ merging."
+ The proposed scenario remains highlv speculative. in the absence of spectroscopically determined ages and metallicities for most GCs in the galaxy.," The proposed scenario remains highly speculative, in the absence of spectroscopically determined ages and metallicities for most GCs in the galaxy."
+"a modified Planck function, which has some given dust temperature and opacity.","a modified Planck function, which has some given dust temperature and opacity."
+" We apply an opacity law with the standard form &οςA~°, where a value of 8=0 would correspond to a perfect blackbody."," We apply an opacity law with the standard form $\kappa \propto \lambda^{-\beta}$, where a value of $\beta=0$ would correspond to a perfect blackbody."
+" In our model we set B—1.5 (Dunne Eales, 2001)."," In our model we set $\beta=1.5$ (Dunne Eales, 2001)."
+" Emission from polycyclic aromatic hydrocarbons (PAHs) is an important spectral feature between 3 and 13 microns (Leger Puget, 1984; Allamandola et al.,"," Emission from polycyclic aromatic hydrocarbons (PAHs) is an important spectral feature between 3 and 13 microns (Leger Puget, 1984; Allamandola et al.,"
+ 1985)., 1985).
+" In our spiral galaxy model we include a fairly schematic representation of the PAH features, designed to broadly resemble the spectra of galaxies observed in the Spitzer Infrared Nearby Galaxies Survey (SINGS; Kennicutt et al.,"," In our spiral galaxy model we include a fairly schematic representation of the PAH features, designed to broadly resemble the spectra of galaxies observed in the Spitzer Infrared Nearby Galaxies Survey (SINGS; Kennicutt et al.,"
+" 2003), presented by Draine et al. ("," 2003), presented by Draine et al. ("
+2007).,2007).
+ We show the IR spectra of our early- and late-types in reffig:modelseds.., We show the IR spectra of our early- and late-types in \\ref{fig:modelseds}.
+" While the modelling of the PAH features in the late-types is somewhat rudimentary, we find it is sufficient to reproduce the observed colours and sources counts of galaxies which sample this part of the spectrum, so we consider it sufficient for the purposes of this paper."," While the modelling of the PAH features in the late-types is somewhat rudimentary, we find it is sufficient to reproduce the observed colours and sources counts of galaxies which sample this part of the spectrum, so we consider it sufficient for the purposes of this paper."
+" reffig:modelseds also includes a third SED, which represents a starburst component modelled on Markarian 33 (Mrk33)."," \\ref{fig:modelseds} also includes a third SED, which represents a starburst component modelled on Markarian 33 (Mrk33)."
+" This is an additional spectral type which we introduce when dealing with observations at A>24um ((see refssec:24,, where the motivation for including this component is discussed)."," This is an additional spectral type which we introduce when dealing with observations at $\lambda \ge 24$ (see \\ref{ssec:24}, where the motivation for including this component is discussed)."
+" We assume a dust temperature of 30K, which is found to be a typical temperature of interstellar dust in star-forming galaxies (Farrah et al.,"," We assume a dust temperature of 30K, which is found to be a typical temperature of interstellar dust in star-forming galaxies (Farrah et al.,"
+" 2003; Pope et al.,"," 2003; Pope et al.,"
+" 2006; Coppin et al.,"," 2006; Coppin et al.,"
+" 2008; Elbaz et al.,"," 2008; Elbaz et al.,"
+ 2010)., 2010).
+" We take Ap=0.3 magnitudes for the normalisation of the 1/A dust absorption law; this is the value determined by Metcalfe et ((2001) when using this model at optical wavelengths, and is a fairly conservative amount of extinction."," We take $A_B=0.3$ magnitudes for the normalisation of the $1/\lambda$ dust absorption law; this is the value determined by Metcalfe et (2001) when using this model at optical wavelengths, and is a fairly conservative amount of extinction."
+" The total integrated flux of the dust emission, in the form of both the PAH features and the blackbody, is normalised to be equal to the total absorbed flux."," The total integrated flux of the dust emission, in the form of both the PAH features and the blackbody, is normalised to be equal to the total absorbed flux."
+" For each galaxy type, we use a luminosity function (LF) defined locally in the B band and correct it into the appropriate IR band using a rest-frame z=0 colour."," For each galaxy type, we use a luminosity function (LF) defined locally in the B band and correct it into the appropriate IR band using a rest-frame $z=0$ colour."
+ The resulting LF parameters for the K band are given in Table 1.., The resulting LF parameters for the $K$ band are given in Table \ref{t-lf}.
+" Then using k+e corrections determined by the BC03 evolution code, the LF is evolved back from the present day."," Then using $k+e$ corrections determined by the BC03 evolution code, the LF is evolved back from the present day."
+" The model is a simplification of the model of Metcalfe et ((2001, 2006) since rather than 5 independent colours, one for each type, we use only two, one for all of the early- (E/SO and Sab) and one for all of the late-types (Sbc, Scd and Sdm)."," The model is a simplification of the model of Metcalfe et (2001, 2006) since rather than 5 independent colours, one for each type, we use only two, one for all of the early-types (E/S0 and Sab) and one for all of the late-types (Sbc, Scd and Sdm)."
+ This is justified on the basis that the band-band colour differences in the NIR are smaller than in the optical., This is justified on the basis that the band-band colour differences in the NIR are smaller than in the optical.
+ The z=0 colours are taken from the BC03 model predictions., The $z=0$ colours are taken from the BC03 model predictions.
+ Table 2 shows the values of M* , Table \ref{t-cols} shows the values of $M^*$ 
+P. the epoch of transit centre Το. the RV. semi-amplituc[4 K. ecosw esinc (e being the eccentricity and w the angle of the periastron). and VsinZcosp Vsindsin. with VsinI being the projection of the stellar equatorial rotation. and 6 the projection of the angle between the stellar spin axis and the planetary orbit axis.,"$P$, the epoch of transit centre $T_0$, the RV semi-amplitude $K$, $e\,cos\,\omega$ $e\,sin\,\omega$ $e$ being the eccentricity and $\omega$ the angle of the periastron), and $V\,sin\,I\,cos \,\beta$ $V\,sin\,I\,sin\,\beta$, with $V\,sin\,I$ being the projection of the stellar equatorial rotation, and $\beta$ the projection of the angle between the stellar spin axis and the planetary orbit axis."
+" Di addition. we employed free normalization factors. for each lightcurve (WASP and Euler) and each set of radial velocity (y,, for and y, for ». which enablec variations to be made in instrumental zero points."," In addition, we employed free normalization factors for each lightcurve (WASP and Euler) and each set of radial velocity $\gamma_\textit{\tiny H}$ for and $\gamma_\textit{\tiny C}$ for ), which enabled variations to be made in instrumental zero points."
+ From these parameters. parameters were derived to characterise the planetary system.," From these parameters, parameters were derived to characterise the planetary system."
+ The best-fit set of parameters that minimize the y? (reduced y is 0.86) are listed in Table 2. as well as their related computed physical parameters., The best-fit set of parameters that minimize the $\chi^2_r$ (reduced $\chi^2$ is 0.86) are listed in Table \ref{tab:params} as well as their related computed physical parameters.
+ With this best-fit solution one computes for the data y=204 with 48 measurements. and for HARPS data y=188 with 82 measurements which implies that additional jittering is present that is not accounted for by the fitted model.," With this best-fit solution one computes for the data $\chi^2=204$ with 48 measurements, and for HARPS data $\chi^2=188$ with 82 measurements which implies that additional jittering is present that is not accounted for by the fitted model."
+ Since the main deviation is related to the data. the uncertainties in the orbital solutions are most likely underestimated.," Since the main deviation is related to the data, the uncertainties in the orbital solutions are most likely underestimated."
+ However. the error bars in the Rossiter parameters are driven mostly by the HARPS on-transit data. and one can assume that they are almost correct.," However, the error bars in the Rossiter parameters are driven mostly by the HARPS on-transit data, and one can assume that they are almost correct."
+ Our best-fit solution corresponds to a giant planet with an eccentric (e= 0.3)) 8.16-day orbit and an additional long-term radial-velocity drift of!.., Our best-fit solution corresponds to a giant planet with an eccentric $e=0.3$ ) 8.16-day orbit and an additional long-term radial-velocity drift of.
+" The planet is dense with 2.25M; and a radius of 1.04 in contrast to the substantial fraction of ""inflated"" hot Jupiters."," The planet is dense with $2.25\,M_j$ and a radius of $1.04\,R_j$, in contrast to the substantial fraction of `inflated"" hot Jupiters."
+"Rj. Surprisingly. the projected angle between the orbital and stellar spin axes is found to be 6=123.3"". indicative of à retrograde orbit."," Surprisingly, the projected angle between the orbital and stellar spin axes is found to be $\beta = 123.3^\circ$, indicative of a retrograde orbit."
+" We note that Vsind21.59kms""! iis in accordance with the line rotation broadening (in Table 1)) derived by the spectral analysts."," We note that $V\,sin\,I=1.59$ is in accordance with the line rotation broadening (in Table \ref{wasp8-params}) ) derived by the spectral analysis."
+ We checked whether the partial defocusing of HARPS during the transit spectroscopic sequence had any effect on our result., We checked whether the partial defocusing of HARPS during the transit spectroscopic sequence had any effect on our result.
+ We divided the series into two subsets and considered for each of them an independent offset (y)., We divided the series into two subsets and considered for each of them an independent offset $\gamma$ ).
+ We obtain a solution with a marginal improvement in the y., We obtain a solution with a marginal improvement in the $\chi^2$.
+ By comparing the solution obtained from these two sets with that for the complete set. the angle6 was changed by 1.5c.," By comparing the solution obtained from these two sets with that for the complete set, the angle $\beta$ was changed by $1.5\,\sigma$."
+ The defocusing problem does not affect the results of this paper., The defocusing problem does not affect the results of this paper.
+ The detection of a hot Jupiter on an eccentric orbit that is misaligned with the. stellar rotation axis and moving m a retrograde directior raises many questions about the origins of this system., The detection of a hot Jupiter on an eccentric orbit that is misaligned with the stellar rotation axis and moving in a retrograde direction raises many questions about the origins of this system.
+ Although the answer is beyond the scope of this paper. the visual faint companion and the drifting y velocity of the system are key components of the puzzle.," Although the answer is beyond the scope of this paper, the visual faint companion and the drifting $\gamma$ velocity of the system are key components of the puzzle."
+ From the observed separation between the A and B components. one can derive a most likely orbital semi-major axis (a=1.35x600 AAU) (Duquennoy&Mayor1991)..," From the observed separation between the A and B components, one can derive a most likely orbital semi-major axis $a=1.35\rho\approx600$ AU) \citep{1991A&A...248..485D}."
+ The observed radial-velocity drift is therefore unlikely to be related to the B component of the binary (Y«GMa*<Ems yr). suggesting that these is an additional closer companion of both unknown mass and period.," The observed radial-velocity drift is therefore unlikely to be related to the B component of the binary $\dot{\gamma}<GMa^{-2}<1$ ), suggesting that these is an additional closer companion of both unknown mass and period."
+ The lack of curvature indicates that the companion is more massive than the transiting planet., The lack of curvature indicates that the companion is more massive than the transiting planet.
+ This intermediate body is very likely to play a significant dynamical role in the system., This intermediate body is very likely to play a significant dynamical role in the system.
+"lightcone, although typically at different redshifts and so with different properties and at offset positions (due to large-scale motions).","lightcone, although typically at different redshifts and so with different properties and at offset positions (due to large-scale motions)."
+" ? suggested applying a series of transformations (rotations, translations and inversions) when tiling space with periodic replications."," \citet{Blaizot2005} suggested applying a series of transformations (rotations, translations and inversions) when tiling space with periodic replications."
+" This does not, of course, prevent multiple appearances of a given object within the lightcone, but these duplicates are then viewed from different directions and no longer fall on a (nearly) regular lattice."," This does not, of course, prevent multiple appearances of a given object within the lightcone, but these duplicates are then viewed from different directions and no longer fall on a (nearly) regular lattice."
+" Unfortunately, this technique also introduces discontinuities in large-scale structure at the boundaries between replications, affecting clustering statistics in a way which is at least as difficult to model as that of the original periodicity."," Unfortunately, this technique also introduces discontinuities in large-scale structure at the boundaries between replications, affecting clustering statistics in a way which is at least as difficult to model as that of the original periodicity."
+" ? showed that for lightcones of relatively small solid angle, the central line-of-sight can be chosen to pass through the lattice of periodic replications in such a direction that multiple images of the same object are minimised or eliminated altogether."," \citet{Kitzbichler2007} showed that for lightcones of relatively small solid angle, the central line-of-sight can be chosen to pass through the lattice of periodic replications in such a direction that multiple images of the same object are minimised or eliminated altogether."
+" The latter is not possible if the comoving volume of the lightcone exceeds that of the simulation, but this technique can still be used to ensure that multiple appearances occur as far apart as possible both on the sky and in redshift."," The latter is not possible if the comoving volume of the lightcone exceeds that of the simulation, but this technique can still be used to ensure that multiple appearances occur as far apart as possible both on the sky and in redshift."
+ We therefore use the method of ? in this paper., We therefore use the method of \citet{Kitzbichler2007} in this paper.
+" Space is filled with periodic replications of the simulation, a position is chosen for the observer, and the central line-of-sight of the survey field is given a previously chosen orientation."," Space is filled with periodic replications of the simulation, a position is chosen for the observer, and the central line-of-sight of the survey field is given a previously chosen orientation."
+ Galaxies whose positions intercept the lightcone are selected and their comoving distance is converted into a redshift., Galaxies whose positions intercept the lightcone are selected and their comoving distance is converted into a redshift.
+" As explained in ?,, the time between stored snapshots for the Millennium Simulation varies between 100 and 380 Myr."," As explained in \citet{Kitzbichler2007}, the time between stored snapshots for the Millennium Simulation varies between 100 and 380 Myr."
+ This means that the intrinsic properties of galaxies are not generally available at the time corresponding to their comoving distance., This means that the intrinsic properties of galaxies are not generally available at the time corresponding to their comoving distance.
+ Rather they must be taken from the stored snapshot which is closest to their light-cone position., Rather they must be taken from the stored snapshot which is closest to their light-cone position.
+" Hence, galaxies with redshift (z+zi-1)/2«z«(%+2i41)/2 are assigned the physical properties stored at z;."," Hence, galaxies with redshift $(z_i+z_{i-1})/2<z<(z_i+z_{i+1})/2$ are assigned the physical properties stored at $z_i$."
+" The resulting discontinuity in galaxy population properties, at the boundaries between snapshots, could be reduced by interpolating, but this works poorly for positions and velocities since the output separation is comparable to orbital times within groups and clusters."," The resulting discontinuity in galaxy population properties, at the boundaries between snapshots, could be reduced by interpolating, but this works poorly for positions and velocities since the output separation is comparable to orbital times within groups and clusters."
+" Moreover, it is not straightforward for other galaxy properties either since these change discontinuously on timescales shorter than the output spacing, for example through mergers and starbursts."," Moreover, it is not straightforward for other galaxy properties either since these change discontinuously on timescales shorter than the output spacing, for example through mergers and starbursts."
+ We thus follow ? and do not attempt any interpolation., We thus follow \citet{Kitzbichler2007} and do not attempt any interpolation.
+ The semi-analytic calculations are perfomed on these intermediate time-steps that vary between 5 and 15 Myr., The semi-analytic calculations are perfomed on these intermediate time-steps that vary between 5 and 15 Myr.
+" This means, for example, that a burst of star formation will have this duration and can happen anywhere between (or at) output snapshots, with the corresponding increase in flux being reflected in galaxy properties at the snapshot."," This means, for example, that a burst of star formation will have this duration and can happen anywhere between (or at) output snapshots, with the corresponding increase in flux being reflected in galaxy properties at the snapshot."
+ Theapparent luminosities and colours of galaxies depend strongly on their redshifts through the conversion between rest- and observed-frame photometric bands and through the inverse square dependence of apparent luminosity on distance., The luminosities and colours of galaxies depend strongly on their redshifts through the conversion between rest- and observed-frame photometric bands and through the inverse square dependence of apparent luminosity on distance.
+ The final redshift of the galaxy in the lightcone is not available at the time observed-frame luminosities are computed in the semi-analytic model., The final redshift of the galaxy in the lightcone is not available at the time observed-frame luminosities are computed in the semi-analytic model.
+" However, there will be two extreme redshifts that bracket it."," However, there will be two extreme redshifts that bracket it."
+" We compute apparent observed-frame luminosities (for fixed intrinsic properties) using these upper and lower limits, and once the galaxy is placed in the lightcone, we interpolate to obtain final observed-frame quantities."," We compute apparent observed-frame luminosities (for fixed intrinsic properties) using these upper and lower limits, and once the galaxy is placed in the lightcone, we interpolate to obtain final observed-frame quantities."
+ For this paper we construct lightcones for square areas of 1.4x1.4 deg? out to high redshift with no faint magnitude cut.," For this paper we construct lightcones for square areas of $1.4\times
+1.4$ $^2$ out to high redshift with no faint magnitude cut."
+ They are however limited by the mass resolution of the dark matter simulation (1.7x10'°h7'Me in halo mass) corresponding to stellar masses of ~109? at z=0., They are however limited by the mass resolution of the dark matter simulation $1.7\times10^{10}h^{-1}\rm{M}_{\sun}$ in halo mass) corresponding to stellar masses of $\sim10^{9.5}\rm{M}_{\sun}$ at z=0.
+" While this does not matter for the questions we studyM in this paper, it should be borne in mind if the lightcones are used for other Semi-analytic models predict intrinsic properties of galaxies, such as stellar mass, star formation history, gas and dust content and metallicity."," While this does not matter for the questions we study in this paper, it should be borne in mind if the lightcones are used for other Semi-analytic models predict intrinsic properties of galaxies, such as stellar mass, star formation history, gas and dust content and metallicity."
+" In order to convert these into observed spectral energy distributions (SED) or broad-band photometry, evolutionary population synthesis and dust models are required."," In order to convert these into observed spectral energy distributions (SED) or broad-band photometry, evolutionary population synthesis and dust models are required."
+" The former predict the evolution of the light associated with a single short burst of star formation of given metallicity and with an assumed Initial Mass Function (IMF), a so-called Simple Stellar Population (SSP)."," The former predict the evolution of the light associated with a single short burst of star formation of given metallicity and with an assumed Initial Mass Function (IMF), a so-called Simple Stellar Population (SSP)."
+ The intrinsic stellar emission from a model galaxy is then represented as a superposition of SSPs weighted, The intrinsic stellar emission from a model galaxy is then represented as a superposition of SSPs weighted
+with masses below 1.3 will on average be unevolved.,with masses below 1.3 will on average be unevolved.
+ The ratio of double degenerate svstems to those with MS companions has been addressed in detailed population synthesis studies (Iben et al., The ratio of double degenerate systems to those with MS companions has been addressed in detailed population synthesis studies (Iben et al.
+ 1997)., 1997).
+ Iben et al., Iben et al.
+ assumed (hat only WDs with companions less than 0.3 wwere detectable. and such low mass svstems would be of the detectable total.," assumed that only WDs with companions less than 0.3 were detectable, and such low mass systems would be of the detectable total."
+ For a flat relative mass function this implies that in the underlving population of all svstems would have MS companions with masses less (han 1.3., For a flat relative mass function this implies that in the underlying population of all systems would have MS companions with masses less than 1.3.
+Af... However. for hotter WDs one would detect à larger fraction of systems.," However, for hotter WDs one would detect a larger fraction of systems."
+ Out of all systems would be undetected (CM> 0.65. )). detected with an MS companion GV< )). and. would be double degenerates.," Out of all systems would be undetected $M_{\rm secondary} >$ 0.65 ), detected with an MS companion $M_{\rm secondary} <$ ), and would be double degenerates."
+ With this assumption an observed relative formation rate of [or binary WDs becomes an intrinsic population of5., With this assumption an observed relative formation rate of for binary WDs becomes an intrinsic population of.
+9%.. This is entirely consistent with the predicted binary fraction of from Duquennoy Mawor (1991) period distribution., This is entirely consistent with the predicted binary fraction of from Duquennoy Mayor (1991) period distribution.
+ We can use the 241ASS data to check on the fraction of systems with low mass companions., We can use the 2MASS data to check on the fraction of systems with low mass companions.
+ With a field binary fraction of we would expect half of the “single” WDs to actually have distant. and unresolved non-interacting binary companions: (hese companions will tvpically be low mass stars., With a field binary fraction of we would expect half of the “single” WDs to actually have distant and unresolved non-interacting binary companions; these companions will typically be low mass stars.
+ We expect of the low mass WDs produced from binary interactions to have close MIS companions and the other to be double degenerate svstems., We expect of the low mass WDs produced from binary interactions to have close MS companions and the other to be double degenerate systems.
+ In the Liebert οἱ al., In the Liebert et al.
+ sample there are 14 low mass WDs with τω>25.000 IX that were used to estimate the formation rate without applving an incompleteness correction.," sample there are 14 low mass WDs with $T_{\rm eff} > 25,000$ K that were used to estimate the formation rate without applying an incompleteness correction."
+ We would expect 2 double degenerate svstems. 5 close binary companions. ancl 3.5 distant binary. companions.," We would expect 2 double degenerate systems, 5 close binary companions, and 3.5 distant binary companions."
+ Three systems are known {ο have degenerate companions. aud we find significant J and Ix band excesses [or 4 of the svstems.," Three systems are known to have degenerate companions, and we find significant J and K band excesses for 4 of the systems."
+ Dased upon a comparison with isochrones. we infer masses ol 0.3-0.35 {for 3 companions and ~0.5 {for one.," Based upon a comparison with isochrones, we infer masses of 0.3-0.35 for 3 companions and $\sim$ 0.5 for one."
+ This indicates a deficit of companions relative to expectations. and in particular we note that there are no hieher mass companions seen even (though a significant number would be naturally produced by binary evolution.," This indicates a deficit of companions relative to expectations, and in particular we note that there are no higher mass companions seen even though a significant number would be naturally produced by binary evolution."
+ If all svstems were [rom interacting binaries. aud the sample was complete. we would expect LO MS companions.," If all systems were from interacting binaries, and the sample was complete, we would expect 10 MS companions."
+ We therefore conclude that even (he hot sample of He WDs is likely to be incomplete due to the presence of nearby companions., We therefore conclude that even the hot sample of He WDs is likely to be incomplete due to the presence of nearby companions.
+ Until a radial velocity survey. is completed. il is also not clear whether the MS companions detected are close binaries (attributed to the binary lormation channel) or distant ones (attributed) to the single Formation channel).," Until a radial velocity survey is completed, it is also not clear whether the MS companions detected are close binaries (attributed to the binary formation channel) or distant ones (attributed to the single formation channel)."
+ Follow-up work of this (wpe will be essential to empirically set the intrinsic space densities., Follow-up work of this type will be essential to empirically set the intrinsic space densities.
+considering that the photon index changes by 0.5 below and above the cooling break frequency. the maximum photon index allowed in the X-ray band is ay«a40.5=2. irrespective of by.,"considering that the photon index changes by 0.5 below and above the cooling break frequency, the maximum photon index allowed in the X-ray band is $\alpha_X < \alpha_X' + 0.5
+= 2$, irrespective of $b_X$."
+ Now we compare these constraints on ay with the range required to 10nize the jet., Now we compare these constraints on $\alpha_X$ with the range required to ionize the jet.
+ The observed X-ray flux at the day 5 should not be much different at the day ~10-13. and we extrapolate the X-ray luminosity down to the UV band and compare to Loree.," The observed X-ray flux at the day 5 should not be much different at the day $\sim$ 10–13, and we extrapolate the X-ray luminosity down to the UV band and compare to $L_{\rm ph, rec}$."
+ Note that only a fraction of the X-ray luminosity is directed to the jet material. and this fraction is given by ~b/by from a geometrical consideration.," Note that only a fraction of the X-ray luminosity is directed to the jet material, and this fraction is given by $\sim b/b_X$ from a geometrical consideration."
+ We found that the spectral index must be extremely soft as ayz 5 or 4 for by= 1 or by=b0.1. respectively. in order that the extrapolated flux down to vp is equal to ἔρμιος.," We found that the spectral index must be extremely soft as $\alpha_X \gtrsim$ 5 or 4 for $b_X = $ 1 or $b_X = b = 0.1$, respectively, in order that the extrapolated flux down to $\nu_T$ is equal to $L_{\rm ph, rec}$."
+ Therefore we can safely exclude the possibility that the nonthermal radiation producing the observed X-rays is tonizing the jet., Therefore we can safely exclude the possibility that the nonthermal radiation producing the observed X-rays is ionizing the jet.
+ Secondly we consider a possibility that a hot. UV-radiating star close to the SN 2002ap may ionize the jet.," Secondly we consider a possibility that a hot, UV-radiating star close to the SN 2002ap may ionize the jet."
+ It is expected that SN 2002ap occurred in a massive star forming region where young massive stars are clustering., It is expected that SN 2002ap occurred in a massive star forming region where young massive stars are clustering.
+ A close binary system ts a candidate for the type Ic supernova progenitors (Nomoto. Filippenko. Shigeyama 1990). and it may provide even stronger ionization source.," A close binary system is a candidate for the type Ic supernova progenitors (Nomoto, Filippenko, Shigeyama 1990), and it may provide even stronger ionization source."
+ However. the total ionization lummosity given in eq. (16))," However, the total ionization luminosity given in eq. \ref{eq:L_rec}) )"
+" is even larger by a factor of several than the ionization flux. ~107""?s7!. above the frequency ry~6«10"" Hz for the most luminous and hottest stars (Schaere de Koter 1997)."," is even larger by a factor of several than the ionization flux, $\sim 10^{49.5} \ \rm s^{-1}$, above the frequency $\nu_T \sim 6 \times 10^{15}$ Hz for the most luminous and hottest stars (Schaere de Koter 1997)."
+ It should be noted that this luminosity is for all direction. but only the radiation within the solid angle of the jet viewed from the tonizing star is available for the Jet ionization. which is expected to be a small fraction.," It should be noted that this luminosity is for all direction, but only the radiation within the solid angle of the jet viewed from the ionizing star is available for the jet ionization, which is expected to be a small fraction."
+ If the region around SN 2002ap is filled up by radiation field with ~Fig) to a radius of rj. the region should have a luminosity of at least Ars corresponding. bolometric luminosity of 2.1«10b16irlΟΕ assuming the spectral energy distribution of the most luminous O stars.," If the region around SN 2002ap is filled up by radiation field with $\sim F_{\rm ion}$ to a radius of $r_{\rm jet}$, the region should have a luminosity of at least $4 \pi F_{\rm ion} r_{\rm
+jet}^2$, corresponding bolometric luminosity of $2.1 \times 10^8 b_{-1}^{-2}
+\zeta_{-1}^{-1} \mu_{14}^{-1} t_{10}^{-3} (f_{\rm el}/0.3) L_\odot$, assuming the spectral energy distribution of the most luminous O stars."
+ Such a huge luminosity is apparently ruled out by the prediscovery image of the SN 2002ap field reported by Smartt et al. (, Such a huge luminosity is apparently ruled out by the prediscovery image of the SN 2002ap field reported by Smartt et al. (
+2002).,2002).
+ To conclude. ionization by nearby young stars Is Impossible.," To conclude, ionization by nearby young stars is impossible."
+ Since photoionization of the jet seems difficult. the only way to ionize the jet is enhanced collisional ionization by external heating.," Since photoionization of the jet seems difficult, the only way to ionize the jet is enhanced collisional ionization by external heating."
+ If the jet is generated at the central compact object. it might include a significant amount of radioactive nuclei such as ??Ni.," If the jet is generated at the central compact object, it might include a significant amount of radioactive nuclei such as $^{56}$ Ni."
+ Asymmetric explosion induced by the jet should also affect nucleosynthesis. and ??Ni production along the jet direction is enhanced (Nagataki 2000: Maeda et al.," Asymmetric explosion induced by the jet should also affect nucleosynthesis, and $^{56}$ Ni production along the jet direction is enhanced (Nagataki 2000; Maeda et al."
+ 2002)., 2002).
+ °°Ni decays by electron capture and gamma-ray emission to ??Co. with an exponential decay time scale of fy;=fj2/1n(2)8.5 d and decay energy is ew;=2.1 MeV. When the material is optically thick. the radioactive heat is quickly thermalized into optical radiation field. as generally seen for supernovae.," $^{56}$ Ni decays by electron capture and gamma-ray emission to $^{56}$ Co, with an exponential decay time scale of $t_{\rm Ni} = t_{1/2} /
+\ln (2) = 8.5$ d and decay energy is $\epsilon_{\rm Ni} = 2.1$ MeV. When the material is optically thick, the radioactive heat is quickly thermalized into optical radiation field, as generally seen for supernovae."
+ On the other hand. if the material is mildly optically thin. the gamma-rays emitted by decaying ?*Ni scatter electrons with a probability ~ and since the gamma-ray energy is comparable with the electronπα. rest mass. the scattered electrons acquire mildly relativistic speed and energy.," On the other hand, if the material is mildly optically thin, the gamma-rays emitted by decaying $^{56}$ Ni scatter electrons with a probability $\sim \tau_{\rm jet}$, and since the gamma-ray energy is comparable with the electron rest mass, the scattered electrons acquire mildly relativistic speed and energy."
+" Such high energy electrons would lose their energy by ionization loss in the jet plasma. with a time scale of where c, 1s the initial velocity of high energy electrons."," Such high energy electrons would lose their energy by ionization loss in the jet plasma, with a time scale of where $\upsilon_e$ is the initial velocity of high energy electrons."
+ Here we used the ionization loss formulae of Longair (1992) and the logarithmic factor is set to be 15., Here we used the ionization loss formulae of Longair (1992) and the logarithmic factor is set to be 15.
+ Therefore the energy deposited by radioactive gamma-rays Is used to ionize the jet material within the time scale of interest. giving an efficient ionization process.," Therefore the energy deposited by radioactive gamma-rays is used to ionize the jet material within the time scale of interest, giving an efficient ionization process."
+" When optical depth is very low. this process would be dominated by positrons emitted from decay of οσο, which has a longer exponential lifetime of fe,= 111.26 days and energy fraction given to positrons is of the total decay energy (Arnett 1979; Woosley. Pinto. Hartmann 1989)."," When optical depth is very low, this process would be dominated by positrons emitted from decay of $^{56}$ Co, which has a longer exponential lifetime of $t_{\rm Co} = $ 111.26 days and energy fraction given to positrons is of the total decay energy (Arnett 1979; Woosley, Pinto, Hartmann 1989)."
+ The tonizing balance is determined by the energy balance between radioactive heating and recombination cooling (see also Graham 1988) as: where fw; is the “CNi mass fraction in the jet. and τν the ionization. potential.," The ionizing balance is determined by the energy balance between radioactive heating and recombination cooling (see also Graham 1988) as: where $f_{\rm Ni}$ is the $^{56}$ Ni mass fraction in the jet, and $w$ the ionization potential."
+ The recombination rate. coefficient depends on the electron gas temperature. which is determined by balance between the radioactive heating and cooling processes.," The recombination rate coefficient depends on the electron gas temperature, which is determined by balance between the radioactive heating and cooling processes."
+" We can estimate the minimum amount of ??Ni by taking the minimum value of o as where the adopted value of aj, is minimum value of doubly ionized oxygen or carbon at temperature of ~10""K (Nahar Pradhan 1997; Nahar 1999) and too=w/(20 eV).", We can estimate the minimum amount of $^{56}$ Ni by taking the minimum value of $\alpha_{\rm rec}$ as where the adopted value of $\alpha_{\rm rec}$ is minimum value of doubly ionized oxygen or carbon at temperature of $\sim 10^4$ K (Nahar Pradhan 1997; Nahar 1999) and $w_{20} = w$ /(20 eV).
+ The jet may include significant amount of heavier nuclei that are difficult to ionize for a fixed value of fy. but on the other hand. it may also include considerable helium that is easier to ionize.," The jet may include significant amount of heavier nuclei that are difficult to ionize for a fixed value of $f_{\rm el}$, but on the other hand, it may also include considerable helium that is easier to ionize."
+ The helium could be mixed from remaining helium layer of the progenitor. or it may be newly synthesized.," The helium could be mixed from remaining helium layer of the progenitor, or it may be newly synthesized."
+ Production of helium is also enhanced along the jet direction in energetic jet-like nucleosynthesis (Maeda et al., Production of helium is also enhanced along the jet direction in energetic jet-like nucleosynthesis (Maeda et al.
+ 2002)., 2002).
+ We also note that highly ionized heavy nuclei. such as “Ni. should produce observable line emission in X-ray bands. and hence the observed weak X-ray flux gives a constraint on the species of ionized elements. (," We also note that highly ionized heavy nuclei, such as $^{56}$ Ni, should produce observable line emission in X-ray bands, and hence the observed weak X-ray flux gives a constraint on the species of ionized elements. ("
+See $5.2. for possible connection to X-ray line features often observed in GRB afterglows.),See \ref{section:X-ray-line} for possible connection to X-ray line features often observed in GRB afterglows.)
+ Whatever the jet composition ts. the above result indicates that. 1f the jet ts kept tonized by radioactive heating. it must have a considerable amount of ??Ni (mass fraction of order unity).," Whatever the jet composition is, the above result indicates that, if the jet is kept ionized by radioactive heating, it must have a considerable amount of $^{56}$ Ni (mass fraction of order unity)."
+ This estimate is. however. very uncertain especially about the composition of the jet. o and electron temperature.," This estimate is, however, very uncertain especially about the composition of the jet, $\alpha_{\rm rec}$ and electron temperature."
+ More sophisticated treatment is necessary to determine the ionization status. but it is beyond the scope of the paper.," More sophisticated treatment is necessary to determine the ionization status, but it is beyond the scope of the paper."
+ Therefore. it 1s difficult to conclude that the jet should be ionized. but it seems the best candidate of ionization process among others.," Therefore, it is difficult to conclude that the jet should be ionized, but it seems the best candidate of ionization process among others."
+ The jet may be ionized by gamma-rays of ?Ni decay leaking. from the photosphere of SN 2002ap. even if the jet does not have radioactive nuclei.," The jet may be ionized by gamma-rays of $^{56}$ Ni decay leaking from the photosphere of SN 2002ap, even if the jet does not have radioactive nuclei."
+ In fact. ionization. of helium envelope above the photosphere. which is required to explain the observed He lines in SN 1987A and type Ib supernovae. is ascribed to the leaking gamma-rays from photosphere (Graham 1988; Lucy 1991).," In fact, ionization of helium envelope above the photosphere, which is required to explain the observed He lines in SN 1987A and type Ib supernovae, is ascribed to the leaking gamma-rays from photosphere (Graham 1988; Lucy 1991)."
+ The mass of *°Ni produced by SN 2002ap Is estimated to be 0.0743:0.02M|. from the light curve modeling by Mazzali et al. (, The mass of $^{56}$ Ni produced by SN 2002ap is estimated to be $0.07 \pm 0.02 M_\odot$ from the light curve modeling by Mazzali et al. (
+2002). which is larger than the jet mass.,"2002), which is larger than the jet mass."
+ However. the efficiency for gamma-rays to hit the jet is reduced by the beaming factor. 5b. and it is further reduced by escaping fraction from photosphere.," However, the efficiency for gamma-rays to hit the jet is reduced by the beaming factor, $b$, and it is further reduced by escaping fraction from photosphere."
+ Although some supernovae. including SN 1998bw. showed evidence that a significant amount of gamma-rays are leaking in late phase (= 30 days) (Nakamura et al.," Although some supernovae, including SN 1998bw, showed evidence that a significant amount of gamma-rays are leaking in late phase $\gtrsim$ 30 days) (Nakamura et al."
+ 2001: Patat et al., 2001; Patat et al.
+ 2001). the leaking fraction should not be large in early phase of ~ 10 days. when the optical luminosity ts still glowing up by," 2001), the leaking fraction should not be large in early phase of $\sim$ 10 days, when the optical luminosity is still glowing up by"
+percentage of events for which the OT would not be observed even without dust extinction. considering the light curves of our simulated GRBs and fixing BR hnuuitiug inaenitudeo of linον=20.5.,"percentage of events for which the OT would not be observed even without dust extinction, considering the light curves of our simulated GRBs and fixing R limiting magnitude of $R_{lim}=20.5$."
+ Solid bold line in Fi, Solid bold line in Fig.
+e.o 8 shows that for the majority of data, \ref{senzaz1} shows that for the majority of data
+Uulike the previous systems. Fomalhaut bas not vet been subject to detailed uumerical simulations.,"Unlike the previous systems, Fomalhaut has not yet been subject to detailed numerical simulations."
+ The Foimallhaut system also differs markedly in that 850 SSCUBA observatious suggest the system is seen nearly edge-on aud consists of a dusty. torus. rather than a thin disk (Hollandetal.1998).," The Fomalhaut system also differs markedly in that 850 SCUBA observations suggest the system is seen nearly edge-on and consists of a dusty torus, rather than a thin disk \citep{holland98}."
+. The unproved spatial resolution offered by the 150 SSCUBA images of Hollandetal.(2003) confirm the torus structure. while revealing a previously uuseen arc of emission near or withliu the torus.," The improved spatial resolution offered by the 450 SCUBA images of \citet{holland03} confirm the torus structure, while revealing a previously unseen arc of emission near or within the torus."
+ This departure from a uuilorm structure is strougly suggestive of a planetary. presence., This departure from a uniform structure is strongly suggestive of a planetary presence.
+ The generation of a single arc of emission can be achieved by a 1:1 resonance. as suggested by Hollandetal.(2003).," The generation of a single arc of emission can be achieved by a 1:1 resonance, as suggested by \citet{holland03}."
+. Inspection of results for resonance occupaucy obtained iu the construction of our synthetic catalogue. however. indicate that the 1:1 resonance is difficult to populate under normal circumstances.," Inspection of results for resonance occupancy obtained in the construction of our synthetic catalogue, however, indicate that the 1:1 resonance is difficult to populate under normal circumstances."
+ An example is shown in Figure 13.. where a Jupiter mass planet populates the 1:1 resonance when parent bodies exist interior to the planet. but does not populate the 1:1 resonance when the pareut bodies are all external to the planet.," An example is shown in Figure \ref{fig:11resonance}, where a Jupiter mass planet populates the 1:1 resonance when parent bodies exist interior to the planet, but does not populate the 1:1 resonance when the parent bodies are all external to the planet."
+ We consider it uulikely that parent bodies with sslightly less than the planets wwould survive for the required length. of time., We consider it unlikely that parent bodies with slightly less than the planet's would survive for the required length of time.
+" As an alternative explanation of the single are feature. we find that systems iucludingo a massive CÀZ,(abt2M{1 j,,). rnoderately eccentric 0.5) planet can iuduce a sugle emission arc over a backgrouud ring by trapping many particles in the n: resonances. where n71."," As an alternative explanation of the single arc feature, we find that systems including a massive $M_{pl} \geq M_{Jup}$ ), moderately eccentric $0.3 \leq e_{pl} \leq 0.5$ ) planet can induce a single emission arc over a background ring by trapping many particles in the $n:1$ resonances, where $n>1$."
+" Alter examining our synthetic catalogue. we chose a model with the followiug parameters: AL,—2.34AL. (appropriateforanΑΝstarsuchasFomalbaut.BarradoyNavascuesetal.1997).. Al=2M yup. ερ=OL. and ay~39 AU."," After examining our synthetic catalogue, we chose a model with the following parameters: $M_{\star} = 2.3 M_{\odot}$ \citep[appropriate for an A3V star such as Fomalhaut -- ][]{bar97}, $M_{pl} = 2 M_{Jup}$ , $e_{pl}=0.4$, and $a_{pl}
+\sim 59$ AU."
+" Parent bodies were distributed. with initial orbital parameters in the following ranges: 0<ey,«0.3. 100<ap,180 AU. 0<7,<25.."," Parent bodies were distributed with initial orbital parameters in the following ranges: $0<e_{pb}<0.3$, $100<a_{pb}<180$ AU, $0<i_{pb}<25$."
+. The higher parent body inclinatious naturally. generates a torus structure. which is believed to exist in the Fomalhaut system.," The higher parent body inclinations naturally generates a torus structure, which is believed to exist in the Fomalhaut system."
+ The simulation used 1000 test particles with 9=0.05. since tle mass of dust erains in the Fomalhaut system with diameters > ((corresponding to values of 8> 0.1) is believed to be negligible (Dentetal.2000).," The simulation used 1000 test particles with $\beta = 0.05$, since the mass of dust grains in the Fomalhaut system with diameters $>$ (corresponding to values of $\beta > 0.1$ ) is believed to be negligible \citep{dent00}."
+. The results of our simulations. which ran until no particles remained after 212 million years. are shown iu Figure 1L.," The results of our simulations, which ran until no particles remained after 212 million years, are shown in Figure \ref{fig:fomalhaut}."
+ The simulated observatious bear a close resemblance to the craw’ image of Fomalhaut presented in Hollandetal.(2003)., The simulated observations bear a close resemblance to the `raw' image of Fomalhaut presented in \citet{holland03}.
+. This planetary configurationOm does egenerates a dust distribution which does uot rotate with the planet. but appears fixed [rom a viewpoint external to the system over au orbital period.," This planetary configuration does generates a dust distribution which does not rotate with the planet, but appears fixed from a viewpoint external to the system over an orbital period."
+ Thus. observations of Fomalhaut over time would slow no eliauge iu emissiou if this model is correct.," Thus, observations of Fomalhaut over time would show no change in emission if this model is correct."
+ The effects of planetary phase are shown in Figure 15.., The effects of planetary phase are shown in Figure \ref{fig:fomalphase}.
+ Whilst the planetary configuration we have preseuted here displays a close similarity to the observations of Fomalhaut to date. as noted above results from our syuthetic catalogue suggest tliat," Whilst the planetary configuration we have presented here displays a close similarity to the observations of Fomalhaut to date, as noted above results from our synthetic catalogue suggest that"
+primary at birth.,primary at birth.
+" The initial conditions described in section (2) use q = wsEi as a definition of the mass-ratio, but since the gainer has Mabecome the most massive component of the Algol-system, we use q = “4 as the definition of mass-ratio of an Algol-system."," The initial conditions described in section \ref{sec_Initial}) ) use q = ${M_{g}\over M_{d}}$ as a definition of the mass-ratio, but since the gainer has become the most massive component of the Algol-system, we use q = ${M_{d}\over M_{g}}$ as the definition of mass-ratio of an Algol-system."
+ The observedMg Algols combine a large fraction of systems where Algol characteristics are produced, The observed Algols combine a large fraction of systems where Algol characteristics are produced
+"orientations probed, as shown in the third row, left panel of Figure 7..","orientations probed, as shown in the third row, left panel of Figure \ref{fig:dist2}."
+" Here, a represents the angle between the angular momentum vectors of 92 and the IBH at the start of the simulations."," Here, $\alpha$ represents the angle between the angular momentum vectors of S2 and the IBH at the start of the simulations."
+" This incidentally argues against the possibility that the observed deviations are due to Kozai oscillations, since the mechanism requires large relative inclinations to operate."," This incidentally argues against the possibility that the observed deviations are due to Kozai oscillations, since the mechanism requires large relative inclinations to operate."
+" However, fits with the same sets of parameters (q,e,a) but different a can have different values of x?."," However, fits with the same sets of parameters $(q,\,e,\,a)$ but different $\alpha$ can have different values of $\chi^2$."
+" This means that the knowledge of (q,e,a) is insufficient to predict the reduced x? but information on the sky position is necessary."," This means that the knowledge of $(q,\,e,\,a)$ is insufficient to predict the reduced $\chi^2$ but information on the sky position is necessary."
+" Finally, the right panel in the third row of Figure 7 shows that the minimum 3D distance between S2 and the IBH also correlates with the reduced x?."," Finally, the right panel in the third row of Figure \ref{fig:dist2} shows that the minimum 3D distance between S2 and the IBH also correlates with the reduced $\chi^2$."
+" In general it holds that the smaller the minimum distance, the worse the corresponding fit."," In general it holds that the smaller the minimum distance, the worse the corresponding fit."
+" Clearly, this parameter is not independent of the semi-major axis."," Clearly, this parameter is not independent of the semi-major axis."
+ The initial parameters adopted for the black hole binary are not sampled homogeneously., The initial parameters adopted for the black hole binary are not sampled homogeneously.
+" This is obvious for the first three panels in Figure 7 showing the reduced x? as a function of the binary parameters (q,e,a)."," This is obvious for the first three panels in Figure \ref{fig:dist2} showing the reduced $\chi^2$ as a function of the binary parameters $(q,\,e,\,a)$."
+ But it also holds (and is less obvious) for the plot investigating the minimum distance between 82 and the IBH., But it also holds (and is less obvious) for the plot investigating the minimum distance between S2 and the IBH.
+ A comparison of the goodness of fit for the runs starting with the IBH at periapsis and apoapsis is shown in Figure 9.., A comparison of the goodness of fit for the runs starting with the IBH at periapsis and apoapsis is shown in Figure \ref{fig:cfr}. .
+" In both cases, the semi-major axis of the black hole orbit is a=10mpc."," In both cases, the semi-major axis of the black hole orbit is $a=10\mpc$."
+ We find a modest worsening of the x? in the case of an IBH initially at the apoapsis of its orbit., We find a modest worsening of the $\chi^2$ in the case of an IBH initially at the apoapsis of its orbit.
+" This can be attributed to the fact that S2's apoapsis, where the star spends most of its time, is about 10mpc."," This can be attributed to the fact that S2's apoapsis, where the star spends most of its time, is about $10\mpc$."
+" Finally, we also investigated the minimum time required for the IBH to become detectable."," Finally, we also investigated the minimum time required for the IBH to become detectable."
+" For this purpose, we repeated the orbital fits for a few cases assuming that the observations span 10,15, 20,25, 30, 35, 40, 45 or 50 years (Figure 10))."," For this purpose, we repeated the orbital fits for a few cases assuming that the observations span 10,15, 20,25, 30, 35, 40, 45 or 50 years (Figure \ref{fig:time}) )."
+" Our initial conditions are such that the first periapse passage of S2 happens after 10 years, the second after 26 years and the third after 42 years."," Our initial conditions are such that the first periapse passage of S2 happens after 10 years, the second after 26 years and the third after 42 years."
+" Figure 10 shows that the reduced X? starts to increase beyond our threshold of 1.22 after the second periapsis passage for fits that show a large reduced x? after 50 simulated years (red/dashed, blue/long-dashed and green/solid curves)."," Figure \ref{fig:time} shows that the reduced $\chi^2$ starts to increase beyond our threshold of 1.22 after the second periapsis passage for fits that show a large reduced $\chi^2$ after 50 simulated years (red/dashed, blue/long-dashed and green/solid curves)."
+ Only for fits that after 50 simulated years have a reduced χ΄<3 is the threshold passed after the third periapse passage(black curve)., Only for fits that after 50 simulated years have a reduced $\chi^2 \simless 3$ is the threshold passed after the third periapse passage(black curve).
+ The discrete nature of periapse passages also is the reason why, The discrete nature of periapse passages also is the reason why
+"For direct comparison purposes we also show MDF data for the halo stars in three locations of the NGC 5128 halo, from Harris et al. (1999, 2000,, 2002)).","For direct comparison purposes we also show MDF data for the halo stars in three locations of the NGC 5128 halo, from Harris et al. \cite{har99}, \cite{har00}, \cite{har02}) )."
+" These were all taken with the HST WFPC2 camera in (V,7) and have rather similar limiting absolute magnitudes to our M87 ACS-based photometry, though they are less affected by crowding."," These were all taken with the HST WFPC2 camera in $(V,I)$ and have rather similar limiting absolute magnitudes to our M87 ACS-based photometry, though they are less affected by crowding."
+ The MDFs were derived in all cases with the same RGB grid of tracks and interpolation code., The MDFs were derived in all cases with the same RGB grid of tracks and interpolation code.
+" The mid-to-outer fields at projected galactocentric distances of 21 and 31 kpc have MDFs that are virtually identical to each other and are combined in the upper panel of Figure 12,, while the inner-halo 8 kpc field is shown by itself in the middle panel."," The mid-to-outer fields at projected galactocentric distances of 21 and 31 kpc have MDFs that are virtually identical to each other and are combined in the upper panel of Figure \ref{fehhisto}, while the inner-halo 8 kpc field is shown by itself in the middle panel."
+ Our corrected MDF for the M87 inner halo — at a mean projected distance of 10 kpc - clearly resembles the 8-kpc NGC 5128 field more closely than the outer fields., Our corrected MDF for the M87 inner halo – at a mean projected distance of 10 kpc – clearly resembles the 8-kpc NGC 5128 field more closely than the outer fields.
+" Over the range [m/H] «—0.3 where we can make the comparison, we conclude that the inner halos of both these giant ellipticals have basically similar MDFs that are broad, predominantly metal-rich, and with very small numbers of metal-poor stars."," Over the range [m/H] $< -0.3$ where we can make the comparison, we conclude that the inner halos of both these giant ellipticals have basically similar MDFs that are broad, predominantly metal-rich, and with very small numbers of metal-poor stars."
+" The MDF for M87 reach a peak near [m/H] ~—0.4, but for the present this metallicity should be viewed as a lower limit to the peak value."," The MDF for M87 reach a peak near [m/H] $\sim -0.4$, but for the present this metallicity should be viewed as a lower limit to the peak value."
+" We have used HST Archive data for an unusually deep set of F606W,F814W ACS images to probe the red-giant stellar population in the inner halo of M87."," We have used HST Archive data for an unusually deep set of $F606W, F814W$ ACS images to probe the red-giant stellar population in the inner halo of M87."
+" Although the crowding levels of the faint halo stars in this field are severe at best, the regions from 115""—155"" (9.3 to 12.5 kpc projected galactocentric distance) give a useful first look at the brightest 1.5 magnitudes of the red-giant branch."," Although the crowding levels of the faint halo stars in this field are severe at best, the regions from $115'' - 155''$ (9.3 to 12.5 kpc projected galactocentric distance) give a useful first look at the brightest 1.5 magnitudes of the red-giant branch."
+" This corresponds to distances of 1.4R,—1.9R,, where R, is the effective radius in the I band of 81"" (6.3 kpc) (Zeilinger, 1993))."," This corresponds to distances of $1.4 R_e -
+1.9 R_e$, where $R_e$ is the effective radius in the I band of 81” (6.3 kpc) (Zeilinger, \cite{zei93}) )."
+ We have used this material to obtain a preliminary TRGB-calibrated distance to M87 for the first time., We have used this material to obtain a preliminary TRGB-calibrated distance to M87 for the first time.
+" We find a distance d=(16.7+0.9) Mpc, in good agreement with the few other relatively direct methods available, including the planetary nebula luminosity function, the Cepheid-calibrated SBF method, and the linear diameters of globular clusters."," We find a distance $d = (16.7 \pm 0.9)$ Mpc, in good agreement with the few other relatively direct methods available, including the planetary nebula luminosity function, the Cepheid-calibrated SBF method, and the linear diameters of globular clusters."
+ Photometry of a more outlying halo field than the one we are able to study here has the potential to give a more precise TRGB determination and would represent the most direct possible distance calibration to the Virgo members at the present time., Photometry of a more outlying halo field than the one we are able to study here has the potential to give a more precise TRGB determination and would represent the most direct possible distance calibration to the Virgo members at the present time.
+" As a second result of our study, we use the brightest 1 magnitude of the RGB stars to measure the metallicity distribution of the M87 inner halo for the first time."," As a second result of our study, we use the brightest 1 magnitude of the RGB stars to measure the metallicity distribution of the M87 inner halo for the first time."
+" We find that it is very broad and predominantly metal-rich, with a peak that may lie very near the completeness limit of our V—band data but is at least as metal-rich as [m/H] ~—0.4."," We find that it is very broad and predominantly metal-rich, with a peak that may lie very near the completeness limit of our $V-$ band data but is at least as metal-rich as [m/H] $\simeq -0.4$."
+" In general, the shape of the MDF strongly resembles the “8-kpc field"" for the inner halo of NGC 5128, the nearby giant elliptical in the Centaurus group."," In general, the shape of the MDF strongly resembles the “8-kpc field” for the inner halo of NGC 5128, the nearby giant elliptical in the Centaurus group."
+" In our view,these results strongly reinforce"," In our view,these results strongly reinforce"
+parameter fitted was the pulse phase.,parameter fitted was the pulse phase.
+ Observations made simultaneously at 3 GHz were used to estimate the dispersion., Observations made simultaneously at 3 GHz were used to estimate the dispersion.
+ The integrated pulse power was not flux calibrated but it is believed that the 50 cm receiver noise is quite stable so the signal to noise ratio is proportional to the pulse flux., The integrated pulse power was not flux calibrated but it is believed that the 50 cm receiver noise is quite stable so the signal to noise ratio is proportional to the pulse flux.
+" As this source is nearby the dispersion is relatively small, even at 685 MHz, but it is not negligible."," As this source is nearby the dispersion is relatively small, even at 685 MHz, but it is not negligible."
+ The autocovariance of is shown in the left panel of Figure 17., The autocovariance of $1/\sqrt{SN}$ is shown in the left panel of Figure 17.
+ It shows1//.$N a quasi-exponential covariance with a time scale of about 20 min (which is the diffractive scale) and very little white noise (which would appear as a delta function at the origin)., It shows a quasi-exponential covariance with a time scale of about 20 min (which is the diffractive scale) and very little white noise (which would appear as a delta function at the origin).
+ The autocovariance of the timing residuals corrected for dispersion is shown in the right panel of Figure 17., The autocovariance of the timing residuals corrected for dispersion is shown in the right panel of Figure 17.
+ This covariance shows a clear white noise component and an exponential component with a time scale of about 30 min., This covariance shows a clear white noise component and an exponential component with a time scale of about 30 min.
+ This is not significantly different from the diffractive scale and is undoubtedly associated with the diffractive intensity variations., This is not significantly different from the diffractive scale and is undoubtedly associated with the diffractive intensity variations.
+ The white noise carries about the same variance as the 30 min variations., The white noise carries about the same variance as the 30 min variations.
+" It should be noted that the rms white noise in these observations, about 1 us, is much larger than the white noise in the normal observations of the Parkes Pulsar Timing Array (PPTA) for several reasons."," It should be noted that the rms white noise in these observations, about 1 $\mu$ s, is much larger than the white noise in the normal observations of the Parkes Pulsar Timing Array (PPTA) for several reasons."
+" The basic PPTA timing is done at 1400 MHz where the effects of the ISM are much smaller, with a much broader bandwidth and a much longer integration time, so the receiver noise is much lower."," The basic PPTA timing is done at 1400 MHz where the effects of the ISM are much smaller, with a much broader bandwidth and a much longer integration time, so the receiver noise is much lower."
+ The cross correlation of the residuals with 1//SN is shown in Figure 18 before and after the scattering correction., The cross correlation of the residuals with $1/\sqrt{SN}$ is shown in Figure 18 before and after the scattering correction.
+ One can see that in advance of the scattering correction the correlation (shown dotted) was about50%., One can see that in advance of the scattering correction the correlation (shown dotted) was about.
+. If one corrects for the white noise component in the residuals the correlation of the diffractive component is about as in the simulations., If one corrects for the white noise component in the residuals the correlation of the diffractive component is about as in the simulations.
+ After the scattering correction the correlation (shown solid) is zero as expected., After the scattering correction the correlation (shown solid) is zero as expected.
+ The scattering correction reduces the total variance of the residuals by about25%., The scattering correction reduces the total variance of the residuals by about.
+. However one can see that the cross correlation persists as a negative component away from the origin., However one can see that the cross correlation persists as a negative component away from the origin.
+ This could perhaps be removed using a more sophisticated removal., This could perhaps be removed using a more sophisticated removal.
+" Rather than subtracting a constant times 1/VS.N one could subtract the reference after being filtered with an FIR filter, as would be done in an adaptive echo canceller."," Rather than subtracting a constant times $1/\sqrt{SN}$ one could subtract the reference after being filtered with an FIR filter, as would be done in an adaptive echo canceller."
+" It has been noted [Verbiest et al, 2009] that the errors in the residuals of J0437—4715 do not behave like white noise."," It has been noted [Verbiest et al, 2009] that the errors in the residuals of $-$ 4715 do not behave like white noise."
+ In particular when residuals sampled at 5 minute intervals are averaged the error on the average does not shrink like the square root of the number of residuals averaged., In particular when residuals sampled at 5 minute intervals are averaged the error on the average does not shrink like the square root of the number of residuals averaged.
+ This is clearly because half the variance is carried in the diffractive component which has a correlation time considerably longer than 5 minutes., This is clearly because half the variance is carried in the diffractive component which has a correlation time considerably longer than 5 minutes.
+ In this pulsar we are able to observe and correct diffractive fluctuations in the TOAs., In this pulsar we are able to observe and correct diffractive fluctuations in the TOAs.
+ The refractive fluctuations were negligible but it seems likely that they too could have been corrected the same way., The refractive fluctuations were negligible but it seems likely that they too could have been corrected the same way.
+" We attempted to analyze the refractive scattering contribution of the pulsar J19394- 2134, as was done by Lestrade et al. ["," We attempted to analyze the refractive scattering contribution of the pulsar $+$ 2134, as was done by Lestrade et al. ["
+1998] however the pulsar is not sampled adequately in the normal PPTA observations to resolve the 2.5 day time scale of the refractive scintillation.,1998] however the pulsar is not sampled adequately in the normal PPTA observations to resolve the 2.5 day time scale of the refractive scintillation.
+ It would be very useful to be able to predict the TOA variations discussed above for a given pulsar at a given frequency., It would be very useful to be able to predict the TOA variations discussed above for a given pulsar at a given frequency.
+" Since we lack an analytical theory, we have created a heuristic model with which we can scale the results of simulations to match the observing parameters."," Since we lack an analytical theory, we have created a heuristic model with which we can scale the results of simulations to match the observing parameters."
+" First we ran simulations over the range of strengths of scattering (m2) accessible to our computing platforms: 3, 5, 10, 20, 30, 60, and 100."," First we ran simulations over the range of strengths of scattering $m_b^2$ ) accessible to our computing platforms: 3, 5, 10, 20, 30, 60, and 100."
+" Then we compared the behavior of the parameters for which we have theoretical scaling models, with the simulations."," Then we compared the behavior of the parameters for which we have theoretical scaling models, with the simulations."
+ This comparison validates both the simulations and the scaling models., This comparison validates both the simulations and the scaling models.
+ The fundamental spatial scales are: the e~°> scale of the electric field (so); the diffractive e! scale of intensity (8418): and the refractive e! scale of intensity (Sret)., The fundamental spatial scales are: the $e^{-0.5}$ scale of the electric field $s_0$ ); the diffractive $e^{-1}$ scale of intensity $s_\text{dif}$ ); and the refractive $e^{-1}$ scale of intensity $s_\text{ref}$ ).
+" The scales of intensity were measured at their half power points, i.e. Τ=(8m?—1)/4 and (m?—1)/4, and then corrected to the e! values using the theoreticalcurve shape."," The scales of intensity were measured at their half power points, i.e. $\Gamma = (3m^2-1)/4$ and $(m^2-1)/4$, and then corrected to the $e^{-1}$ values using the theoreticalcurve shape."
+ The scale of the field is measured from the autocovariance of the field at εδ., The scale of the field is measured from the autocovariance of the field at $e^{-0.5}$.
+ The results are plotted vs m? in Figure 19 as symbols with error bars., The results are plotted vs $m_b^2$ in Figure 19 as symbols with error bars.
+ The error bars are derived from multiple simulations., The error bars are derived from multiple simulations.
+ For a given m? we can use equation (3) to find so., For a given $m_b^2$ we can use equation (3) to find $s_0$.
+" This expression, which is valid in any strength of scattering, is plotted as a solid line and agrees well with the simulations."," This expression, which is valid in any strength of scattering, is plotted as a solid line and agrees well with the simulations."
+ The strong scattering approximation for diffractive scale sai;=so agrees less well but improves in stronger scattering., The strong scattering approximation for diffractive scale $s_\text{dif} = s_0$ agrees less well but improves in stronger scattering.
+" The strong scattering approximation for refractive scale sje;=77/80, which is also plotted as a solid line, has weak agreement with the simulations."," The strong scattering approximation for refractive scale $s_\text{ref} = r_\text{f}^2 / s_0$, which is also plotted as a solid line, has weak agreement with the simulations."
+" A heuristic expression for s,,; derived from numerical solutions to the moment equations (Goodman and Narayan, 2006), which is plotted as dashed line, agrees much better."," A heuristic expression for $s_\text{ref}$ derived from numerical solutions to the moment equations (Goodman and Narayan, 2006), which is plotted as dashed line, agrees much better."
+ It is reassuring that the numerical solution agrees with the simulations for νου., It is reassuring that the numerical solution agrees with the simulations for $s_\text{ref}$ .
+" The intensity variance is a fundamental parameter,"," The intensity variance is a fundamental parameter,"
+"lines from the ((0-0) to the ((4-0) band were fitted simultaneously, while the ((5-0) band was ignored during the fitting process due to the presence of a clear blending.","lines from the (0-0) to the (4-0) band were fitted simultaneously, while the (5-0) band was ignored during the fitting process due to the presence of a clear blending."
+" We got N(CO)=14.74+ 0.07, b=1.1€0.1 kmss! and T4 (CO)log= K. For the remaining two systems (toward SDSS J08572641855247.807 and JJ170542+354340), the rotational levels lines are hardly and the signal-to-noise ratio is not high enough to measure the column densities in each rotational level independently."," We got $\log N($ $)=14.74\pm0.07$ , $b=1.1\pm0.1$ $^{-1}$ and $T_{\rm ex}($ $)=7.8^{+0.7}_{-0.6}$ K. For the remaining two systems (toward SDSS J085726+185524 and J170542+354340), the rotational levels lines are hardly and the signal-to-noise ratio is not high enough to measure the column densities in each rotational level independently."
+" Instead, we used a technique similar to that presented in Burghetal.(2007) and Prochaskaetal.(2009):: a single excitation temperature was used as an external parameter to model the CO profile directly, again, using a Boltzmann distribution of the rotational level populations."," Instead, we used a technique similar to that presented in \citet{Burgh07} and \citet{Prochaska09}: a single excitation temperature was used as an external parameter to model the CO profile directly, again, using a Boltzmann distribution of the rotational level populations."
+ The best fit-model was chosen from the minimum X? obtained for a grid of total column densities and excitation temperatures (see right-hand side panels of Figs., The best fit-model was chosen from the minimum $\chi^2$ obtained for a grid of total column densities and excitation temperatures (see right-hand side panels of Figs.
+ 2 and 3))., \ref{0857} and \ref{1705}) ).
+ The excitation temperatures for the five CO-bearing systems are given in Table 1.., The excitation temperatures for the five CO-bearing systems are given in Table \ref{tco}.
+" At high redshift, the excitation of CO is dominated by radiative excitation as predicted in diffuse interstellar clouds (Warinetal.,1996;Burgh 2007)."," At high redshift, the excitation of CO is dominated by radiative excitation as predicted in diffuse interstellar clouds \citep{Warin96,Burgh07}."
+". Indeed, the excitation temperatures we measured are well above the mean temperature measured in the Galaxy for similar CO column densities (Τε)= =3.6 K,2007)."," Indeed, the excitation temperatures we measured are well above the mean temperature measured in the Galaxy for similar CO column densities \citep[$\avg{T_{\rm ex}}= =3.6 K,."
+" This is further supported by the low volume density of the gas derived from the analysis of H5 and C? lines in two of these high-z absorption systems (Srianandetal.,2008;Noterdaemeetal., 2010)."," This is further supported by the low volume density of the gas derived from the analysis of $_2$ and $^{\rm 0}$ lines in two of these $z$ absorption systems \citep{Srianand08, Noterdaeme10co}."
+". Therefore, T.,(CO) must be a good proxy for Tcmp at high redshift."," Therefore, $T_{\rm ex}$ (CO) must be a good proxy for $T_{\rm CMB}$ at high redshift."
+" The CMB temperature derived from the rotational excitation of CO in five absorption systems (three studied here plus two previously published, see Table 1)) are presented in Fig."," The CMB temperature derived from the rotational excitation of CO in five absorption systems (three studied here plus two previously published, see Table \ref{tco}) ) are presented in Fig."
+ 4 together with measurements and upper-limits obtained from the analysis of the populations of the fine-structure energy levels of atomic carbon or from the S-Z effect ingalaxy clusters., \ref{TCMBz} together with measurements and upper-limits obtained from the analysis of the populations of the fine-structure energy levels of atomic carbon or from the S-Z effect ingalaxy clusters.
+ An upper-limit has also been obtained from the analysis of millimetre absorption lines from different molecules in the, An upper-limit has also been obtained from the analysis of millimetre absorption lines from different molecules in the
+we obtain the following integral representation Lor o: The variation of the (transmission probability ΤΙ=exp(—2c) is represented. as a function of the parameter ¢=fine+h?fe and for different values of the temperature and particle energy ralio 2a7T7/a7 in Fig.,"we obtain the following integral representation for $\sigma $: The variation of the transmission probability $\left| \mathcal
+T\right| ^{2}=\exp
+\left( -2\sigma \right) $ is represented, as a function of the parameter $%
+\zeta =\sqrt{m_{e}^{2}+k_{\perp }^{2}}/\omega and for different values of the temperature and particle energy ratio $2\pi
+^{2}T^{2}/\omega ^{2}$ in Fig."
+ 2., 2.
+ For large values of the particle energy. ho40 and the transmission probability is equal to 1. ΤΙ—1.," For large values of the particle energy, $\omega >>\sqrt{%
+m_{e}^{2}+k_{\perp }^{2} , $\zeta \rightarrow 0$ and the transmission probability is equal to $1$, $\left| \mathcal
+T\right| ^{2}\rightarrow 1$."
+ The transmission probability also increases with the temperature of the electrosphere., The transmission probability also increases with the temperature of the electrosphere.
+ In the process of particle production in the electrosphere of quark stars due to the (unuelling from the negative energv state. a positive energy particle (an electron) is created. leavinge a hole in the negativee enereve. continuum. which can be taken to be an anti-particle (a positron) moving in a direction opposite to the direction of (he created particle.," In the process of particle production in the electrosphere of quark stars due to the tunnelling from the negative energy state, a positive energy particle (an electron) is created, leaving a hole in the negative energy continuum, which can be taken to be an anti-particle (a positron) moving in a direction opposite to the direction of the created particle."
+ The created pair of particles is characterized by the energy. (frequency) w., The created pair of particles is characterized by the energy (frequency) $\omega $.
+ For such a pair. one cannot. strictly speaking. specilv a particular point as the location where (he pair is produced 19388).," For such a pair, one cannot, strictly speaking, specify a particular point as the location where the pair is produced \citep{Wa88}."
+. One can only sav that the pair of particles begins to emerge between the point 2 and z., One can only say that the pair of particles begins to emerge between the point $z_{-}$ and $z_{+}$.
+ Nevertheless. one can associate the point z. which is the solution of the equation w@=V(z). as Che location in the vicinity of which a pair of particles is produced.," Nevertheless, one can associate the point $z$, which is the solution of the equation $\omega =V(z)$, as the location in the vicinity of which a pair of particles is produced."
+ With (his approximate association. (he enerev of the produced particle is then identified by an approximate location. and the energy interval also can be approximately related to a spatial interval at which the pair of particles is produced via the relation ο...," With this approximate association, the energy of the produced particle is then identified by an approximate location, and the energy interval also can be approximately related to a spatial interval at which the pair of particles is produced via the relation $d\omega =\left( \partial V/\partial z\right) dz=eE\left(
+z,T\right) dz$ ."
+ Moreover. we shall introduce polar coordinates in the momentum space so that and fhy=ksin@.," Moreover, we shall introduce polar coordinates in the momentum space so that $k_{x}=k\cos \theta $ and $k_{y}=k\sin \theta $."
+ Therefore the electron-positron pair production rate »-. giving the nunmber of electron-positron pairs created per unit time ancl per unit volume by (he electric field of the electrosphere at the surface of quark stars. (V)/MXrNg Nz. can be written as where 5s describes the spin degrees of freedom of the produced particles (s=1 for bosons and s2 for fermions).," Therefore the electron-positron pair production rate $\dot{n_{\pm
+}}$, giving the number of electron-positron pairs created per unit time and per unit volume by the electric field of the electrosphere at the surface of quark stars, $%
+\left\langle N\right\rangle /\Delta t\Delta x\Delta y\Delta z, can be written as where $s$ describes the spin degrees of freedom of the produced particles $s=1$ for bosons and $s=2$ for fermions)."
+panels of Fig.,panels of Fig.
+ 4. show the Stokes-V images without and with pointing and squint corrections., \ref{EX1} show the Stokes-V images without and with pointing and squint corrections.
+ The algorithm was also tested for Stokes-I and -V imaging using VLA 1.4 GHz observations of the superthin galaxy 1C2233 (2)., The algorithm was also tested for Stokes-I and -V imaging using VLA 1.4 GHz observations of the superthin galaxy IC2233 \citep{Matthews_Uson_08}.
+ The field contains two strong sources mJy/beam and ~145 mJy/beam) on opposite sides of the pointing center. located at positions of ~75% and ~35% primary beam response levels respectively.," The field contains two strong sources mJy/beam and $\sim145$ mJy/beam) on opposite sides of the pointing center, located at positions of $\sim 75$ and $\sim 35$ primary beam response levels respectively."
+ The observations were made in spectral mode with channels of width ~24 kHz for a total of ~11.6 hours in 2 passes with well distributed uv-coverage., The observations were made in spectral mode with channels of width $\sim 24$ kHz for a total of $\sim 11.6$ hours in 2 passes with well distributed uv-coverage.
+ The line-free channels (11 from the second “IF pat”) were used for the tests described here., The line-free channels (11 from the second “IF pair”) were used for the tests described here.
+ The aperture illumination pattern for each antenna was assumed to be the same and computed using the model for, The aperture illumination pattern for each antenna was assumed to be the same and computed using the model for
+line from the T Tau complex demonstrating that in this svstem most of the enission originates in jets/outflows with neon atoms partly ionizecl by stellar A-ravs.,line from the T Tau complex demonstrating that in this system most of the emission originates in jets/outflows with neon atoms partly ionized by stellar X-rays.
+ Our aim here is to enlarge the sample of spectrally resolved lines and understand in which svstems emission originates in a disk., Our aim here is to enlarge the sample of spectrally resolved lines and understand in which systems emission originates in a disk.
+ In the following subsections we describe our observational campaign aud (he data reduction (Sects. ??.. ??.. ?7)).," In the following subsections we describe our observational campaign and the data reduction (Sects. \ref{sect:targets}, , \ref{sect:observations}, \ref{sect:dataredu}) )."
+ Observations were carried out. wilh the high-resolution (R~30.000) spectrograph. VISIR mounted on the VLT telescopeMelipal," Observations were carried out with the high-resolution $\sim$ 30,000) spectrograph VISIR mounted on the VLT telescope (Sect."
+ (Sect. ?? for more details)., \ref{sect:observations} for more details).
+ In addition to the already published Spitzer/IRS spectra. we reduced and present here new archival IRS spectra that aid the interpretation of the VISIR observations (Sect. ??)).," In addition to the already published Spitzer/IRS spectra, we reduced and present here new archival IRS spectra that aid the interpretation of the VISIR observations (Sect. \ref{sect:dataredu}) )."
+ Because of the lower-sensitivitv of VISIR. with respect to Spitzer/1hS. we restricted our sample to disks with bright/unresolved lines (Spitzer [lixes greater than 1 uflux)).," Because of the lower-sensitivity of VISIR with respect to Spitzer/IRS, we restricted our sample to disks with bright/unresolved lines (Spitzer fluxes greater than $^{-14}$ )."
+ In addition. we selected disks brighter than ~50mmJv in the continuum because Larget acquisition for fainter objects becomes challenging with VISIR.," In addition, we selected disks brighter than $\sim$ mJy in the continuum because target acquisition for fainter objects becomes challenging with VISIR."
+ The 6 targets we selected for this campaien are TW Ilva. CS Cha. VW Cha. T Cha. Sz 73. and Sz 102 (see Table D. and 4 for their main properties).," The 6 targets we selected for this campaign are TW Hya, CS Cha, VW Cha, T Cha, Sz 73, and Sz 102 (see Table \ref{table:prop} and \ref{table:modelparameters}
+ for their main properties)."
+ At the beginning of the second night we also observed the double-Iined spectroscopic binary WD 34700AÀ. which is the brightest infrawed T Tauri couple in a quadruple svstem (Sterziketal.2005).," At the beginning of the second night we also observed the double-lined spectroscopic binary HD 34700A, which is the brightest infrared T Tauri couple in a quadruple system \citep{sterzik05}."
+. Although WD 347004 did not have a published line detection. its near- and mid-infrared spectra present unusually strong PAIT emission bands (e.g.. Smithetal. 2004)). possibly hinting to a high stellar far-UV flux (e.g.. Geersetal. 2006)). whichmay be indicative of a high EUV flux.," Although HD 34700A did not have a published line detection, its near- and mid-infrared spectra present unusually strong PAH emission bands (e.g., \citealt{smith04}) ), possibly hinting to a high stellar far-UV flux (e.g., \citealt{geers06}) ), whichmay be indicative of a high EUV flux."
+ 72 7. in dwarf galaxy. starbursts like NCC 1569 (2). in interacting eas-rich galaxies like NGC 1038/39 (the Auteunae) (???)..,"\cite*{Larsen04}, , \cite*{Mora+07}, in dwarf galaxy starbursts like NGC 1569 \citep{Anders+04b}, in interacting gas-rich galaxies like NGC 4038/39 (the Antennae) \citep{Whitmore+95,Whitmore+05,Anders+07}."
+ Slightly older auc iuteiuediate-age SCs ire observed in post-starburst merger remnants like NGC 7252 (2277) and ανασααν voung ellipticals like NGC 1316 οη. respectively.," Slightly older and intermediate-age SCs are observed in post-starburst merger remnants like NGC 7252 \citep{Whitmore+93,FB95,Miller+97,SchweizerSeitzer98} and dynamically young ellipticals like NGC 1316 \citep{Goudfrooij+01b,Goudfrooij+04,Goudfrooij+07}, respectively."
+ Thev can form all over the main body of a galaxy. as e.g. in the Auteunae or NGC. 7252. in aud around a starburs imeleus. as ee. in Arp 220 or NGC 6210 (??).. all along some but not all extended. tida features (?77).. as well as in group cuvirommeuts like Stephanus quintett (?)..," They can form all over the main body of a galaxy, as e.g. in the Antennae or NGC 7252, in and around a starburst nucleus, as e.g. in Arp 220 or NGC 6240 \citep{Shioya+01,Pasquali+03}, all along some -- but not all – extended tidal features \citep{Knierman+03,deGrijs+03a,Trancho+07a}, , as well as in group environments like Stephan's quintett \citep{Gallagher+01}."
+ These cuviromuents cover a huge range in terms of density. kinetic temperature. chemical abundances aud it is by no means obvious whether or not all these SCs are similar or different. mdividuallv or as a population.," These environments cover a huge range in terms of density, kinetic temperature, chemical abundances and it is by no means obvious whether or not all these SCs are similar or different, individually or as a population."
+ Related questious are where. when. aud how GC are ormmed and what a vouug GC looks like.," Related questions are where, when, and how GC are formed and what a young GC looks like."
+ Or how to tell apart YSCs iuto long-lived aud short-lived ones — bx nass. concentration. mass function. ...?," Or how to tell apart YSCs into long-lived and short-lived ones – by mass, concentration, mass function, ...?"
+ Current cluster formation models require exceptionally veh SE cficiencies SFE:=NLλεω>30X as a xerequisite for the formation of massive stronely pound and loue-term stable SCs. ic. for the formation of voung GCs (222?)..," Current cluster formation models require exceptionally high SF efficiencies ${\rm SFE~:=~M_{\ast}/M_{gas}~>30~\%}$ as a prerequisite for the formation of massive strongly bound and long-term stable SCs, i.e. for the formation of young GCs \citep{Brown+95,Burkert+96,ElmegreenEfremov97,Li+04}."
+ Ou a global scale. SF efficiencies ii normal spiral and regular galaxies. as well as iu starbursting dwarf galaxies are of order 0.11 ," On a global scale, SF efficiencies in normal spiral and irregular galaxies, as well as in starbursting dwarf galaxies are of order $0.1 - 3 \% $ \citep{Krueger+95}."
+On the smaller scale of molecular clouds in the Mile Wav. the SF efficiency is of the same order of magnitude and so is the mass ratio between the molecular cloud core and the entire molecular cloud.," On the smaller scale of molecular clouds in the Milky Way, the SF efficiency is of the same order of magnitude and so is the mass ratio between the molecular cloud core and the entire molecular cloud."
+ No GC formation 1s therefore expected iun spirals. irregulars or star-burstiug dwarf galaxies by today.," No GC formation is therefore expected in spirals, irregulars or star-bursting dwarf galaxies by today."
+ In giant eas-rich interacting ealaxies. on the other hand. SE effcieucies of order 1050 are reported on elobal scales. aud of order 30.9054 on nuclear scales ofa few hundred pe up to —1 kpc.," In giant gas-rich interacting galaxies, on the other hand, SF efficiencies of order $10 - 50 \%$ are reported on global scales, and of order $30 - 90 \%$ on nuclear scales of a few hundred pc up to $\sim 1$ kpc."
+ Iu those svstems. GC formation should be possible.," In those systems, GC formation should be possible."
+ The fact that different ματά lines (CO(1-0). IICN(I-0). CS(1-0)) trace molecular gas at different densities (nc100.0c35-10Ln~10 an). has allowed to see tha while in the Milkv Way aud other ον ealaxies onlv a small fraction (0.1 10 lass OD a lulecular cloud makes up its hieh deusitv core. the situation is drastically different iu Ultraluuinous Iufared galaxies (ULIRGs). which all are late stages of inassive eas-rich mergers with strong unclear (few 100 pc) starbursts.," The fact that different submm lines (CO(1-0), HCN(1-0), CS(1-0)) trace molecular gas at different densities ${\rm n \geq 100, n~\geq 3 \cdot 10^4, n~\sim 10^5~cm^{-3}}$ ), has allowed to see that while in the Milky Way and other nearby galaxies only a small fraction $0.1 - 3 \%$ ) of the mass of a molecular cloud makes up its high density core, the situation is drastically different in Ultraluminous Infrared galaxies (ULIRGs), which all are late stages of massive gas-rich mergers with strong nuclear (few 100 pc) starbursts."
+ Iu those ULIRGs. almost all the molecular gas iu the central starburst region is at the ligh densities of molecular cloudcores. indicating that the moleculay cloud structure must be verv different from what we know in our Calaxy.," In those ULIRGs, almost all the molecular gas in the central starburst region is at the high densities of molecular cloud, indicating that the molecular cloud structure must be very different from what we know in our Galaxy."
+ The eutire unclear region is just one superent molecular cloud core. seriously raising the question whether the star and SC formation processes cau be the same as m normal ealaxies. not to mention the situation in exteuded. expanding. low-density tidal structures in the outskirts of other interacting galaxies.," The entire nuclear region is just one supergiant molecular cloud core, seriously raising the question whether the star and SC formation processes can be the same as in normal galaxies, not to mention the situation in extended, expanding, low-density tidal structures in the outskirts of other interacting galaxies."
+ Iu anv case. before ALATA becomes operational. the YSCs forming in these different tvpes of environments are our best proxy to the molecular cloud. structure.," In any case, before ALMA becomes operational, the YSCs forming in these different types of environments are our best proxy to the molecular cloud structure."
+ Iu the Milkv. Wav. molecular cloud cores. molecular clouds. aud YSC's all feature power-law mass functions. sugeesting scale-free sclfsimilar evolution.," In the Milky Way, molecular cloud cores, molecular clouds, and YSCs all feature power-law mass functions, suggesting scale-free self-similar evolution."
+ Not even for the closest massive mereer. the Áutenuae. cam we prescutly determine the molecular cloud or cloud core lnass functions (cf. 7)).," Not even for the closest massive merger, the Antennae, can we presently determine the molecular cloud or cloud core mass functions (cf. \cite*{Wilson+03}) )."
+ The masses of YSC's and the shape of their mass fiction is all we can access (??)..," The masses of YSCs and the shape of their mass function is all we can access \citep{Wilson+06,Anders+07}."
+ 7? and 7 have shown that for all galaxies from Blue Compact Dwarfs to spirals aud ULIRGs there is a fiel correlation between SFR. as derived. from far-lnfrarec DIunuuositv. aud the mass in molecular cloud cores. as doerivec from the TICN luminosity.," \cite*{GaoSolomon04} and \cite*{Solomon+92} have shown that for all galaxies – from Blue Compact Dwarfs to spirals and ULIRGs – there is a tight correlation between SFR, as derived from far-infrared luminosity, and the mass in molecular cloud cores, as derived from the HCN luminosity."
+ They also find the SF efficiency. to be proportional to the mass ratio of molecular gas at core and nonual deusitics. Le. to the ratio between Πο or CS huuimositv aud CO luminosity.," They also find the SF efficiency to be proportional to the mass ratio of molecular gas at core and normal densities, i.e. to the ratio between HCN or CS luminosity and CO luminosity."
+ The highest density molecular eas iu all these euvirouments is transformed iuto stars with almost 1060 efficiency. aud it is the amount of gas at hose high densities hat coutrols SE., The highest density molecular gas in all these environments is transformed into stars with almost 100 efficiency and it is the amount of gas at those high densities that controls SF.
+ The fraction of uolecular gas at the highest deusities therefore defines he SF efficiency., The fraction of molecular gas at the highest densities therefore defines the SF efficiency.
+ The high ambicut pressure building up in the course of massive eas-vich mergers cau drive up SE efficicucies wil2 orders of magnitude by compressing molecular clouds. Increasing their masses and. duo particular. heir core mass fractious.," The high ambient pressure building up in the course of massive gas-rich mergers can drive up SF efficiencies by $1-2$ orders of magnitude by compressing molecular clouds, increasing their masses and, in particular, their core mass fractions."
+ ?? have shown that the ISAL pressure during mergers can easily become 3.1 ines licher than the typical internal molecular cloud xessure. raising the SE cfiiciency to 7090%.," \cite*{JogDas92,JogDas96} have shown that the ISM pressure during mergers can easily become $3-4$ times higher than the typical internal molecular cloud pressure, raising the SF efficiency to $70-90 \%$."
+ This cads us to expect that the relative amount of SF that eoes into the formation of massive. strongly bound voung GCs in relation to the amount of SF that goes into field stars and low-mass. short-lived clusters is chhauced in massive gas-rich 1uergers.," This leads us to expect that the relative amount of SF that goes into the formation of massive, strongly bound young GCs in relation to the amount of SF that goes into field stars and low-mass, short-lived clusters is enhanced in massive gas-rich mergers."
+ Before I turn to the results obtained so far for SCs and SC populations in different euvironments. I briefly recall our GALEV evolutionary svuthesis iuodels aud he dedicated analysis tools we use iu our analysis of SC systems.," Before I turn to the results obtained so far for SCs and SC populations in different environments, I briefly recall our GALEV evolutionary synthesis models and the dedicated analysis tools we use in our analysis of SC systems."
+GALEV models in the first place deseribethe spectral evolution of SCs of various netallicities 1.7€|Fe/TI]<10.1 over the age range,GALEV models in the first place describethe spectral evolution of SCs of various metallicities ${\rm -1.7 \leq [Fe/H] \leq +0.4}$ over the age range
+density inhomogencities.,density inhomogeneities.
+ Figure 4 also shows the temporal evolution of ej and σι assuming that the initial perturbed surface density is lower than in Sect., Figure 4 also shows the temporal evolution of $\sigma _0$ and $\tilde \sigma _1$ assuming that the initial perturbed surface density is lower than in Sect.
+ 2.3.2. namely σον) instead. of συέρως.," 2.3.2, namely $\sigma _0(t_{em})$ instead of $\sigma _0(t_{em})$."
+ Despite these different initial values. cach panel displays almost identical σι(/) curves in both cases.," Despite these different initial values, each panel displays almost identical $\tilde \sigma _1(t)$ curves in both cases."
+" Owing to the initial transverse Lows. the initial clumps quickly unclergo a replenishment in shell material leacing afterwards to very similar temporal evolutions of the. perturbed surface density whatever the initial perturbed surface density,"," Owing to the initial transverse flows, the initial clumps quickly undergo a replenishment in shell material leading afterwards to very similar temporal evolutions of the perturbed surface density whatever the initial perturbed surface density."
+" Obviously. this one is not a key parameter of the fragmentation process,"," Obviously, this one is not a key parameter of the fragmentation process."
+ All the results presented here above assume the spherical svnunetry of the supershell., All the results presented here above assume the spherical symmetry of the supershell.
+ Lt is clear however that such a system is not expected to remain perfectly spherical., It is clear however that such a system is not expected to remain perfectly spherical.
+ The deviations from spherical svnimetry can arise. for instance. from the non point-like nature of the energy input. namely all the stars of the first generation cluster are not located exactly at the centre of the cloud/shell.," The deviations from spherical symmetry can arise, for instance, from the non point-like nature of the energy input, namely all the stars of the first generation cluster are not located exactly at the centre of the cloud/shell."
+ A rough estimate of the initial amplitude of the transverse motions within the shell can therefore be derived from the size of this cluster., A rough estimate of the initial amplitude of the transverse motions within the shell can therefore be derived from the size of this cluster.
+ A star located at a distance r from the shell centre will induce a transverse velocity. e such that: In order to estimate the size rr of the cluster of massive stars hosted by a POGCC. we refer to 111136. the dense core of the DDoradus Nebula located in the Large Magellanic Cloud.," A star located at a distance $r$ from the shell centre will induce a transverse velocity $v$ such that: In order to estimate the size $r$ of the cluster of massive stars hosted by a PGCC, we refer to 136, the dense core of the Doradus Nebula located in the Large Magellanic Cloud."
+ “Phe DDoraclus nebula shows an impressive example of a two-stage stellar formation., The Doradus nebula shows an impressive example of a two-stage stellar formation.
+ The energetic activitv of a very compact bright. cluster. 11136. which includes several tens of O stars. triggers the formation of a new stellar. generation revealed. by numerous infrared sources in or near some bright filaments west and northeast of h136 (Rubio et al.," The energetic activity of a very compact bright cluster, R136, which includes several tens of O stars, triggers the formation of a new stellar generation revealed by numerous infrared sources in or near some bright filaments west and northeast of R136 (Rubio et al."
+ 1998)., 1998).
+ 3ased on Hubble: Space Telescope photometry. Campbell et al. (," Based on Hubble Space Telescope photometry, Campbell et al. ("
+1992). detected about 160 stars more massive than MM. in. 101136 which they define as a region of ppc ppc.,1992) detected about 160 stars more massive than $_{\odot}$ in 136 which they define as a region of pc $\times$ pc.
+ This number of massive stars being remarkably similar to the numbers of SNell used in our self-enrichment. model. we adopt 101136 as the most similar example in the Local Croup of what may have been the first. generation cluster.," This number of massive stars being remarkably similar to the numbers of SNeII used in our self-enrichment model, we adopt 136 as the most similar example in the Local Group of what may have been the first generation cluster."
+ A radius of ppc appears therefore as a reasonable estimate of the size of this cluster. the source of the energy. input.," A radius of pc appears therefore as a reasonable estimate of the size of this cluster, the source of the energy input."
+ Considering an average background pressure οἱ 10Lo ddvnecm and the corresponding: cloud: radius.: (2 ppc. see Eq.," Considering an average background pressure of $^{-10}$ $^{-2}$ and the corresponding cloud radius $\simeq$ pc, see Eq."
+" 3. 7Paper 1). the transverse velocity when the shell reaches the cloud boundary. is: Figure 5 displavs the evolution. with time ofσυ and 6, assuming three cdillerent values for the initial transverse velocity. namely VG). Velho) and Voto). while keeping all the other parameters to their previous values."," 3, Paper I), the transverse velocity when the shell reaches the cloud boundary is: Figure 5 displays the evolution with time of$\sigma _0$ and $\tilde \sigma _1$ assuming three different values for the initial transverse velocity, namely $V_s(t_{em})$, $V_s(t_{em})$ and $V_s(t_{em})$, while keeping all the other parameters to their previous values."
+ In. sharp contrast with σι(fan). CU) appears to be an important parameter of the shell transverse collapse.," In sharp contrast with $\tilde \sigma _1(t_{em})$, $\tilde v(t_{em})$ appears to be an important parameter of the shell transverse collapse."
+ For instance. considering the bottom panel in Fig.," For instance, considering the bottom panel in Fig."
+" 5 (3,25.10I 7 andN 2200). the [ragmoentation takes place less than 15 million vears after the first SN explosion if 0(/,,,)=0.03 VC). whereas it is prevented if (Eau)=OOLVCo, )."," 5 $P_h=5 \times 10^{-10}$ $^{-2}$ and$N$ =200), the fragmentation takes place less than 15 million years after the first SN explosion if $v(t_{em})=0.03 \,V_s(t_{em})$ , whereas it is prevented if $v(t_{em})=0.01 \,V_s(t_{em})$ ."
+caleulatious of uncertainties in the eectron collision data is noeelieije.,calculations of uncertainties in the electron collision data is negligible.
+ Iudeed. together with tie collision data for the charge exchange process Li(3s)II—LiLII now available. and barring the existeicc of an uuknown important collisional process. the οςJlisional data in general is not a source of significant 1neertainty iu non-LTE Li line formation calculations.," Indeed, together with the collision data for the charge exchange process $\mathrm{Li(3s)} + \mathrm{H} \rightleftharpoons \mathrm{Li}^+ + \mathrm{H}^-$ now available, and barring the existence of an unknown important collisional process, the collisional data in general is not a source of significant uncertainty in non-LTE Li line formation calculations."
+We use (he code (Melianiοἱal.2007:IXeppenset2011).. which solves sets ol near-conservation laws on an adaptive mesh.,"We use the code \citep{metal07,keppens11}, which solves sets of near-conservation laws on an adaptive mesh."
+ For (his particular study. we added a new module to the code which calculates the behavior of dust grains as coupled to (he lamiliar eas dvnamic equations.," For this particular study, we added a new module to the code which calculates the behavior of dust grains as coupled to the familiar gas dynamic equations."
+ The dust and gas are linked (rough drag forces which {μον exert on each other., The dust and gas are linked through drag forces which they exert on each other.
+ Besides conservation of mass. the moment and energv equations for the gas become with p the gas density. v the gas velocity. p the thermal pressure and € the total energy densitv for the gas.," Besides conservation of mass, the momentum and energy equations for the gas become with $\rho$ the gas density, $\vecv$ the gas velocity, $p$ the thermal pressure and $e$ the total energy density for the gas."
+ We assume a standard ideal gas law lor closure., We assume a standard ideal gas law for closure.
+ The right hand side source ternis for the energy evolution describe optically (hin radiative cooling aud (he work done bv the drag forces fy., The right hand side source terms for the energy evolution describe optically thin radiative cooling and the work done by the drag forces $\fdrag$.
+ Racliative losses depend on the hydrogen ancl electron particle densities (derived from p assuming full ionization with hydrogen mass jg) and involve a temperature dependent cooling curve A(T)., Radiative losses depend on the hydrogen and electron particle densities (derived from $\rho$ assuming full ionization with hydrogen mass $m_{\rm h}$ ) and involve a temperature dependent cooling curve $\Lambda(T)$.
+ As most of the gas in our simulation is at Comparatively low temperature. we use a cooling curve A(Z) that extends down to 1 Ik. with tabulated info based on numerical caleulations done with the code (Ferlandetal.1993).," As most of the gas in our simulation is at comparatively low temperature, we use a cooling curve $\Lambda(T)$ that extends down to 1 K, with tabulated info based on numerical calculations done with the code \citep{fetal98}."
+. The drag force for species d per unit volume fy is specilied in Eq. 5:, The drag force for species ${\rm d}$ per unit volume $\fdrag$ is specified in Eq. \ref{eq:drag};
+ since we (treat niultiple dust species we sum over the individual drag forces., since we treat multiple dust species we sum over the individual drag forces.
+ We follow the prescription [rom Paarcekooper&Mellema(2006)... treating the dust as a gas without pressure. since (he internal energy. of a cdustgrain only inlluences its surface temperature. but has no influence on ils movement.," We follow the prescription from \citet{pm06}, treating the dust as a gas without pressure, since the internal energy of a dustgrain only influences its surface temperature, but has no influence on its movement."
+ Therefore. (he motion of each dustspecies can be treated with the equations of conservation of mass and momentum: where pq and v4 are the mass density and velocity of the dust species.," Therefore, the motion of each dustspecies can be treated with the equations of conservation of mass and momentum: where $\rhod$ and $\vd$ are the mass density and velocity of the dust species."
+ The drag force fy is given by a combination of Epstein’s drag law for the subsonic regime anc Stokes’ crag law lor the supersonic regime (οκ 1975).. which in closed form writes as," The drag force $\fdrag$ is given by a combination of Epstein's drag law for the subsonic regime and Stokes' drag law for the supersonic regime \citep{k75}, , which in closed form writes as"
+only by4.45... which is negligible compared with uncertainty of the Crab-based PSF due to Crab statisties.,"only by, which is negligible compared with uncertainty of the Crab-based PSF due to Crab statistics."
+ Finally. we note that the width of the PSF decreases with energv.," Finally, we note that the width of the PSF decreases with energy."
+ We verified this by comparing Crab images in 35 GeV and 510 GeV bands. and confirming that the former is broader than the latter.," We verified this by comparing Crab images in 3–5 GeV and 5–10 GeV bands, and confirming that the former is broader than the latter."
+ Since the spectrum of the nearby AGN is harder than that of Crab. the instrumental svstematics can only make the AGN image sharper. nol broader. but the opposite is interred [rom Fig. 6..," Since the spectrum of the nearby AGN is harder than that of Crab, the instrumental systematics can only make the AGN image sharper, not broader, but the opposite is inferred from Fig. \ref{fig:crab}."
+ Therefore. it is conservative to ignore (he small svstematic uncertainties due lo spectrum dependence of PSF.," Therefore, it is conservative to ignore the small systematic uncertainties due to spectrum dependence of PSF."
+ Acloptinge the Crab profile as a calibrated PSF. we quantitatively investigatee the excess of nearby/hard AGN profiles identified al 6?>0.2 deg? in Fig. 6..," Adopting the Crab profile as a calibrated PSF, we quantitatively investigate the excess of nearby/hard AGN profiles identified at $\theta^2 \gtrsim 0.2$ $^2$ in Fig. \ref{fig:crab}."
+ The PSF and AGN proliles are normalized to each other such (hat. they give the same brightness in the innermost angular bin.. 6?9<.0.045 > deg?.," The PSF and AGN profiles are normalized to each other such that they give the same brightness in the innermost angular bin, $\theta^2 <
+0.045$ $^2$ ."
+ |The excess photon counts are Nc(205)y(93n?>125d30(stat)+ 2i(svs).," The excess photon counts are $N_{\rm excess}^{(z<0.5)}(\theta^2 > 0.225
+~\mathrm{deg^2}) = 125 \pm 30({\rm stat}) \pm 21({\rm sys})$ ."
+ By takine square root of quadratic sum of the statistical and svstematic errors as a total error. we find that this excess is of 3.50 significance.," By taking square root of quadratic sum of the statistical and systematic errors as a total error, we find that this excess is of $3.5\sigma$ significance."
+ For the distant/solt AGN population. on the other hand. the excess is NGO>0.2255ELT(slal) 2ZO(svs). consistent with null hypothesis.," For the distant/soft AGN population, on the other hand, the excess is $N_{\rm excess}^{(z>0.5)}(\theta^2 > 0.225 ~\mathrm{deg^2}) = -5 \pm
+27({\rm stat}) \pm 29({\rm sys})$ , consistent with null hypothesis."
+" Dividing N44 bv the total number ol PSF counts. we obtain the values of fia, for both AGN populations: Clearly. this conclusion usinge the Crab-calibrated PSF agreese with that based on the pre-launch calibration."," Dividing $N_{\rm excess}$ by the total number of PSF counts, we obtain the values of $f_{\rm halo}$ for both AGN populations: Clearly, this conclusion using the Crab-calibrated PSF agrees with that based on the pre-launch calibration."
+ One can go even further and design (wo separate Crab-calibrated PSFs for two classes of photons. namely those that convert in the front. laver and (hose in the back laver of the detector.," One can go even further and design two separate Crab-calibrated PSFs for two classes of photons, namely those that convert in the front layer and those in the back layer of the detector."
+ While all of these photons must be used in an analysis. allowing for the differences in PSF offers vet another opportunity to find and eliminate some unexpected instrumental effects.," While all of these photons must be used in an analysis, allowing for the differences in PSF offers yet another opportunity to find and eliminate some unexpected instrumental effects."
+" To this end. we introduce another statistical quantity O44,NEI NRE. where all the Ns with sell-explanatory. superscripts and subscripts refer to photon counts αἱ 02?>0.225 deg? after the homogeneous backgrounds were subtracted."," To this end, we introduce another statistical quantity $\delta_{\rm
+excess} \equiv N_{\rm excess}^{\rm front} / N_{\rm psf}^{\rm front} + 
+N_{\rm excess}^{\rm back} / N_{\rm psf}^{\rm back}$ , where all the $N$ 's with self-explanatory superscripts and subscripts refer to photon counts at $\theta^2 > 0.225$ $^2$ after the homogeneous backgrounds were subtracted."
+ This way. we explicitly include any PSF differences between Iront. and back-converted. photons.," This way, we explicitly include any PSF differences between front and back-converted photons."
+" The meaninge of desees,CCS Is clear: a value consistent with zero corresponds to absence of physical halos.", The meaning of $\delta_{\rm excess}$ is clear: a value consistent with zero corresponds to absence of physical halos.
+ We obtain Oz=1.4zEO.5(stat)£0.2(svs). 2.070 away from the null hypothesis.," We obtain $\delta_{\rm excess}^{(z<0.5)} = 1.4 \pm 0.5 ({\rm stat}) \pm
+0.2 ({\rm sys})$, $2.7\sigma$ away from the null hypothesis."
+ We also fine that the individual values of Oexcess for the front and back photons are consistent with each other. within errors.," We also find that the individual values of $\delta_{\rm excess}$ for the front and back photons are consistent with each other, within errors."
+ We have also performed the same analvsis for LO100 GeV. IHere. we renormalized the Crab and AGN profiles using 67<0.025 deg?bin. and counted the excess photons over Crab- PSF in 9?=0.075 0.25 deg?.," We have also performed the same analysis for 10–100 GeV. Here, we renormalized the Crab and AGN profiles using $\theta^2 < 0.025$ $^2$bin, and counted the excess photons over Crab-calibrated PSF in $\theta^2 = 0.075$ –0.25 $^2$."
+ We obtain NU0.5)=19+13(stat)&I5(sys) and," We obtain $N_{\rm excess}^{(z<0.5)} = 19 \pm 13 ({\rm stat}) \pm 15
+({\rm sys})$ and"
+SED.,SED.
+ The four optical-UV data points appear to be roughly on (he extrapolation of the soft X-ray SED. suggesting (he same origin of the optical-UV. and soft X-ray emission.," The four optical-UV data points appear to be roughly on the extrapolation of the soft X-ray SED, suggesting the same origin of the optical-UV and soft X-ray emission."
+ We perform a detailed spectral and temporal analvsis lor the second oobservation of S5 0T164-714., We perform a detailed spectral and temporal analysis for the second observation of S5 0716+714.
+ Most of our results are in agreement will previous results obtained from the first oobservation (FEOG) and other X-ray observations (e.g.. Giommi et al.," Most of our results are in agreement with previous results obtained from the first observation (FE06) and other X-ray observations (e.g., Giommi et al."
+ 1999: Taglialerri et al., 1999; Tagliaferri et al.
+ 2003)., 2003).
+ Nevertheless. we also [ind some new phenomena. adding new clues to better understand the underlying physical processes taking place in (he source.," Nevertheless, we also find some new phenomena, adding new clues to better understand the underlying physical processes taking place in the source."
+ The concave X-ray spectra of $5 07164-1714 can be disentangled into two power law components., The concave X-ray spectra of S5 0716+714 can be disentangled into two power law components.
+ The steep power law (I~ 2.6) component is interpreted as the hieh energy tail of the svuehrotron emission. whereas the flat power law (DL 1.2) component is ascribed {ο the low energy side of the IC emission.," The steep power law $\Gamma \sim 2.6$ ) component is interpreted as the high energy tail of the synchrotron emission, whereas the flat power law $\Gamma \sim1.2$ ) component is ascribed to the low energy side of the IC emission."
+ EEO06 obtained similar results with the first oobservation., FE06 obtained similar results with the first observation.
+ Ht is worth notüng that the photon indices of the steep power law component are similar to those of IIBLs in the hard. X-ray band. (e.g.. PINS 2155304: Zhang 2003). supporting the interpretation of the steep power law component as the svichrotron tail.," It is worth noting that the photon indices of the steep power law component are similar to those of HBLs in the hard X-ray band (e.g., PKS 2155–304: Zhang 2008), supporting the interpretation of the steep power law component as the synchrotron tail."
+ The X-rav variability amplitude of WBLs monotonically increases with hieher energy (e.g.. PINS 2155304: Zhang οἱ al.," The X-ray variability amplitude of HBLs monotonically increases with higher energy (e.g., PKS 2155–304: Zhang et al."
+ 2005; Mrk 421: Zhang et al., 2005; Mrk 421: Zhang et al.
+ 2010). which is thought to be the signature of svnchrotron emission.," 2010), which is thought to be the signature of synchrotron emission."
+ However. LBLs are highly variable in the soft X-ravs. whereas (hey show little variability in the hard. X-rays (e.g.. Ciommi et al.," However, LBLs are highly variable in the soft X-rays, whereas they show little variability in the hard X-rays (e.g., Giommi et al."
+ 1999)., 1999).
+ Our results demonstrate that S5 O716+714 is indeed strongly variable in the soft. N-ravs. showing the maximum variability by a [actor of ~4 throughout the whole observation and several episodes of rapid variations on timescales of hours.," Our results demonstrate that S5 0716+714 is indeed strongly variable in the soft X-rays, showing the maximum variability by a factor of $\sim4$ throughout the whole observation and several episodes of rapid variations on timescales of hours."
+ In a sharp-cut contrast. the had X-ray [Iuxes of the source are niuch less variable. exhibiting only ~50% change between the minimun and maximum count rates and no rapid events.," In a sharp-cut contrast, the hard X-ray fluxes of the source are much less variable, exhibiting only $\sim50\%$ change between the minimum and maximum count rates and no rapid events."
+ For the first time. we quantify the enerev dependence of the variability amplituce for 55 07162714.," For the first time, we quantify the energy dependence of the variability amplitude for S5 0716+714."
+ The variability aaplitude increases [rom (he 0.30.5 t0 0.57.5 keV band. but [rom the 0.751 keV bane. it decreases with hieher energv.," The variability amplitude increases from the 0.3–0.5 to 0.5–7.5 keV band, but from the 0.75–1 keV band, it decreases with higher energy."
+ Moreover. the OM data suggest lower variability amplitude in (he UV band than in the soft X-ray baud.," Moreover, the OM data suggest lower variability amplitude in the UV band than in the soft X-ray band."
+ The energy. dependence of the variability amplitude of $5 OT1G+714 is thus clearly different. from those of IIBEs., The energy dependence of the variability amplitude of S5 0716+714 is thus clearly different from those of HBLs.
+shown in the bottom panels of Fig. 1..,shown in the bottom panels of Fig. \ref{spmapclv}.
+ The positive Stokes V signals at the red wing are of opposite to the polarity compared to the spot., The positive Stokes V signals at the red wing are of opposite to the polarity compared to the spot.
+ Examples of Stokes profiles with the positive enhancements at the red wing are shown in Figs., Examples of Stokes profiles with the positive enhancements at the red wing are shown in Figs.
+ 2 (b) and 3 (b)., \ref{sprof1} (b) and \ref{sprof2} (b).
+ The Stokes V profiles in these regions are characterized as three lobe profiles which are completely different from a regular antisymmetric Stokes V profile., The Stokes V profiles in these regions are characterized as three lobe profiles which are completely different from a regular antisymmetric Stokes V profile.
+" The red lobe of the Stokes profiles is produced by a strongly red-shifted Stokes V profile with the polarity opposite to the major polarity of the spot (e.g.SánchezAlmeida&Ichimoto,2009).", The red lobe of the Stokes profiles is produced by a strongly red-shifted Stokes V profile with the polarity opposite to the major polarity of the spot \cite[e.g.][]{sanchesalmeida2009}.
+. The patches with the positive Stokes V signals are mainly observed near the boundary of the penumbra when the spot was located near the disk center., The patches with the positive Stokes V signals are mainly observed near the boundary of the penumbra when the spot was located near the disk center.
+" They are attributed to downflows of the Evershed flows along magnetic field lines returning into the photosphere atthe penumbral boundary (WestendorpPlazaetal.,2001; 2008b)."," They are attributed to downflows of the Evershed flows along magnetic field lines returning into the photosphere at the penumbral boundary \citep{westendorp2001,bellotrubio2004,ichimoto2007,
+shimizu2008b}."
+. It is realized that similar positive patches are also seen even in the middle of the penumbra in the bottom panels of Fig. 1.., It is realized that similar positive patches are also seen even in the middle of the penumbra in the bottom panels of Fig. \ref{spmapclv}.
+ These were studied by SainzDalda&BellotRubio(2008) who showed that these can be associated with the sea-serpent field lines in the mid-penumbra., These were studied by \cite{sainzdalda2008} who showed that these can be associated with the sea-serpent field lines in the mid-penumbra.
+" The third ones, which have not been realized before, are enhancements of Stokes V signals at the red wing similar with the previous one, but with the magnetic polarity same as the spot."," The third ones, which have not been realized before, are enhancements of Stokes V signals at the red wing similar with the previous one, but with the magnetic polarity same as the spot."
+ The patches having negative enhancements at the red wing are indicated by solid arrows in the bottom panels of Fig. 1.., The patches having negative enhancements at the red wing are indicated by solid arrows in the bottom panels of Fig. \ref{spmapclv}.
+" The patches are less frequent than the patches with the positive enhancements, and their sizes are smaller than 0.5 arcseconds."," The patches are less frequent than the patches with the positive enhancements, and their sizes are smaller than 0.5 arcseconds."
+ This is probably why the features have not been realized before., This is probably why the features have not been realized before.
+" Figs. 2,,"," Figs. \ref{sprof1},"
+ 3 (c) and (d) show examples of Stokes profiles observed inside the patches marked with the solid arrows., \ref{sprof2} (c) and (d) show examples of Stokes profiles observed inside the patches marked with the solid arrows.
+ The Stokes V profiles have humps at the red wing with the same polarity., The Stokes V profiles have humps at the red wing with the same polarity.
+ The Stokes V profile in Figs., The Stokes V profile in Figs.
+" 3 (d) does not have a hump at the red wing, but it has a red tail causing the enhancements in the red wing images of Stokes V, and the Stokes V amplitude of the red lobe is smaller than that of the blue lobe."," \ref{sprof2} (d) does not have a hump at the red wing, but it has a red tail causing the enhancements in the red wing images of Stokes V, and the Stokes V amplitude of the red lobe is smaller than that of the blue lobe."
+ The enhancements at the red wing are not so large in the Stokes Q and U profiles though there are weak tails at the red wing in the Stokes Q profiles in Fig., The enhancements at the red wing are not so large in the Stokes Q and U profiles though there are weak tails at the red wing in the Stokes Q profiles in Fig.
+ 2 (c) and (d)., \ref{sprof1} (c) and (d).
+ The Stokes profiles are easily distinguishable from the known profiles which are supposed to be associated with the Evershed flow., The Stokes profiles are easily distinguishable from the known profiles which are supposed to be associated with the Evershed flow.
+ The enhanced Stokes V signals at the red wing can be caused by either temperature changes or magnetic, The enhanced Stokes V signals at the red wing can be caused by either temperature changes or magnetic
+ , 
+broken scale invariant or BCDAL model.,broken scale invariant or BCDM model.
+ Earlier ciscussions of this model and an analytic fit to the spectrum are given in 7).., Earlier discussions of this model and an analytic fit to the spectrum are given in \citet{KMGMR95}.
+ We perform simulations with 3005 particles in 600° cells. and we simulate large boxes of (500IMpe)? volume.," We perform simulations with $300^3$ particles in $600^3$ cells, and we simulate large boxes of $(500
+\mhmpc)^3$ volume."
+ μονο simulations are described in more detail in 2).. where we studied the cluster power spectrum. and in ?) αμα 7). where we simulated: overdensity regions in the LORS that. correspond to. superclusters of galaxies.," These simulations are described in more detail in \citet{RBGKM98}, where we studied the cluster power spectrum, and in \citet{MDRT98} and \citet{DMRT99}, where we simulated overdensity regions in the LCRS that correspond to superclusters of galaxies."
+ l'herefore. the present investigation should. provide. complementary information.," Therefore, the present investigation should provide complementary information."
+ In using a large box size. we have a sullicient volume to simulate reliably voids. with sizes of up to 60 and to find a reasonable representation of the void hierarchy.," In using a large box size, we have a sufficient volume to simulate reliably voids with sizes of up to 60 and to find a reasonable representation of the void hierarchy."
+ “Phe price to be paid. for thesebox sizes is a particle mass of (13)104TAL..., The price to be paid for thesebox sizes is a particle mass of $(1 - 3) \times 10^{11}$.
+ In other words. we must identify galaxies with single mass points.," In other words, we must identify galaxies with single mass points."
+ For galaxy identification. we cmplov the ideas of 7) to dillerentiate the simulation particles in voids Lor low environmental densities and in clustered galaxies [or densities higher than a critical density δν.," For galaxy identification, we employ the ideas of \citet{EJS80} to differentiate the simulation particles in voids for low environmental densities and in clustered galaxies for densities higher than a critical density $\delta_{th}$."
+ To this aim. we determine the density around cach simulation particle at a fixed radius of tAlpe.," To this aim, we determine the density around each simulation particle at a fixed radius of 1 ."
+. Then. we identify no galaxies if," Then, we identify no galaxies if"
+which is sometimes referred to as the reduced shear.,which is sometimes referred to as the reduced shear.
+ Furthermore. (he average ellipticities of background galaxies is equal to the complex distortion. 0=29g/(1+[g[?).," Furthermore, the average ellipticities of background galaxies is equal to the complex distortion, $\delta
+\equiv 2 g /(1+|g|^2)$."
+ Solving this quadratic equation for g results in g=(14V1jo?)/o*., Solving this quadratic equation for $g$ results in $g=(1 \pm \sqrt{1-|\delta|^2})/\delta^{*}$.
+ Since there are two roots to this equation. one cannot tell fom purely local measurements which root to choose.," Since there are two roots to this equation, one cannot tell from purely local measurements which root to choose."
+ This reflects an inherent local degeneracy (ο anv weak lensing measurements., This reflects an inherent local degeneracy to any weak lensing measurements.
+ For weak fields. one should lake (he negative root. for which 5~gy90/2.," For weak fields, one should take the negative root, for which $\gamma \simeq g
+\simeq \delta/2$."
+ The choice of sien is related to image parity: ib changes when crossing critical curves in (he lens plane. where the magnification diverges and ares are present.," The choice of sign is related to image parity; it changes when crossing critical curves in the lens plane, where the magnification diverges and arcs are present."
+ In practice. it is onlv in the cores of massive clusters where one mich make the wrong choice of parity. and (his is unlikely (ο present a major problem for the statistical method of weak lensing we pursue here.," In practice, it is only in the cores of massive clusters where one might make the wrong choice of parity, and this is unlikely to present a major problem for the statistical method of weak lensing we pursue here."
+ These rare. extreme regions are not the places where weak lensing will be the most useful probe of the mass distribution.," These rare, extreme regions are not the places where weak lensing will be the most useful probe of the mass distribution."
+ The other continuous degeneracy is not as problematic: one can show that if gq is measured. one can invert for & up to the mass sheet degeneracy (kaiser1995)..," The other continuous degeneracy is not as problematic: one can show that if $g$ is measured, one can invert for $\kappa$ up to the mass sheet degeneracy \citep{kaiser:nonlinear}."
+ For our azimuthally svannmnetric case. we can simply substitute g(1—&) for ? in Equation 4.. resulting in so that In(1—&) is determined up to a constant.," For our azimuthally symmetric case, we can simply substitute $g (1-\kappa)$ for $\gamma$ in Equation \ref{eq:delta-kappa}, resulting in so that $\ln (1-\kappa)$ is determined up to a constant."
+ Buposing the boundary. condition K(oc)=0. one can integrate this ordinary differential equation: which in the weak-field limit becomes In section 5.. we will discuss the truncation of integrals of this kind to a finite region.," Imposing the boundary condition $\kappa(\infty)=0$, one can integrate this ordinary differential equation: which in the weak-field limit becomes In section \ref{section:practical}, we will discuss the truncation of integrals of this kind to a finite region."
+ since Equation 9 is linear in sy. one could simply multiply it by X4 to obtain an expression for XR).," Since Equation \ref{eq:kappa-weak} is linear in $\gamma_T$, one could simply multiply it by $\Sigma_{crit}$ to obtain an expression for $\Sigma(R)$."
+ This is suitable for an analvsis such as that in Sheldonetal.(2004)... where AN=Mont is averaged.," This is suitable for an analysis such as that in \cite{sheldon:gmcf}, where $\Delta\Sigma \equiv \Sigma_{crit} \gamma$ is averaged."
+ DBevond the weak field regime. the situation is more complicated. because of the 1—g term in the denominator of G(H): in this case. an average of the ensing distortion does not simply vield an average of X(R).," Beyond the weak field regime, the situation is more complicated, because of the $1-g$ term in the denominator of $G(R)$; in this case, an average of the lensing distortion does not simply yield an average of $\Sigma(R)$."
+ Nonetheless. as noted. above. (his should only be a problem for small scales and high. densities. i.e.. for the cores of very nassive clusters. and in (hat regime one could apply a correction or an iterative process to determine X(A).," Nonetheless, as noted above, this should only be a problem for small scales and high densities, i.e., for the cores of very massive clusters, and in that regime one could apply a correction or an iterative process to determine $\Sigma(R)$."
+ Moreover. there is only a small spatial range where (his correction is important and where one can still safely assume (hal 5=gy790/2 is the correct root.," Moreover, there is only a small spatial range where this correction is important and where one can still safely assume that $\gamma = g \simeq \delta/2$ is the correct root."
+ The uethod of averaging shear profiles mav simply not be the most prudent method lor probing, The method of averaging shear profiles may simply not be the most prudent method for probing
+the star formation mode is expected to occurs at 1073-5Ze (Mackeyetal.2003).,the star formation mode is expected to occurs at $10^{-3.5}\ Z_{\odot}$ \citep{Mack03}.
+". In this case, the CMB limits the lower masses of stars to a few 10s of Me (Smithetal.2008;Schneider&Omukai 2010)."," In this case, the CMB limits the lower masses of stars to a few $10$ s of $M_{\odot}$ \citep{Smi08, Sch10}."
+". Formation history of galaxies considering the time-dependent IMF from the top-heavy to the Salpeter-like IMF and taking into account the production and destruction of dust by PISNe, will be explored in the forthcoming paper (Yamasawa et al."," Formation history of galaxies considering the time-dependent IMF from the top-heavy to the Salpeter-like IMF and taking into account the production and destruction of dust by PISNe, will be explored in the forthcoming paper (Yamasawa et al."
+ 2010 in preparation)., 2010 in preparation).
+" Throughout this paper we adopt the models by Nozawaetal.(2003,2007) for dust formation and destruction."," Throughout this paper we adopt the models by \citet{Noz03, Noz07} for dust formation and destruction."
+" Nozawaetal.(2003) investigated the dust production in the ejecta of primordial SNe II as well as PISNe, applying a theory of non-steady state nucleation and grain growth."," \citet{Noz03} investigated the dust production in the ejecta of primordial SNe II as well as PISNe, applying a theory of non-steady state nucleation and grain growth."
+ They revealed the grain species formed in the ejecta and their size distributions for the unmixed and mixed elemental compositions within the He core., They revealed the grain species formed in the ejecta and their size distributions for the unmixed and mixed elemental compositions within the He core.
+" In what follows, we apply the results of calculation for the unmixed ejecta of SNe IIwith the progenitor mass m=13,20,25,and30Mo and the explosion energy 10°!erg, and extrapolate the data to the mass range from 8 to 40Mo."," In what follows, we apply the results of calculation for the unmixed ejecta of SNe IIwith the progenitor mass $m=13, 20, 25,\ {\rm and}\ 30\ M_{\odot}$ and the explosion energy $10^{51}\ {\rm erg}$ and extrapolate the data to the mass range from $8$ to $40\ {M_{\odot}}$."
+" The solid line in Figure 1 shows the IMF-averaged mass distribution of dust formed in the ejecta of SNe II, M9(a), whichis weighted by the Salpeter IMF and is summed up over all the grain species as,where M5;(a,m)da is the mass of the j-th dust species produced in a SN II with radii between a and a+da as a function of progenitor mass m, and superscript 0 means the case with no destruction by a reverse shock."," The solid line in Figure \ref{fig:prod} shows the IMF-averaged mass distribution of dust formed in the ejecta of SNe II, $\overline{{\cal M}_{{\rm d}}^{0}}(a)$, whichis weighted by the Salpeter IMF and is summed up over all the grain species as,where ${\cal M}_{{\rm d},j}^{0}(a,m){\rm d}a$ is the mass of the$j$ -th dust species produced in a SN II with radii between $a$ and $a+{\rm d}a$ as a function of progenitor mass $m$, and superscript 0 means the case with no destruction by a reverse shock."
+" In Figure 1,, we plot aM9(a) in the vertical axis to make clear the mass fraction in each logarithmic bin."," In Figure \ref{fig:prod}, we plot $a\overline{{\cal M}_{{\rm d}}^{0}}(a)$ in the vertical axis to make clear the mass fraction in each logarithmic bin."
+ We can see that the grain radii range from a few uup to a few wm and that the size spectrum of dust in mass has a peak at a~0.1jum., We can see that the grain radii range from a few up to a few $\mu{\rm m}$ and that the size spectrum of dust in mass has a peak at $a\sim0.1\ \mu{\rm m}$.
+" In the course of their injection into ISM, dust grains formed in the ejecta are destroyed due to sputtering in the hot gas between the reverse and forward shocks, which is hereafter referred to as the destruction by reverseshock."," In the course of their injection into ISM, dust grains formed in the ejecta are destroyed due to sputtering in the hot gas between the reverse and forward shocks, which is hereafter referred to as the destruction by reverseshock."
+" Nozawa investigated the survival of the newly formed dust in the shocked gas within the SNRs expanding into the uniform ISM with hydrogen number densities of Ίαν=0.1, 1.0 and 10.0cm-?, and showed that the destruction efficiency of newly formed dust is not only sensitive to the initial size distribution but also strongly depends on Ίαν."," \citet{Noz07} investigated the survival of the newly formed dust in the shocked gas within the SNRs expanding into the uniform ISM with hydrogen number densities of $n_{\rm SN}=0.1$, $1.0$ and $10.0\ {\rm cm^{-3}}$, and showed that the destruction efficiency of newly formed dust is not only sensitive to the initial size distribution but also strongly depends on $n_{\rm SN}$."
+" To investigate the dependence of destruction of dust on the ISM densities, we extend their models to six cases of ngw=0.03, 0.1, 0.3, 1.0, 3.0, and 10.0 cm?."," To investigate the dependence of destruction of dust on the ISM densities, we extend their models to six cases of $n_{\rm SN}=0.03$, $0.1$, $0.3$, $1.0$, $3.0$, and $10.0$ ${\rm cm}^{-3}$."
+" Figure 1 shows the IMF-averaged mass distribution of dust, M?""(a), injected into ISM after destruction by the reverse shock for ng= 0.03, 0.1, 0.3, 1.0, 3.0, and 10.0 cm-?, which is weighted by the Salpeter IMF and is summed up over all the grain species as in Equation (2))."," Figure \ref{fig:prod} shows the IMF-averaged mass distribution of dust, $\overline{{\cal M}_{{\rm d}}^{n_{\rm SN}}}(a)$, injected into ISM after destruction by the reverse shock for $n_{\rm SN}=0.03$ , $0.1$ ,$0.3$ , $1.0$ $3.0$ , and $10.0$ ${\rm cm}^{-3}$ , which is weighted by the Salpeter IMF and is summed up over all the grain species as in Equation \ref{proddustnoreverse}) )."
+ The mass distribution of the j-th dust species for the case of, The mass distribution of the $j$ -th dust species for the case of
+Chebyshev polynomials are defined for XC|1.1].,"Chebyshev polynomials are defined for $X\in [-1, 1]$."
+ Differential equations iu eoueral will be defined on a different interval. C[a.b].," Differential equations in general will be defined on a different interval $x\in [a,b]$."
+ Tn order to use Chebyshey polvuomials. one introduces a 1iappliug that maps thecoordinate c outo the X.," In order to use Chebyshev polynomials, one introduces a mapping that maps the $x$ onto the $X$."
+ Oue could explicitly substitute this mapping iuto the PDE under consideration., One could explicitly substitute this mapping into the PDE under consideration.
+ Derivatives would be multiplied by a Jacobian. aud we would obtain the PDE ou the interval |1.1].," Derivatives would be multiplied by a Jacobian, and we would obtain the PDE on the interval $[-1, 1]$."
+ For example. the differeutial equation in the variable vr becomes the followine differential equation iu the variable X: where X=OX/Or aud X=OPX042.," For example, the differential equation in the variable $x$ becomes the following differential equation in the variable $X$: where $X'=\partial X/\partial x$ and $X^{\prime\prime}=\partial^2X/\partial x^2$."
+ Now oue could expand «4(X) in Chebyshev polvuouials. compute derivatives 0/0.X via the reciurence relation and code Eq.," Now one could expand $u(X)$ in Chebyshev polynomials, compute derivatives $\partial/\partial X$ via the recurrence relation and code Eq."
+ in terms of du/ONX., in terms of $\partial u/\partial X$.
+ This approach is conmuou iu the literature [?.?]..," This approach is common in the literature \cite{Boyd:2001,Kidder-Finn:2000}."
+ However. it has several disadvantages: As one can already see from this sinple example. the equations become longer aud oue has to code and debug more ternis.," However, it has several disadvantages: As one can already see from this simple example, the equations become longer and one has to code and debug more terms."
+ Second. aud more miportaut. it is flexible. since for cach different map one has to derive and code a mapped equation(22).," Second, and more important, it is inflexible, since for each different map one has to derive and code a mapped equation."
+.. one might not know the appropriate map for a differential equation. aud in order to try several maps. one has to code the mapped equation several times.," one might not know the appropriate map for a differential equation, and in order to try several maps, one has to code the mapped equation several times."
+ Also. for domain decomposition. a different map is needed for cach subdomain.," Also, for domain decomposition, a different map is needed for each subdomain."
+ We propose a different approach., We propose a different approach.
+ We still expand in terms of Chebyshey polynomials on XC|1.1| aud obtain the physical solution via a mapping Vr). and we still compute Ou6X)/0.X. and OOuCXMOX? via the recurrence relation(LL).," We still expand in terms of Chebyshev polynomials on $X\in [-1, 1]$ and obtain the physical solution via a mapping $X(x)$, and we still compute $\partial u(X)/\partial X$ and $\partial^2u(X)/\partial X^2$ via the recurrence relation."
+.. However. now we do substitute QuCX)/0X and d2a(NX)/aN? into the mapped differential equation. Eq.," However, now we do substitute $\partial u(X)/\partial X$ and $\partial^2u(X)/\partial X^2$ into the mapped differential equation, Eq."
+(22).. Iustead we compute first nunericallv and substitute these values iuto the original plivsical differeutial equation(21)., Instead we compute first numerically and substitute these values into the original physical differential equation.
+. The collocation poiuts are thus mapped to the physical coordinates This approach separates the code into three distinct parts:, The collocation points are thus mapped to the physical coordinates This approach separates the code into three distinct parts:
+and Substitution of equations (17))-(21)) into equation (16)) gives the surface density. where c=mi/(GzBDyj) and Finally. it follows from equation (15)) that the emission measure of the thin shell (as a function of position) is given by with EM=σοριωτ’τω:GO23e.,"and Substitution of equations \ref{eq17}) \ref{eq21}) )into equation \ref{eq16}) ) gives the surface density, where $\sigma_0= \dot m_1/ (2\pi \beta D v_1)$ and Finally, it follows from equation \ref{eq15}) ) that the emission measure of the thin shell (as a function of position) is given by with $EM_0= \sigma_0\, \rho_{1,0}\, v_{1,0}^{2}/(2 \bar{m}^2 c_s^2)$."
+ To calculate the radio-continuum emission from the whole system. it is necessary to compute the total optical depth along each line of sight. which will have the contribution from the stellar winds and also from the thin shell.," To calculate the radio-continuum emission from the whole system, it is necessary to compute the total optical depth along each line of sight, which will have the contribution from the stellar winds and also from the thin shell."
+ Then we estimate the intensity emerging from each direction and calculate the flux by integrating the intensity over the solid angle., Then we estimate the intensity emerging from each direction and calculate the flux by integrating the intensity over the solid angle.
+" Given the emission measure EM(89.04) of the shocked layer (eq.(24))). the optical depth perpendicular to the WCR is caleulated. by το(0.06)=81)EM.v(v).. where ιν) 28436 x 107v-! with the frequency v in Hz (for an electron temperature 7,= 10'K)."," Given the emission measure $EM(\theta, \theta_1)$ of the shocked layer \ref{eq24}) )), the optical depth perpendicular to the WCR is calculated by $\tau_{WCR,\perp}(\theta, \theta_1)= EM(\theta, \theta_1)\,
+\chi(\nu)$, where $\chi(\nu)=$ 8.436 $\times$ $^{-7} \, \nu^{-2.1}$ with the frequency $\nu$ in Hz (for an electron temperature $T_e=$ $^4$ K)."
+ By defining a critical frequency. such that it can be shown from equation (24)) that where D= D/Rpo.," By defining a critical frequency, such that it can be shown from equation \ref{eq24}) ) that where $\tilde D= D/ R_0$ ."
+ For a line of sight intersecting the thin at an cos’!($-f) from the normal (being $2cosQRshell.—sin 80). anglethe optical depth is then given by where R is in units of Ro.," For a line of sight intersecting the thin shell at an angle $\mbox{cos}^{-1}\,(\hat{z} \cdot \hat{n})$ from the normal (being $\hat{z}= \mbox{cos}\,\theta \,\hat{R} - \mbox{sin}\,\theta
+\,\hat{\theta}$ ), the optical depth is then given by where $\tilde{R}$ is in units of $R_0$."
+ Let us now caleulate the contribution of the unshocked stellar winds to the optical depth., Let us now calculate the contribution of the unshocked stellar winds to the optical depth.
+ Consider a line of sight that intersects the thin shell at a point (7ctg8.7). as shown in Figure I.," Consider a line of sight that intersects the thin shell at a point $(\tilde{r}\, \mbox{ctg}\, \theta, \tilde{r})$, as shown in Figure 1."
+ According to Panagia & Felli (1975) and Wright & Barlow (1975). the optical depth along the line of sight ofthe wind source located at z=0 is obtained by where where nio.=pio/f (=my[AgriORs ) Is the number density of the flow at the stagnation point.," According to Panagia $\&$ Felli (1975) and Wright $\&$ Barlow (1975), the optical depth along the line of sight ofthe wind source located at $z= 0$ is obtained by where where $n_{1,0} = \rho_{1,0} / \bar{m}$ $= \dot{m_1}/ 4 \pi \bar{m}
+v_{1,0} R_0^2$ ) is the number density of the flow at the stagnation point."
+ From equations (25)). (28)). and (29)). it follows that with Analogously. the optical depth of the wind source located at >=Dis calculated by where where 29=poo/m (—mis/AnmvsoRg) is the number density of the wind at the stagnation point.," From equations \ref{eq25}) ), \ref{eq28}) ), and \ref{eq29}) ), it follows that with Analogously, the optical depth of the wind source located at $z= D$ is calculated by where where $n_{2,0} =\rho_{2,0}/\bar{m}$ $=\dot{m_2}/ 4 \pi \bar{m} v_{2,0} R_0^2$ ) is the number density of the wind at the stagnation point."
+ Substitution of (32)) into equation (31)) gives with We assumed that the system is far enough (D<L. being L the distance to the observer) so that the lines of sight intersecting the central stars can be ignored.," Substitution of \ref{eq32}) ) into equation \ref{eq31}) ) gives with We assumed that the system is far enough $D \ll L$, being $L$ the distance to the observer) so that the lines of sight intersecting the central stars can be ignored."
+" Then. the radio-continuum fluxfrom the colliding wind binary can be calculated by where r(0.0|)=τοῦ.01)+TCO)T2084) is the total optical depth along the line of sight. B,(=2kT17/c. with k the Boltzmann's constant. and ο the light speed) is the Planck function in the Rayleigh-Jeans approximation. and 7(94,) is the impact parameter at the asymptotic angle €, (corresponding to R25 oo) of the thin shell. which can be found from (see Cantó.. Raga & Wilkin 1996)."," Then, the radio-continuum fluxfrom the colliding wind binary can be calculated by where $\tau(\theta,\theta_1)= \tau_{WCR}(\theta, \theta_1) + \tau_{w,1} (\theta)
++ \tau_{w,2} (\theta)$ is the total optical depth along the line of sight, $B_{\nu}\,(= 2 k T \nu^2/c^2$ with $k$ the Boltzmann's constant, and $c$ the light speed) is the Planck function in the Rayleigh-Jeans approximation, and $\tilde{r}(\theta_{\infty})$ is the impact parameter at the asymptotic angle $\theta_{\infty}$ (corresponding to $R \rightarrow
+\infty$ ) of the thin shell, which can be found from (see Cantó,, Raga $\&$ Wilkin 1996)."
+ Defining the parameter f.=cLAnKT(u/v7. we finally obtain Equation (35)) represents an self-similar solution for the free-free radiation from colliding stellar winds from massive binary systems.," Defining the parameter $f_c =
+4 \pi k T (R_0/ c L)^2 \nu_c^2$, we finally obtain Equation \ref{eq35}) ) represents an self-similar solution for the free-free radiation from colliding stellar winds from massive binary systems."
+ We apply the model for computing the emission from a colliding wind binary using stellar wind parameterssimilar to those from a typical WR+O binary system., We apply the model for computing the emission from a colliding wind binary using stellar wind parameterssimilar to those from a typical WR+O binary system.
+" That is. we have adopted vj=I0? km ντι, in=1.25x107 Μ. yr?! for the O type star. and for the WR component. v»=10? km sl. dis=5x107 Μ. !'."," That is, we have adopted $v_1= 10^{3}$ km $^{-1}$ , $\dot m_{1} = 1.25
+\times 10^{-5}$ $_{\odot}$ $^{-1}$ for the O type star, and for the WR component, $v_2= 10^{3}$ km $^{-1}$ , $\dot m_{2} = 5 \times 10^{-5}$ $_{\odot}$ $^{-1}$ ."
+ In addition. we have assumed," In addition, we have assumed"
+optical emission is intrinsic faint due to low efficiency (e.¢.. Jakobsson οἱ al.,"optical emission is intrinsic faint due to low efficiency (e.g., Jakobsson et al."
+ 2005: Urata οἱ al., 2005; Urata et al.
+ 2007) or is heavily extinetecl by the dust (e.g.. Zheng et al.," 2007) or is heavily extincted by the dust (e.g., Zheng et al."
+ 2009: Hashimoto et al., 2009; Hashimoto et al.
+ 2010)., 2010).
+ It is emphasized that none of the GRBs listed in Table 1. can be classified as a “dark” burst., It is emphasized that none of the GRBs listed in Table 1 can be classified as a “dark” burst.
+ Table 2 lists (he radio ancl optical properties of three typical “dark” GRBs: 9970823. 9990506. and 0020819.," Table 2 lists the radio and optical properties of three typical “dark” GRBs: 970828, 990506, and 020819."
+ For each burst. we list the GRB identification. measured redshift. observing frequency in radio. observation (nme elapsed since the trigger. observed fhix density in radio. upper limit of R-band fIux density. aud (he caleulated monochromatic luminosity αἱ the correspondinge radio frequenev.," For each burst, we list the GRB identification, measured redshift, observing frequency in radio, observation time elapsed since the trigger, observed flux density in radio, upper limit of R-band flux density, and the calculated monochromatic luminosity at the corresponding radio frequency."
+ The data are compiled from literature as well., The data are compiled from literature as well.
+" Along with the A-correction. 1) the flux densities (and.also pL,) al 5GGllz are transformed. from (hat at GGllz by assuming a powerlaw spectral shape with slope a=—1/3 lor 9990506 and 0020319: 2) the upper limits of R-band flux densities are (ransformecl to V-band values by. assuming a spectral slope a=1.2 for all (he three evenis."," Along with the $k$ -correction, 1) the flux densities (andalso $\nu L_\nu$ ) at GHz are transformed from that at GHz by assuming a powerlaw spectral shape with slope $\alpha=-1/3$ for 990506 and 020819; 2) the upper limits of R-band flux densities are transformed to V-band values by assuming a spectral slope $\alpha=1.2$ for all the three events."
+ The three bursts are then over-plotted in Figure 1 bv (he green open triangles., The three bursts are then over-plotted in Figure 1 by the green open triangles.
+ The over-plotted. vertical arrows show the calculated spectral slope oo are underestimated because of the R-band upper limits., The over-plotted vertical arrows show the calculated spectral slope $\alpha_{\mathrm{RO}}$ are underestimated because of the R-band upper limits.
+ One can clearly see from Figure 1 that all the three “dark” bursts significantly deviate from the well established correlation (both optically bright GRB afterglows anc Blazars) αἱ a significant. level larger than 20 due to their large ao., One can clearly see from Figure 1 that all the three “dark” bursts significantly deviate from the well established correlation (both optically bright GRB afterglows and Blazars) at a significant level larger than $\sigma$ due to their large $\alpha_{\mathrm{OR}}$.
+ The deviation means these bursts are faint in optical bands with respect to not only X-ray. emission. bul also radio radiation.," The deviation means these bursts are faint in optical bands with respect to not only X-ray emission, but also radio radiation."
+ Brielly. there are (wo approaches thal are widely adopted to interpret the optically “dark” burst puzzle: one is to reduce the value of σος by enhancing (he A-ray emission. including long-livecl reverse shock lor N-rav. (e.g... Dai 2004: Genet et al.," Briefly, there are two approaches that are widely adopted to interpret the optically “dark” burst puzzle: one is to reduce the value of $\beta_{\mathrm{OX}}$ by enhancing the X-ray emission, including long-lived reverse shock for X-ray (e.g., Dai 2004; Genet et al."
+ 2007). late A-ray prompt emission (Ghisellini et al.," 2007), late X-ray prompt emission (Ghisellini et al."
+ 2007). ancl inverse Compton scattering (Pannaitescu 2003): an alternative approach is to suppress the optical radiation through heavy dust extinction or intrinsic faintness in optical bands.," 2007), and inverse Compton scattering (Pannaitescu 2008); an alternative approach is to suppress the optical radiation through heavy dust extinction or intrinsic faintness in optical bands."
+" The larger ago, suggests that both A-ray and radio enussion should be enhanced simultaneously for (he optically ""dark bursts in the first approach.", The larger $\alpha_{\mathrm{OR}}$ suggests that both X-ray and radio emission should be enhanced simultaneously for the optically “dark” bursts in the first approach.
+ In contrast. (he faint optical eniission in the “dark” bursts could be plausibly explained by the heavy dust extinction (e.g. Djorgovski οἱ al.," In contrast, the faint optical emission in the “dark” bursts could be plausibly explained by the heavy dust extinction (e.g., Djorgovski et al."
+ 2001). because the radio emission is believed to be [ree of the dust extinction.," 2001), because the radio emission is believed to be free of the dust extinction."
+ The extinction scenario was recently supported by Zheng et al. (, The extinction scenario was recently supported by Zheng et al. (
+"2009 and references (herein) who found the natural hydrogen column densitw is svstematically higher in oplically ""dark bursts than optically bright events.",2009 and references therein) who found the natural hydrogen column density is systematically higher in optically “dark” bursts than optically bright events.
+ The heavy extinction is also identified in the host galaxies of several “dark” GRBs (e.g. Levesque et al.," The heavy extinction is also identified in the host galaxies of several “dark” GRBs (e.g, Levesque et al."
+ 2010)., 2010).
+" In the context of the dust extinction scenario. we can roughlv estimate (he required extinction according to the ago—pL, correlation."," In the context of the dust extinction scenario, we can roughly estimate the required extinction according to the $\alpha_{\mathrm{RO}}-\nu L_\nu$ correlation."
+" Taking into account of the definition of the slope anoRCxlog(f,OUmj/Fie: f,,). we can express the GRB alterglow5 extinction;\ as a"," Taking into account of the definition of the slope $\alpha_{\mathrm{RO}}\propto\log(f_{\nu_1}/f_{\nu_2})$ , we can express the GRB afterglow extinction $A_V$ as a"
+Soft gamma-ray repeaters (SGRS) are believed to be a voung neutron star with strong magnetic fields (107 €). namely he magnetar (see. e.g.. Woods&Thompson2006:: Mereghetti 2008.. for review).,"Soft gamma-ray repeaters (SGRs) are believed to be a young neutron star with strong magnetic fields $\sim 10^{15}~\mathrm{G}$ ), namely the magnetar (see, e.g., \citealt{2006csxs.book..547W}; \citealt{2008A&ARv..15..225M}, for review)."
+ The magnetic fields inside the magnetar are amplified by the dvnamo mechanism at the birh of the neutron star., The magnetic fields inside the magnetar are amplified by the dynamo mechanism at the birth of the neutron star.
+ The Lorentz force stressing the crust of the magnetar dances wih the rigidity of the crust., The Lorentz force stressing the crust of the magnetar balances with the rigidity of the crust.
+ When the critical twist is accumulated. the magnetic twist injected into the magnetar magnetosphere will trigger the expansion of magnetic oops (Phompson&Duncan2001)..," When the critical twist is accumulated, the magnetic twist injected into the magnetar magnetosphere will trigger the expansion of magnetic loops \citep{2001ApJ...561..980T}."
+ Phe magnetic reconnection taking ace inside the expanding magnetic loops can be responsible for SCIL [lares (Woosetal.2001:Lyutikoy2006)..," The magnetic reconnection taking place inside the expanding magnetic loops can be responsible for SGR flares \citep{2001ApJ...552..748W, 2006MNRAS.367.1594L}."
+ Recently. relativistic simulations have been perforpmied to study the cdvnamies of the magnetospheres of neutron stars (Ixomissarov.2002: Asanoetal. 2005: Spitkovsky20)6:: Ixomissarov» 0062).," Recently, relativistic simulations have been performed to study the dynamics of the magnetospheres of neutron stars \citealt{2002MNRAS.336..759K}; \citealt{2005PASJ...57..409A}; ; \citealt{2006ApJ...648L..51S}; \citealt{2006MNRAS.367...19K}) )."
+ Spitkovsky(200!5) reported the results of relativistic force-free simulations o ‘the magnetar Dares triggered by th1e injection of the magnetic twis sat the ootpolnts ¢X the loops., \citet{2005KITP...CONF..HP} reported the results of 2-dimensional relativistic force-free simulations of the magnetar flares triggered by the injection of the magnetic twists at the footpoints of the loops.
+ WV‘hen the critical twist is accumulaCL. the magnetic looOS ονxui velativistically.," When the critical twist is accumulated, the magnetic loops expand relativistically."
+" Asano(2007) carried out 2-dimensional relativistic force-free simulations of exp:uxding magnetic OODS and srowed that the Lorenz factor defined bv the drift veloci yu,=eBD)/D? exceeds 10 (scanl""chida1997.fortvedeinitior1ofthedriftvelocity)."," \citet{2007Chiba...1..1} carried out 2-dimensional relativistic force-free simulations of expanding magnetic loops and showed that the Lorentz factor defined by the drift velocity $\bmath v_d = c(\bmath E \times \bmath B)/B^2$ exceeds 10 \citep[see,][for the definition of the drift
+velocity]{1997PhRvE..56.2181U}."
+ These simulations indicate that the magnetic loops expand self-simiarVe, These simulations indicate that the magnetic loops expand self-similarly.
+ Assuming relativistic orce-ree dynamics. Lyutikovy&BleIHieWord(2003). obtained se[-simiar solutions of the spherically expanding magnetic shell.," Assuming relativistic force-free dynamics, \citet{2003astro.ph.12347L} obtained self-similar solutions of the spherically expanding magnetic shell."
+ Prendergast(2005).} found selt-simil:ir solutions o£ the relativisic Loreοος field., \citet{2005MNRAS.359..725P} found self-similar solutions of the relativistic force-free field.
+ In these sucies of [orce-[ree dynamics. gas pressure and inertial terms are neglecος.," In these studies of force-free dynamics, gas pressure and inertial terms are neglected."
+ In the framework of the! pelativistic magnetohycdrocdvnamices (MILD). Lvutikov(2002) ouncl self-similar solutions of the spherically expandinge magneoic sheIs.," In the framework of the relativistic magnetohydrodynamics (MHD), \citet{2002PhFl...14..963L}
+ found self-similar solutions of the spherically expanding magnetic shells."
+ Low(1982). obtained self-similar MID solutions of the expanding maerjetic loops in solar Hares or supernovae explosion by assuming axisvmmetry., \cite{1982ApJ...261..351L} obtained non-relativistic self-similar MHD solutions of the expanding magnetic loops in solar flares or supernovae explosion by assuming axisymmetry.
+ Subsequentv. Low(1984) extended his. mode to the case including toroidal magnetic fields ancl applied. it to solar coronal mass ejecions (CALs).," Subsequently, \citet{1984ApJ...281..392L} extended his model to the case including toroidal magnetic fields and applied it to solar coronal mass ejections (CMEs)."
+ The latter model was emploved by Stonectal.(1992) as à testproblem to check, The latter model was employed by \cite{1992ApJ...388..415S} as a testproblem to check
+paper.,paper.
+ We employ RAMISES (Tevssier 2002). an adaptive mesh relinement cosmology code that solves (he equations of hydrodynamies using a second-order unsplit Godunov method and integrales particles using parlicleamesh techniques.," We employ RAMSES (Teyssier 2002), an adaptive mesh refinement cosmology code that solves the equations of hydrodynamics using a second-order unsplit Godunov method and integrates particles using particle-mesh techniques."
+" The code has been augmented to follow nonequilibrium chemistry for a nine-species gas: e. HH. WI. Wel. Hell. Well]. . Ibs. and IL,."," The code has been augmented to follow nonequilibrium chemistry for a nine-species gas: e, HI, HII, HeI, HeII, HeIII, $^{\rm -}$ , $_2$ , and $_2^+$."
+ Cooling due to collisions between II» and the other species is modeled according to Glover Abel (2008)., Cooling due to collisions between $_2$ and the other species is modeled according to Glover Abel (2008).
+ The Dubois Tevssier (2003) star formation recipe is modified to capture pop ILE star formation ancl supernova leedback., The Dubois Teyssier (2008) star formation recipe is modified to capture pop III star formation and supernova feedback.
+ Wherever the gas number density is ereater than the critical value ny=I0?em7. there is a chance to form a star particle.," Wherever the gas number density is greater than the critical value $n_0=10^5 \rm cm^{-3}$, there is a chance to form a star particle."
+ The local SFE is set to per cdvnamical time., The local SFE is set to per dynamical time.
+" If the metallicity of the gas is above Zi=5x10""Z.. the resulting star particle represents a distribution of population IL stars."," If the metallicity of the gas is above $Z_{\rm thresh}=5\times10^{-7} Z_{\odot}$, the resulting star particle represents a distribution of population II stars."
+ This threshold is lower than suggested by Dromm Loeb (2003) because we wanted to explore whether very low-metallicity population ll stars could form if given the opportunity., This threshold is lower than suggested by Bromm Loeb (2003) because we wanted to explore whether very low-metallicity population II stars could form if given the opportunity.
+ Ten percent of (he star particle's mass is injected back into medium alter 10 Myr as supernova ejecta. with 0.1 of the ejecta mass in metals.," Ten percent of the star particle's mass is injected back into medium after 10 Myr as supernova ejecta, with 0.1 of the ejecta mass in metals."
+ Below Zia. the particles represent a distribution of population LI stars.," Below $Z_{\rm thresh}$, the particles represent a distribution of population III stars."
+ The IMF for population III stars remains unclear., The IMF for population III stars remains unclear.
+ The minimum mass may be 2LOA. (e.g. Dromm et al.," The minimum mass may be $\gtrsim 10 M_{\odot}$ (e.g, Bromm et al."
+ 2009). but this remains highly uncertain.," 2009), but this remains highly uncertain."
+ Likewise. we know neither ihe maxinum mass lor a pop IH star nor (he importance of pair-instabilitv. supernovae io the enrichment of the second. generation of stars.," Likewise, we know neither the maximum mass for a pop III star nor the importance of pair-instability supernovae to the enrichment of the second generation of stars."
+ Tumlinson et al. (, Tumlinson et al. (
+2004) argue that the observed. element ratios of old pop IL stars are inconsistent wilh enrichment. by pair- supernovae.,2004) argue that the observed element ratios of old pop II stars are inconsistent with enrichment by pair-instability supernovae.
+ Moreover. (he IME shows little sensiivily (to metallicity (Ixroupa 2001).," Moreover, the IMF shows little sensitivity to metallicity (Kroupa 2001)."
+ Whether (his extends to extremely low metallicities remains unknown., Whether this extends to extremely low metallicities remains unknown.
+" Iu (he absence of better constraints on pop HT star formation. we assume a Salpeter LMF (4.N/dM~AL 779) with a low-mass cutoff AL,CO (see below)."," In the absence of better constraints on pop III star formation, we assume a Salpeter IMF $dN/dM\sim M^{-2.35}$ ) with a low-mass cutoff $M_{\rm co}$ (see below)."
+" Whenever a star particle is formed. (he particle's nass Ma, is distributed among a random distribution of stars."," Whenever a star particle is formed, the particle's mass $M_{\rm sp}$ is distributed among a random distribution of stars."
+" The probability that a star is between masses M, and M, is given by P=ACA,EF?—M,I). where d ⇀⋃∙⊲⊓↓↝⋮⋡⋝⊽↴↕⋝↓≀↧↴∐≺⇂⇀∪↿⊓↓⋝∶⊥∩∪∪⇀∪⋅↕⊳∖⇁⊔∐↲∐⋯⊳∖⊽↥∐↓≀↧↪∖⊽⊳∖⇁↕∖↽≼↲⊳∖⇁↥≀↧↴↕⋅∐↥⊔∐↲≼∐⊳∖⊽∏⋅↕∣↽≻∏∐∪∐⋅↼≚≸≟↕∖↽≼↲∐⋝∖⊽↥≀↧↴↕⋅ ∐↓≀↧↪∖⇁⋟∖⇁↥⋟∖⇁⋟∖⇁≼↲↥≼↲≺∢∥↲≼⇂∣↽≻∡∖"," The probability that a star is between masses $M_0$ and $M_1$ is given by $P=A(M_1^{-1.35}-M_0^{-1.35})$, where $A=(M_{\rm top}^{-1.35}-M_{\rm co}^{-1.35})^{-1}$ and $M_{\rm top}=1000 M_{\odot}$ is the most massive star in the distribution."
+⇁⋟∖⇁≼↲⊔↕∐≸≟⇀⋃⊓↿⋥⋃⋅∶≼↜∣⋮∕∕∕⇀∣↥⊹⇀∪∙⊲⊓↓⋡⋮⋝⋝⊽↴↕⋟↓∡⋮⋝⋝⊽↴⋅∖∖⊽↥∐↲↕⋅≼↲∣⋮↕⋝∖⊽≀↧↴↕⋅≀↧↴∐≼⇂∪↕∐∖↽≀↧↴↕⋅↕≀↧↴∣↽≻↥≼↲∖∖⊽↕⊔↥ ≀↕↴∐≼↲∖↽≼↲∐≺∐⋟∖⊽∏⋅↕∣↽≻∏∐∪∐∣↽≻≼↲↥∖∖⊽≼↲≼↲∐∩≀↧↴∐≺⇂↓⋅↴∏∐⋟∖⊽↕↽≻↕⋅," A given star mass is selected by setting $M_{\rm star}=(r/A+M_{\rm co}^{-1.35})^{-1/1.35}$, where $r$ is a random variable with an even distribution between 0 and 1."
+∪≺∢≼↲⋟∖⊽⋟∖⊽≺∢∪∐∐∐⋯↲⋟∖⊽∏∐↥∐⊔∐↲↥∪↥≀↧↴↥∐⋯⊔∖∷∖⊽↕∐⋟∖⊽↥≀↕↴↕⋅⋟∖⊽↕⋟∖⊽ ≼↲≺↥∏≀↧↴⊓∪⇀∪⊓⋝⋅∖∖⊽≼↲⊳∖, This process continues until the total mass in stars is equal to $M_{\rm sp}$.
+⇁≼↲↥⇀∪⋯∶⊑↽⊰⇀∪⋅⋅∖∖↽∐↕≺∢∐↕⊳∖⇁∣↽≻≀↧↪∖⇁≼↲≼⇂∪∐⊔∐↲⊔∐↲∪↕⋅≼↲∐≺∢≀↧↴↥≺∢≀↧↴↥≺∢↕∐≀↧↴∐∪∐⊳∖⇁∪↓∟∖↡≀↧↴↳↽≀↧↴∐↓⋯⋅≀↧↴ Umenmura (2001). who find that the smallest Jeans mass for which the gas becomes opaque to II» lines is ~LAL..," We set $M_{\rm co}=3 M_{\odot}$, which isbased on the theoretical calculations of Nakamura Umemura (2001), who find that the smallest Jeans mass for which the gas becomes opaque to $_2$ lines is $\sim 1 M_{\odot}$."
+ For each population HI star in (he distribution. its contribution to ejecta. metal enrichment. long-lived stars and remnants. pair-instabilitv supernovae. aud black holes is tracked.(seeTable 1).," For each population III star in the distribution, its contribution to ejecta, metal enrichment, long-lived stars and remnants, pair-instability supernovae, and black holes is tracked(seeTable 1)."
+" The remnant neutron star or black hole mass is determined. when applicable. by Αμ=MANALgOLEMMes135M., )."," The remnant neutron star or black hole mass is determined, when applicable, by $M_{\rm R/BH}={\rm MAX}( M_{\rm star} 0.1^{13.5 M_{\odot}/M_{\rm star}},\ 1.35 M_{\odot} )$ ."
+ This prescription roughly follows the results, This prescription roughly follows the results
+The projected velocity dispersion is calculated using (Woodley οἱ al.,The projected velocity dispersion is calculated using (Woodley et al.
+ 2007). where N is (he number of clusters in each group of GCs. found as à result of the CA. vy; is the GC's radial velocity after subtraction of rotational component determined with equation (C1) and σι is the projected velocity dispersion.," 2007), where N is the number of clusters in each group of GCs, found as a result of the CA, $v_{fi}$ is the GC's radial velocity after subtraction of rotational component determined with equation (C1) and $\sigma_{v}$ is the projected velocity dispersion."
+"Together with the timing analysis on the evolution of the rms along the outburst and its energy dependence, we have also studied the evolution of the main ΟΡΟ properties.","Together with the timing analysis on the evolution of the rms along the outburst and its energy dependence, we have also studied the evolution of the main QPO properties."
+ In table 2 we present the fits for the 43 PDS in which Type-C QPOs were observed and the 9 PDS with a type-B QPO., In table \ref{tab:QPO} we present the fits for the 43 PDS in which Type-C QPOs were observed and the 9 PDS with a type-B QPO.
+" Only in one observation (95118-01-06-00) we see a possible type-A QPO, although its significance is low and we will not consider it in our analysis."," Only in one observation (95118-01-06-00) we see a possible type-A QPO, although its significance is low and we will not consider it in our analysis."
+ PDS fitting was carried out with the standard fitting package by using a one-to-one energy-frequency conversion and a unit response., PDS fitting was carried out with the standard fitting package by using a one-to-one energy-frequency conversion and a unit response.
+" Following ?,, we fitted the noise components with three Lorentzians, one zero-centred and other two centred at a few Hz."," Following \cite{Belloni2002}, we fitted the noise components with three Lorentzians, one zero-centred and other two centred at a few Hz."
+" The QPOs were fitted with one Lorentzian each, only occasionally needing the addition of a Gaussian component to better approximate the shape of the narrow peaks and to reach values of reduced x? close to |."," The QPOs were fitted with one Lorentzian each, only occasionally needing the addition of a Gaussian component to better approximate the shape of the narrow peaks and to reach values of reduced $\chi ^2 $ close to 1."
+" The behaviour of both Type-B and Type-C QPOs is similar to that observed in other BHT (see e.g., ?))."," The behaviour of both Type-B and Type-C QPOs is similar to that observed in other BHT (see e.g., \citealt{Belloni2010}) )."
+" We see the type-C frequency increasing with hardness, whereas type-B are always observed within the frequency range ~ 2-4 Following ? and Motta et al. ("," We see the type-C frequency increasing with hardness, whereas type-B are always observed within the frequency range $\sim$ 2–4 Following \cite{Casella2004} and Motta et al. ("
+2011 in prep.),2011 in prep.)
+ we have plotted total rms as a function of the QPO frequency (Fig. 7))., we have plotted total rms as a function of the QPO frequency (Fig. \ref{qpo}) ).
+" As it was found in those works, Type-C QPOs follow a clear negative correlation."," As it was found in those works, Type-C QPOs follow a clear negative correlation."
+ This correlation seems to saturate around ~7.5 Hz., This correlation seems to saturate around $\sim 7.5$ Hz.
+" Interestingly, if we only consider the Type-C QPOs observed during the soft-to-hard transition (open circles in Fig. 7))"," Interestingly, if we only consider the Type-C QPOs observed during the soft-to-hard transition (open circles in Fig. \ref{qpo}) )"
+ a slightly higher slope is observed in the Type-B QPOs are unequivocally separated from Type-C using the rms-frequency representation and they are solely observed when the rms is «10 per cent., a slightly higher slope is observed in the Type-B QPOs are unequivocally separated from Type-C using the rms-frequency representation and they are solely observed when the rms is $\leq 10$ per cent.
+" Time-lags between soft and hard variability were computed for the first observation (95358-01-02-00), the only one long enough for this purpose."," Time-lags between soft and hard variability were computed for the first observation (95358-01-02-00), the only one long enough for this purpose."
+ Following ? we used the energy ranges ~ 2-4 keV and ~ 8-13 keV for the soft and hard bands., Following \cite{Pottschmidt2000} we used the energy ranges $\sim$ 2–4 keV and $\sim$ 8–13 keV for the soft and hard bands.
+" They corresponds to the STD2 channels 0-5 and 15-27, respectively."," They corresponds to the STD2 channels 0–5 and 15–27, respectively."
+ Fast Fourier transforms of each band were computed and cross spectra produced., Fast Fourier transforms of each band were computed and cross spectra produced.
+ A positive lag means hard variability lagging soft variability., A positive lag means hard variability lagging soft variability.
+ The obtained time-lags (At) as a function of frequency (v) are shown in Fig., The obtained time-lags $\Delta t$ ) as a function of frequency $\nu$ ) are shown in Fig.
+ 8 (black circles)., \ref{tlags} (black circles).
+" As observed in previous works, the time-lag decreases with frequency within the 0.1-10 Hz band to which we are sensitive."," As observed in previous works, the time-lag decreases with frequency within the 0.1-10 Hz band to which we are sensitive."
+ ? noted that in Cyg X-1 the time-lags were larger during the transition between hard and soft states as compared to those observed in the canonical states., \cite{Pottschmidt2000} noted that in Cyg X-1 the time-lags were larger during the transition between hard and soft states as compared to those observed in the canonical states.
+ The system was in the HIMS during observation #11 and in agreement with the above work the lags we measure are larger than those typically observed in BHT during LHS., The system was in the HIMS during observation 1 and in agreement with the above work the lags we measure are larger than those typically observed in BHT during LHS.
+ For a direct comparison we have computed time-lags for a LHS observation (diamonds in Fig.8)) and two HIMS observations at the same hardness (triangles in Fig.8)) and same fractional rms (open circles in Fig.8)) corresponding to the 2007 outburst of GX 339-4., For a direct comparison we have computed time-lags for a LHS observation (diamonds in \ref{tlags}) ) and two HIMS observations at the same hardness (triangles in \ref{tlags}) ) and same fractional rms (open circles in \ref{tlags}) ) corresponding to the 2007 outburst of GX 339-4.
+" Whereas the time-lags corresponding to the LHS observation follows the relation At~0.009xv~°7 (dashed line in Fig.8)), consistent with the one observed in other systems during that state (see e.g., ? for Cyg X-1 and XTE1752-223) a normalization at least two times bigger (~0.02; dashed line in Fig.8)) is needed to account for the time- observed in MAXI J1659-152."," Whereas the time-lags corresponding to the LHS observation follows the relation $\Delta t \sim 0.009\times\nu^{-0.7}$ (dashed line in \ref{tlags}) ), consistent with the one observed in other systems during that state (see e.g., \citealt{tmd2010} for Cyg X-1 and XTE1752-223) a normalization at least two times bigger $\sim 0.02$; dashed line in \ref{tlags}) ) is needed to account for the time-lags observed in MAXI J1659-152."
+ We note that deviations from the power-law seems to be present at low frequencies., We note that deviations from the power-law seems to be present at low frequencies.
+ The evolution of MAXI J1659-152 during its 2010 outburst is consistent with that usually observed in black hole transients., The evolution of MAXI J1659-152 during its 2010 outburst is consistent with that usually observed in black hole transients.
+" The hardness-intensity, rms-hardness and rms-intensity diagrams are rather typical, with a hard-to-soft transition, a flux decay during a soft (accretion disc dominated) state and a final soft-to-hard transition toward the quiescence."," The hardness-intensity, rms-hardness and rms-intensity diagrams are rather typical, with a hard-to-soft transition, a flux decay during a soft (accretion disc dominated) state and a final soft-to-hard transition toward the quiescence."
+ Variability is in general correlated with hardness., Variability is in general correlated with hardness.
+" Its energy spectrum is hard during the HIMS and the HSS and flat during the If we compare the hardness values with those observed in other systems, we see that MAXI J1659-152 has rather hard spectrum."," Its energy spectrum is hard during the HIMS and the HSS and flat during the If we compare the hardness values with those observed in other systems, we see that MAXI J1659-152 has a rather hard spectrum."
+ This is clear when looking at the spectral parametersa that we obtain, This is clear when looking at the spectral parameters that we obtain
+with (he available data?,with the available data?
+ Inverting the question. can the mass distribution be inferred from the kinematic data?," Inverting the question, can the mass distribution be inferred from the kinematic data?"
+ The point is not to find the best fit involving every. possible kind οἱ data (see Widrow. Pym. Dubinski 2008 for such an exercise in the conventional context).," The point is not to find the best fit involving every possible kind of data (see Widrow, Pym, Dubinski 2008 for such an exercise in the conventional context)."
+ Rather. I seek to use one tvpe of data to an entirely independent set of observations.," Rather, I seek to use one type of data to an entirely independent set of observations."
+ The paper is organized as follows., The paper is organized as follows.
+ In section 82.. 1 describe MOND as applied here. considering a variety of interpolation functions.," In section \ref{mus}, I describe MOND as applied here, considering a variety of interpolation functions."
+ In 83. I describe (he mass distribution of the Milky Way. adopting the mmocdel (Flynn 22006) for the stellar mass distribution. and utilizing the tabulation of Olling Alerrifield (1999) for the eas distribution.," In \ref{MWmass} I describe the mass distribution of the Milky Way, adopting the model (Flynn 2006) for the stellar mass distribution, and utilizing the tabulation of Olling Merrifield (1999) for the gas distribution."
+ In 84. I present the results of applying MOND to the adopted mass model., In \ref{MWRC} I present the results of applying MOND to the adopted mass model.
+ In 85. Linvert the question. ancl infer the deviations from a smooth exponential huninositwv profile implied by the bumps ancl wigegles in the terminal velocity curve.," In \ref{MW_emp} I invert the question, and infer the deviations from a smooth exponential luminosity profile implied by the bumps and wiggles in the terminal velocity curve."
+ The conclusions are summarized in &G.., The conclusions are summarized in \ref{conc}.
+ The basic idea of MOND (Milerom 1933a) is that rather (han invoking dark matter to explain mass discrepancies in extragalactic svstems. one mocdifies the force law such that where gy is (he acceleration calculated using Newtonian dynamics. and a is the actual resultant acceleration.," The basic idea of MOND (Milgrom 1983a) is that rather than invoking dark matter to explain mass discrepancies in extragalactic systems, one modifies the force law such that where $_N$ is the acceleration calculated using Newtonian dynamics, and $\textbf{a}$ is the actual resultant acceleration."
+ The modification occurs not at some length scale. but al the acceleration scale αμ. with cc=a/ay (here a= jal}.," The modification occurs not at some length scale, but at the acceleration scale , with $x = a/\anot$ (here $a = |\textbf{a}|$ )."
+ The interpolation function j(r) has the property Γι Lin the limit of large accelerations @Xαμ so that Newtonian behavior is recovered., The interpolation function $\mu(x)$ has the property $\mu \rightarrow 1$ in the limit of large accelerations $a \gg \anot$ so that Newtonian behavior is recovered.
+ In the limit of small accelerations the modification applies. with por lor a«ety.," In the limit of small accelerations the modification applies, with $\mu \rightarrow x$ for $a \ll \anot$."
+ The acceleration scale iis very small., The acceleration scale is very small.
+ Degeman uu ↜⊥≤∍≤∐⊔∪∏∐≼⇂∩∣∣∶−⋡⋅−≻⋗⋖⊥∩↓∎∎∐↓⋟∖⇁−∶≡↽⊰↙∩∩↳↽∐↓⋟∖⊽−↳↽↕↽≻≺∢↓↕⋅∪∐↓∐↥⋟∖⊽↥∪↥∐↖≺↽↔↴∐≺⇂∏≀↧↴∐↥⋡∖↽↕⋅∪↥≀↕↴⊔∪∐.2ye9 Lp . ⋅ ⋅ ⋅ ∢∏↕⋅∖↽≼↲⊳∖⊽⋅⇀∖↕∪↕⋅≼↲↕⋅≼↲≺∢≼↲∐," Begeman (1991) found $\anot = 1.2 \times 10^{-10}\;
+\textrm{m}\,\textrm{s}^{-2} = 
+3700\;\textrm{km}^2\,\textrm{s}^{-2}\,\textrm{kpc}^{-1}$ from fits to high quality rotation curves."
+↥∡∖⇁⋅∐∪⊔≼↲∐↓≀↧↴≼↲↥≀↧↴↥⋅ ⋝⊋∩∪∃⇄⋝≼↲⊳∖⇁∐∐↓≀↧↴∩↲≼⇂∩∣∣∶⊑↽⊰∪∪∪↳↽↕∐⇉⋝∖⊽⊋↳↽↕↽≻≺∢↓≀↧↴∐≼⇂⇀∖↕≺∢⊂↽⊐≀↧↴∏≸≟∐⊔∪∪⊥⋟∩∣∣≈∔∩∩∪↕≶∐↓⇄⊳∖⇁⊋↳↽↕↽≻≺∢ ∐↲↥≀↧↴⊔≼↲↕⋅∏∐∐∠↕∐≸↽↔↴↕↽≻∪↕↽≻∏↥≀↧↴∐∪∐⊳∖⊽∡∖↽∐⊔∐↲⋝∖⊽↕⊳∖⇁≼↲⊳∖⊽∐∐↓≀↧↴∥↲⋝∖⊽∪↓⋟⊔∐↲⊳∖⇁∥↲∐≀↧↴↕⋅↕∐≀↧↪∖⋱∖⊽⋅∐≼↲↕⋅≼↲↕⋯⇂∪↕↽≻↥⊔∐↲ ∎∐∩↲↕⋅∐∐↲≺∐≀↕↴∩↲∐≼↲↖≺↽↔↴≼↲∐," More recently, Bottema (2002) estimated $\anot = 3000\;\textrm{km}^2\,\textrm{s}^{-2}\,\textrm{kpc}^{-1}$ and McGaugh (2004) $\anot \approx 4000\;\textrm{km}^2\,\textrm{s}^{-2}\,\textrm{kpc}^{-1}$, the latter utilizing population synthesis estimates of the stellar mass."
+↓≀↧↴∐≼↲↥≀↕↴↥⋅ ⋝⊥≤∍≤∐⇄⋝∖↽≀↧↴↥⋯↲≀↧↴∐≼∐↘↽≼↲≼↲↕↽⊓⇃∣∣ fixed while considering different possible interpolation functions., Here I adopt the intermediate Begeman (1991) value and keep fixed while considering different possible interpolation functions.
+ Thereader should bear in wind that in addition to the uncertaintv in, Thereader should bear in mind that in addition to the uncertainty in
+xobabilitv of having JV eravitationally bouud structures on a scale larger than Π is given bv the probability law PUNE)—NiRyxRἨ.,probability of having $N$ gravitationally bound structures on a scale larger than $R$ is given by the probability law ${\cal P}(N_{R})\equiv N(>R)\propto R^{-D}$.
+ In Fourier space. this simply 110115 hat he uuuber of fluctuations of wavenuuber ko1/R is proportional to the volume of a fluctuation. Lo. ING)xdk=bPlak.," In Fourier space, this simply means that the number of fluctuations of wavenumber $k\sim 1/R$ is proportional to the volume of a fluctuation, i.e. $dN(k)\propto d{\vec k}=k^{D-1}dk$."
+ The density probability. 1.6. he probability to have a structure of size C[RR|dH]. lius reads as drA homogeneous volume iu 3; dieusious obviously corresponds to D=3. Le. a uniform density probability IN(kh)dk=.constant. as assunied esso d Pacoan Nordlund (2002).," The density probability, i.e. the probability to have a structure of size $\in[R,R+dR]$, thus reads as A homogeneous volume in 3 dimensions obviously corresponds to $D=3$, i.e. a uniform density probability $dN(k)/d{\vec k}=constant$, as assumed e.g. in Padoan Nordlund (2002)."
+ From Eqs. CD) , From Eqs. \ref{Mturb}) )
+"and (5)). the probability of having a bound structure of mass C.LALOR).AL|dAI(RC dR). which defines the CAIF/IME. reads As scon from Eq.(6)). for D=5. pure thermal support (9=3.4 0) yields a power-law IME D~AP1, substantially stecper than the Salpeter value. a=2.351."," and \ref{N}) ), the probability of having a bound structure of mass $\in[M(R),M+dM(R+dR)]$ , which defines the CMF/IMF, reads As seen from \ref{Nturb}) ), for $D=3$, pure thermal support $n=3,\eta =0$ ) yields a power-law IMF $\frac{dN}{dM}\propto M^{-4}$, substantially steeper than the Salpeter value, $\alpha=2.35$."
+", Iu the two liniting cases of incompressible (Nolmoeorov) (0=1l/3.4g 1/3) aud pressurcless (Burgers) (50=lj= 1/2) turbulence. the turbulent support vields a=28 and a=2.5. respectively."," In the two limiting cases of incompressible (Kolmogorov) $n=11/3,\eta =1/3$ ) and pressureless (Burgers) $n=4,\eta =1/2$ ) turbulence, the turbulent support yields $\alpha=2.8$ and $\alpha=2.5$, respectively."
+ In molecular clouds. urbuleuce is assumed to scale according to the observe Larson (1981) velocity dispersion - size relation with a vpieal value jj~ 0.10.5.," In molecular clouds, turbulence is assumed to scale according to the observed Larson (1981) velocity dispersion - size relation with a typical value $\eta\sim 0.4$ -0.5."
+ This value is well recoverce or the aforementioned power spectrum iudex iduferre roni nunierical simulations of supersonic turbulence. 3.8.," This value is well recovered for the aforementioned power spectrum index inferred from numerical simulations of supersonic turbulence, $n=3.8$ ."
+" According to equ.(6)). this vields a slope for the CMFE/IMEF «aΞ-2.66, between the Iohnuogorov ai Bureers values."," According to \ref{Nturb}) ), this yields a slope for the CMF/IMF $\alpha=2.66$, between the Kolmogorov and Burgers values."
+ There are observational sugeestious.OO however. that uolecular clouds have a fractal dimension (Ehueereeun Falearoue 1996). with D~ 2.5-2.7 (Samuchez ct al.," There are observational suggestions, however, that molecular clouds have a fractal dimension (Elmegreen Falgarone 1996), with $D\approx 2.5$ -2.7 (Sánnchez et al."
+ 2005)., 2005).
+ Herschel observations of eravitationally bound prestellar cores. on the other haud. suggest a uiass-size relation AJxwBU. with zm 1-2 (Fig.," Herschel observations of gravitationally bound prestellar cores, on the other hand, suggest a mass-size relation $M\propto R^\beta$, with $\beta\approx 1$ -2 (Fig."
+ doof André et al, 4 of André et al.
+ 2010 and Exóunyvves et al., 2010 and Könnyves et al.
+ 2010)., 2010).
+ Combined with a SalpeteLad nass distribution. this vields dN/dRXRUNE) wit1 Dzz 1.1-2.7.," Combined with a Salpeter mass distribution, this yields $dN/dR \propto R^{-(1+D)}$ with $D\approx 1.4$ -2.7."
+ Althoueh these results should be taken with caution. they suggest a structured space with D<3 all the wav from clouds to cores.," Although these results should be taken with caution, they suggest a structured space with $D<3$ all the way from clouds to cores."
+" This is consistent with turbulence being the dominant mechanisin that determines the distribution of both uibound aud bound structures,", This is consistent with turbulence being the dominant mechanism that determines the distribution of both unbound and bound structures.
+ It arises. however. froin two differeut properties of turbulence.," It arises, however, from two different properties of turbulence."
+ As demonstrated in ICs. the structure of theunbound clouds/clumps arises from the structure of turbulence. characterized by a lognormal PDF alc a power spectruii iudex i’.," As demonstrated in HC08, the structure of the clouds/clumps arises from the structure of turbulence, characterized by a lognormal PDF and a power spectrum index $n^\prime$."
+ Ou the other haud. the distribution of eravitationallybound structures is set up by the virial coudition aud thus involves the properties of the structure of turbulence (see Eq. (2))).," On the other hand, the distribution of gravitationally structures is set up by the virial condition and thus involves the properties of the structure of turbulence (see Eq. \ref{vir}) )),"
+ characterized by a power spectrum index b., characterized by a power spectrum index $n$.
+" As mentioned above. both indexes appear to be similar for supersonic turbulence, with voznc3.8 (Ixitsuk et al."," As mentioned above, both indexes appear to be similar for supersonic turbulence, with $n^\prime\simeq n\simeq 3.8$ (Kritsuk et al."
+ 2007)., 2007).
+ Although the mass-size relation of (evolved) prestellar cores cannot be directly applied to the one characteristic of the iuitia mass reservolrs. it is interesting. however. to note that a value n=3.8 in Eq. (1))," Although the mass-size relation of (evolved) prestellar cores cannot be directly applied to the one characteristic of the initial mass reservoirs, it is interesting, however, to note that a value $n=3.8$ in Eq. \ref{Mturb}) )"
+ vields a relation for the bouud structures supported by turbulence APxRES (uoo Eq.(13) and below in WCO0s for the colplete relation). consistent with the aforemeutioned observationa determination from Ilerschel.," yields a relation for the bound structures supported by turbulence $M\propto R^{1.8}$ (see Eq.(43) and below in HC08 for the complete relation), consistent with the aforementioned observational determination from Herschel."
+ luterestingly. using the aforementioned observationallv iuforred value Dzx2.5-2.7 vields a=3.5-3.7. according to Eq.(6)). in the case of purely thermal support. which is still significantly steeper than the Salpeter value. a=2.25-2.35 and 2.5-2.6 in the Durgers and Eohuosgorov limits. respectively. and a=2.39-2.5 when using the aforciuentioned value for the supersonic turbulence index.," Interestingly, using the aforementioned observationally inferred value $D\approx 2.5$ -2.7 yields $\alpha=3.5$ -3.7, according to \ref{Nturb}) ), in the case of purely thermal support, which is still significantly steeper than the Salpeter value, $\alpha=2.25$ -2.35 and 2.5-2.6 in the Burgers and Kolmogorov limits, respectively, and $\alpha=2.39$ -2.5 when using the aforementioned value for the supersonic turbulence index."
+ The complete calculations. ax done in IICOS aud IICU. vield of course more accurate determinations of the slope of the CALF. both in the thermal aud. nonthermal case.," The complete calculations, as done in HC08 and HC09, yield of course more accurate determinations of the slope of the CMF, both in the thermal and nonthermal case."
+ It is nuportaut to stress that. iu the preseut simplistic approach. one a space distribution given by Eq. (5)).," It is important to stress that, in the present simplistic approach, one a space distribution given by Eq. \ref{N}) ),"
+ which leads to the (D-3) correction in Eq. (6))., which leads to the $D$ $3$ ) correction in Eq. \ref{Nturb}) ).
+ Tn he rigorous (Press-Schechter based) statistical approach developed in ΠοΟδ and IICO9. the dimension D. both or the (inbound) chumps aud the (bound) cores is cutering the DIC theory.," In the rigorous (Press-Schechter based) statistical approach developed in HC08 and HC09, the dimension $D$, both for the (unbound) clumps and the (bound) cores is entering the HC theory."
+ There is 10 need to explicitly assume a fractal index. so there is 10 explicit dependence of the Salpeter slope upon such an iudex.," There is no need to explicitly assume a fractal index, so there is no explicit dependence of the Salpeter slope upon such an index."
+ On the contrary. the size distribution naturally arises frou the calculations with (as easily interred frou eqn.(12) of [COs aud the mass-size relation 1)) - detailed in HCUS - derived from the viral coudition): dN/dRxRO= RR ie Do-2d consistent with the aforementioned Ierschel determinations.," On the contrary, the size distribution naturally arises from the calculations with (as easily inferred from eqn.(42) of HC08 and the mass-size relation \ref{Mturb}) ) - detailed in HC08 - derived from the viral condition): $ dN/dR \propto R^{-{n+3\over 2}}\approx R^{-3.4}$ , i.e. $D\approx 2.4$, consistent with the aforementioned Herschel determinations."
+ lu contrast. a mre thermal support (0qu.(3)) aud η= 3) iu the complete IIC theory vields 4N/d4RxR2773 Rie D=2.," In contrast, a pure thermal support \ref{therm}) ) and $n=3$ ) in the complete HC theory yields $ dN/dR \propto R^{-{3n-3\over 2n-4}}\propto R^{-3}$, i.e. $D=2$."
+" The sometimes advocated depeudence of the IAIF upon he spatial dimension D. aud the fractional value of the . thus simply reflect the (stabilizing) impact of the urbuleut velocity Ποια ou large-scale bow structures,"," The sometimes advocated dependence of the IMF upon the spatial dimension $D$, and the fractional value of the latter, thus simply reflect the (stabilizing) impact of the turbulent velocity field on large-scale bound structures."
+ The simpledimensional relations derived in this short rote thus enable us to erasp the crucial impact of urbuleut support. combined with eravitv through the virial condition. ou the shape of the IME.," The simpledimensional relations derived in this short note thus enable us to grasp the crucial impact of turbulent support, combined with gravity through the virial condition, on the shape of the IMF."
+ Indeed. based on simple scaling arguineuts. these relations demonstrate hat nonthermal support is mandatory to vield the correct," Indeed, based on simple scaling arguments, these relations demonstrate that nonthermal support is mandatory to yield the correct"
+In some of these cases both Μία=0.3) ancl f(a=0.5) exceed unity by about 20 or ereater.,In some of these cases both $R(a=0.3)$ and $R(a=0.5)$ exceed unity by about $\sigma$ or greater.
+ Furthermore. the inferred. values for e lor the 2 binning intervals are in agreement. within the errors.," Furthermore, the inferred values for $\epsilon$ for the 2 binning intervals are in agreement, within the errors."
+ Although the results of this analysis do not present strong evidence for anisotropy of the Local Cloud turbulence (our best cases are. after all. shown in Figures 2 and 4). thev are not inconsistent. with 7 pointing in the direction (A=40°&207.550x207). and an anisotropy. parameter e=0.5—0.7.," Although the results of this analysis do not present strong evidence for anisotropy of the Local Cloud turbulence (our best cases are, after all, shown in Figures 2 and 4), they are not inconsistent with $\hat{b}$ pointing in the direction $(\lambda=40^{\circ} \pm 20^{\circ},\beta=50^{\circ} \pm 20^{\circ})$, and an anisotropy parameter $\epsilon = 0.5 - 0.7$."
+ Before leaving tliis section. (wo points should be Muphasized.," Before leaving this section, two points should be emphasized."
+ First. (he modest indications of anisotropy in Table 2 are completely dependent on excluding the measurements of Alkaid and ¢ Dor. which so prominently depart [rom the model curves in Figures 2 and 4.," First, the modest indications of anisotropy in Table 2 are completely dependent on excluding the measurements of Alkaid and $\zeta$ Dor, which so prominently depart from the model curves in Figures 2 and 4."
+ Second. the measurements of £? vs. |cos4| do not adhere loselv to the relationship given by Equation (3). but show a dispersion about Chat curve which is larger (han the measurement error.," Second, the measurements of $\xi^2$ vs. $|\cos A|$ do not adhere closely to the relationship given by Equation (3), but show a dispersion about that curve which is larger than the measurement error."
+ If anisotropy is present in these data. there nist be another. unnamed physical process responsible for variation in £ from one line of sight to another.," If anisotropy is present in these data, there must be another, unnamed physical process responsible for variation in $\xi$ from one line of sight to another."
+ , 
+concentrate on the svstem at the highest redshift. which we will adopt as belonging to the host galaxy of CRB 010222.,"concentrate on the system at the highest redshift, which we will adopt as belonging to the host galaxy of GRB 010222."
+ The host galaxy’s absorption lines anc their identification are compiled in Table 2. together with the rest-frame equivalent width (EW) and some atomic parameters (the oscillator strength.f ancl the ionization potential of the ion. LP. in eV).," The host galaxy's absorption lines and their identification are compiled in Table 2, together with the rest-frame equivalent width (EW) and some atomic parameters (the oscillator strength, and the ionization potential of the ion, IP, in eV)."
+ For comparison. we have listed the EW published by Jha et al. (," For comparison, we have listed the EW published by Jha et al. ("
+2001) and. Masetti et al. (,2001) and Masetti et al. (
+2001).,2001).
+ Phe lines that we could not. identify (neither at redshift 1.4276. nor at the other two redshifts) are Listed in Table 3.," The lines that we could not identify (neither at redshift 1.476, nor at the other two redshifts) are listed in Table 3."
+ lt is important to determine whether the lines are saturated or not. which is cillicult given the low resolution of our spectrum.," It is important to determine whether the lines are saturated or not, which is difficult given the low resolution of our spectrum."
+ Most. of the absorption lines are expected to have a doppler broadening less than a [ew tens of5. which is much less than the spectral resolution (~220 133.," Most of the absorption lines are expected to have a doppler broadening less than a few tens of, which is much less than the spectral resolution $\sim
+220$ )."
+ Therefore. we have compared the observed ratio of the line intensities to the theoretical one. which is eiven by the ratio of their oscillator strengths.," Therefore, we have compared the observed ratio of the line intensities to the theoretical one, which is given by the ratio of their oscillator strengths."
+ Furthermore. the absorption lines with an equivalent. width larger than ~0.5 are most Likely saturated (see for example Pettini ct al.," Furthermore, the absorption lines with an equivalent width larger than 0.5 are most likely saturated (see for example Pettini et al."
+ 2001)., 2001).
+ Xecordinglv. most of the absorption lines detected in the afterglow spectrum of GRB 010222. are actually saluratecd.," Accordingly, most of the absorption lines detected in the afterglow spectrum of GRB 010222 are actually saturated."
+ The strongest. line in the spectrum of GIU. 010222's, The strongest line in the spectrum of GRB 010222's
+The authors would like to thank the auonviuous referee for insightful comments.,The authors would like to thank the anonymous referee for insightful comments.
+ DLE and MCN are grateful for support from the Cordon and Betty Moore Foundation., BLF and MCN are grateful for support from the Gordon and Betty Moore Foundation.
+ GL acknowledges financial support frou NSF Craut AST 07-08819., GL acknowledges financial support from NSF Grant AST 07-08849.
+"reasonable. with SO arguing for c,zz0.05. and most of our computations assume such a value.","reasonable, with SO arguing for $\epsilon_w \approx 0.05$, and most of our computations assume such a value."
+ Llowever. other authors have argued that the power in the mechanical outflow. can substantially exceed Zi (Churazov οἱ 22002. 2005: Peterson Fabian 2006): this situation may be generally the case for raciatively inellicient accretion onto black holes (e.g... Hopkins. Naravan Llernquist 2006).," However, other authors have argued that the power in the mechanical outflow can substantially exceed $L_{\rm bol}$ (Churazov et 2002, 2005; Peterson Fabian 2006); this situation may be generally the case for radiatively inefficient accretion onto black holes (e.g., Hopkins, Narayan Hernquist 2006)."
+ The outlowing wind carries enough momentum to carve out a low density cavity or bubble. but it is also depositing some mass into that bubble.," The outflowing wind carries enough momentum to carve out a low density cavity or bubble, but it is also depositing some mass into that bubble."
+ To provide an estimate of the density of this bubble through which the jets later will be propagating in our scenario. we must introduce another parameter. which can be taken to be the wind speed. co.," To provide an estimate of the density of this bubble through which the jets later will be propagating in our scenario, we must introduce another parameter, which can be taken to be the wind speed, $v_w$."
+" Then we have an expression for the rate at which mass is ejected in the non-relativistic wind A, with eyοςkms+) and τι—nra/10vr."," Then we have an expression for the rate at which mass is ejected in the non-relativistic wind $\dot{M}_w$ with $v_{w,4} = v_w/(10^4 {\rm km ~s}^{-1})$ and $\tau_8 = \tau_{\rm act}/10^8 {\rm yr}$."
+ A plausible. but. by. no means rigid. constraint on. the rate at which mass outllows in this wind is that it should be less than the mass accretion rate needed: to power the AGN.," A plausible, but by no means rigid, constraint on the rate at which mass outflows in this wind is that it should be less than the mass accretion rate needed to power the AGN."
+ Using the typical parametrization for an standard. (raciatively ellicient) accretion disk. where Lince Meet. where Cree0.1. we find that if we adopt the above constraint Thus the minimum wind speed for a raciatively ellicient accretion [low ds. PeauCdO'kms[2 where we have evaluated ((7) taking e;=0.01 and εως=0.057. the latter being the minimum value for thin disk accretion onto a non-rotating black hole (c.g. Shakura Sunvaev 1973).," Using the typical parametrization for an standard (radiatively efficient) accretion disk, where $L_{\rm bol} = \epsilon_{\rm acc} \dot{M}_{\rm acc} c^2$ , where $\epsilon_{\rm acc} \sim 0.1$, we find that if we adopt the above constraint Thus the minimum wind speed for a radiatively efficient accretion flow is, $v_{w,\rm min} \simeq 10^4 {\rm km ~s}^{-1}$, where we have evaluated (7) taking $\epsilon_{k} = 0.01$ and $\epsilon_{\rm acc} = 0.057$, the latter being the minimum value for thin disk accretion onto a non-rotating black hole (e.g. Shakura Sunyaev 1973)."
+ Lt is. however. important to note that substantially lower wind speeds are allowed if the aceretion rate is sullicienthy low so that a raciativelv inellicient accretion Dow (ee. Naravan Xi 1994) is established.," It is, however, important to note that substantially lower wind speeds are allowed if the accretion rate is sufficiently low so that a radiatively inefficient accretion flow (e.g., Narayan Yi 1994) is established."
+" Iis under these circumstances that L, can execed Lisa.", It is under these circumstances that $L_w$ can exceed $L_{\rm bol}$.
+ Since the total mass ejecteck in the wind up to time [uua ds Mau the mean density in a wincd-filled bubble is where we have used ((4).," Since the total mass ejected in the wind up to time $t < \tau_{\rm act} $ is $\dot{M}_w t$ , the mean density in a wind-filled bubble is where we have used (4)."
+ llowever. once ἐς2Tua. Bs is better approximated by ((1) and no more mass is injected. eiving ομμ and thus So the bubble density declines faster after a time Tact ἐκ ) than it did before that time Cx£ P7) even though the radius expands more slowly for /2Tact.," However, once $t > \tau_{\rm act} $, $R_S$ is better approximated by (1) and no more mass is injected, giving $M_w = \dot{M}_w \tau_{\rm act} = 2E_w / v_w^2$ and thus So the bubble density declines faster after a time $\tau_{\rm act}$ $\propto t^{-6/5}$ ) than it did before that time $\propto t^{-4/5}$ ) even though the radius expands more slowly for $t > \tau_{\rm act}$."
+ Vhis gas density is usually extremely low. but the temperature of this gas is tvpically hot enough to produce some N-ray. emission (e.g. Ixrause 2005).," This gas density is usually extremely low, but the temperature of this gas is typically hot enough to produce some X-ray emission (e.g., Krause 2005)."
+ We now assume that at time /;. measured [rom the onset of the wind outflow. a pair of jets of equal kinetic power is launched: the total power in the jets is £j.," We now assume that at time $t_j$, measured from the onset of the wind outflow, a pair of jets of equal kinetic power is launched; the total power in the jets is $L_j$."
+ We will further assume that £; remains constant throughout the duration of the jets’ ejection. 7;.," We will further assume that $L_j$ remains constant throughout the duration of the jets' ejection, $\tau_j$."
+ Because the jets are both focussed ancl ploughing through a low density medium in the wind-swept bubble. they will propagate outward: substantially faster than would a quasi-spherical wind of similar power.," Because the jets are both focussed and ploughing through a low density medium in the wind-swept bubble, they will propagate outward substantially faster than would a quasi-spherical wind of similar power."
+ ‘Two-dimensional hydrodynamic simulations of a particular case of a pair of very powerful jets propagating through a supersonic bubble produced by a supernova driven galactic wind have recently been performed by Krause (2005)., Two-dimensional hydrodynamic simulations of a particular case of a pair of very powerful jets propagating through a supersonic bubble produced by a supernova driven galactic wind have recently been performed by Krause (2005).
+ These computations were designed. to explain. the clumpy HI absorption seen in many high-z RCs in terms of instabilities produced. in the wind shell upon penetration by the jet., These computations were designed to explain the clumpy HI absorption seen in many $z$ RGs in terms of instabilities produced in the wind shell upon penetration by the jet.
+ They clearly show that a powerful jet launched a long time (ἐν= NO Myr in this case) after the wind. initiation will. as we assume. produce a narrow cocoon within the bubble.," They clearly show that a powerful jet launched a long time $t_j =$ 80 Myr in this case) after the wind initiation will, as we assume, produce a narrow cocoon within the bubble."
+ Then. after. passing though the bubble wall. the jet. will continue to propagate more slowlv in a denser ambient medium. but will still produce extended overpressured lobes (Ixrause 2005).," Then, after passing though the bubble wall, the jet will continue to propagate more slowly in a denser ambient medium, but will still produce extended overpressured lobes (Krause 2005)."
+ These simulations neglected. magnetic field in the bubble. whereas we expect the bubble to become magnetized. even if it was not originally so. due to some mixing with the svnchrotron plasma of the radio lobes blown by the jets during their passage through the bubble.," These simulations neglected magnetic field in the bubble, whereas we expect the bubble to become magnetized, even if it was not originally so, due to some mixing with the synchrotron plasma of the radio lobes blown by the jets during their passage through the bubble."
+ Any magnetization should enhance the bubble’s stability against surface distortions (c.g. Ixaiser et 22005: Lyutikoy 2006).," Any magnetization should enhance the bubble's stability against surface distortions (e.g., Kaiser et 2005; Lyutikov 2006)."
+ Following the basic arguments of Ixaiser Alexander (1997) and calibrating the models through catalogs of powerful radio galaxies (Blundell. Rawlings Willott 1999: Darai Wiita 2006. 2007) we have a lower bound on the distance out to which a jet of average collimation will have propagated within the bubble. at time /2/;j. as long as we ignore the spatial gradient in density within the bubble. which is an approximation adequate for our purposes.," Following the basic arguments of Kaiser Alexander (1997) and calibrating the models through catalogs of powerful radio galaxies (Blundell, Rawlings Willott 1999; Barai Wiita 2006, 2007) we have a lower bound on the distance out to which a jet of average collimation will have propagated within the bubble, at time $t > t_j$, as long as we ignore the spatial gradient in density within the bubble, which is an approximation adequate for our purposes."
+" Phe above expression. could be an equality i py did not continue to decline after /;: however. if the velocity at which the jet moves is sulliciently high compared to the rate of expansion of the bubble after the launch of the jet (as usually will be the case in the situations of interest to us). then this continuing temporal decline of p, has only a small effect. which we shall ignore."," The above expression could be an equality if $\rho_b$ did not continue to decline after $t_j$; however, if the velocity at which the jet moves is sufficiently high compared to the rate of expansion of the bubble after the launch of the jet (as usually will be the case in the situations of interest to us), then this continuing temporal decline of $\rho_b$ has only a small effect, which we shall ignore."
+ Hence the numerical results are computed assuming ((10) is an equality., Hence the numerical results are computed assuming (10) is an equality.
+" Only iff;Tact and Ad, is large do we have such high values of ο that the jets will have inordinate cilliculty in catching up to the wind under the approximation we emplov: under these circumstances the continued. decline in. p, with time really should. be considered.", Only if $t_j \ll \tau_{\rm act}$ and $\dot{M_w}$ is large do we have such high values of $\rho_b(t_j)$ that the jets will have inordinate difficulty in catching up to the wind under the approximation we employ; under these circumstances the continued decline in $\rho_b$ with time really should be considered.
+ La the following exploratory calculation we also ignore the slowing down the jet. would suller. while penetrating the reverse shock and the denser shell of material swept up by the expanding bubble., In the following exploratory calculation we also ignore the slowing down the jet would suffer while penetrating the reverse shock and the denser shell of material swept up by the expanding bubble.
+ ba addition. we ignore the modest acceleration produced in the boundary of the bubble through the extra pressure injected by the jet.," In addition, we ignore the modest acceleration produced in the boundary of the bubble through the extra pressure injected by the jet."
+ We note that this slowing of the jet's relative speed. with respect to that of the wind arising from these ellects is roughly olfset by the relative acceleration produced by the continuing temporal decline of the bubble density., We note that this slowing of the jet's relative speed with respect to that of the wind arising from these effects is roughly offset by the relative acceleration produced by the continuing temporal decline of the bubble density.
+ In 1l we clisplay the expansion. of both the wind-blown bubbleand the jets launched. later for four representative combinations of parameters., In 1 we display the expansion of both the wind-blown bubbleand the jets launched later for four representative combinations of parameters.
+ The decline of the bubble expansion. rate is seen at 100. Myr., The decline of the bubble expansion rate is seen at 100 Myr.
+ In. one, In one
+is a low-huumositv (Mpli. lag) E2 ealaxy belonging to the Virgo cluscr.,is a low-luminosity $_B = -17.77$ mag) E2 galaxy belonging to the Virgo cluster.
+ Its catalogue properties are sununuariZe in Table 1.., Its catalogue properties are summarized in Table \ref{tab:properties}.
+ I is one of he four galaxies projected wit ina few arciu of the coutral cD. (?)..," It is one of the four galaxies projected within a few arcmin of the central cD, \citep{prugniel1987}."
+ It haboiws a spectacular clear disk of stars and dust (7).., It habours a spectacular nuclear disk of stars and dust \citep{kormendy2005}.
+ This alinost edge-on disk. deteced up to a radius of 250-350 x (3 ...o»proxinatelv half the effective radius. )). Is marginalv duer than he surroundis population. sugecsting that it is at leas 2 Gyr vounger.," This almost edge-on disk, detected up to a radius of 250-350 pc (3–4, approximately half the effective radius, ), is marginally bluer than the surrounding population, suggesting that it is at least 2 Gyr younger."
+ It is aliguec aloug the major axis of he ealaxy., It is aligned along the major axis of the galaxy.
+ Tus disk ds heieved to have beei formed frou accreted eas that fieed down to the central region., This disk is believed to have been formed from accreted gas that funneled down to the central region.
+ Althoteh it is à bright ealaxy. it is not often observed )ecatise of a foreerouud xielt star located 2.5 away froli he iucleus C223...," Although it is a bright galaxy, it is not often observed because of a foreground bright star located 2.5 away from the nucleus \citep{rc1,devauc1959}."
+ How«vor. fus particular circumstance nakes this oh.ject a first choice tarect for adaptive optics observatious.," However, this particular circumstance makes this object a first choice target for adaptive optics observations."
+ ? used he adaptive optics svsteu of the ΕΠ teescope to reach a spatial resolution of Lor (FWIIM) after a Lucy deconvolution., \citet{kormendy2005} used the adaptive optics system of the CFH telescope to reach a spatial resolution of 0.07 (FWHM) after a Lucy deconvolution.
+ ? use he jear-iufrared inteera field spectrograph SINFONI study the central kinenatics., \citet{nowak2007} used the near-infrared integral field spectrograph SINFONI to study the central kinematics.
+" Using the Sclawarzschi orbit superposition nmerod o fit the two-dimensiona oeformaion. they found a central superauassive black liole of mass Af,=1.25ς Εν and rejected anv mode without a black hole at re [1.5 0 confidence level."," Using the Schwarzschild orbit superposition method to fit the two-dimensional information, they found a central super-massive black hole of mass $
+M_\bullet = 1.25 \times 10^7$ $_\odot$ and rejected any model without a black hole at the 4.5 $\sigma$ confidence level."
+ Tn tUs paper. we stdv the iuterual kinematics ai the stelar population «X this galaxy to investigate if its nuclear particularity is related to its global properties.," In this paper, we study the internal kinematics and the stellar population of this galaxy to investigate if its nuclear particularity is related to its global properties."
+ The paper is organized as follows: in Sect., The paper is organized as follows: in Sect.
+ 2 we prescut the observations aud the data reduction. in Sect.," 2 we present the observations and the data reduction, in Sect."
+ 3 the analysis. aud in Sect.," 3 the analysis, and in Sect."
+ Lowe discuss the results aud conclude., 4 we discuss the results and conclude.
+ Iu the frame of another project (?).. we discovered observations of LLSGA in the Ceili Science Archive.," In the frame of another project \citep{koleva2011}, we discovered observations of 4486A in the Gemini Science Archive."
+ They were taken with the GAIOS spectrograph attacred to the S.1 in CemimiSouth telescope in loue- mode., They were taken with the GMOS spectrograph attached to the 8.1 m Gemini-South telescope in long-slit mode.
+ The setup ane journal of these observations are reported in Table 2.., The setup and journal of these observations are reported in Table \ref{tab:setup}.
+" Two sliehtlv different erating oricutations were used to patch the holes iithe wavelength coverage, Which are cause by the separaion between the three CCD detectors."," Two slightly different grating orientations were used to patch the holes in the wavelength coverage, which are caused by the separation between the three CCD detectors."
+ T1C οserver probably ignored the presence of the star and inadvertently ceu1ος the slit on it. nufortiately σα]ig fre centre of the ealaxy by about 1 (Fig. 1)).," The observer probably ignored the presence of the star and inadvertently centred the slit on it, unfortunately missing the centre of the galaxy by about 1 (Fig. \ref{fig:slit}) )."
+ Nevertheless. these spectre| provide a claico fo proο the stellar population aud the lijeniaties of this object.," Nevertheless, these spectra provide a chance to probe the stellar population and the kinematics of this object."
+ The data reduction was nde using the €(NOS pipeline iu the IRAF euvironment exactly as described iu 7.., The data reduction was made using the GMOS pipeline in the IRAF environment exactly as described in \citet{koleva2011}.
+ The instrumental broacleing. or linc-spread. function (LSF). was found to be acceptably modelled with a Gaussian with saudard deviation Jin.61|.," The instrumental broadening, or line-spread function (LSF), was found to be acceptably modelled with a Gaussian with standard deviation $\sigma_{\rm ins} = 61$."
+ Some pixels within r=0(67 from t1e peak of lieit are saturated., Some pixels within $r = 0.6$ from the peak of light are saturated.
+ The secing measured oninages taken curing the sane observiug bocks is for the two mights 0.9 and 1.6 FWIIM wit La typical error of £0.06.., The seeing measured on images taken during the same observing blocks is for the two nights 0.9 and 1.6 FWHM with a typical error of $\pm 0.06$.
+". The radial positiol. F. of a spectrin with respect to the galaxys kinenatical ceutre is related to the positiou aloie the sli with res])oct o the peak of light... asi LOT?qp(028c0: ""7."," The radial position, $r$, of a spectrum with respect to the galaxy's kinematical centre is related to the position along the slit with respect to the peak of light, $x$, as: $r^2 = (x+2.15\pm0.07)^2 + (1.28\pm0.07)^2$ $^2$."
+ Iu order ο leasure the distance of the galaxy nucleus orthogoua o the slit. ali ACS image of LLISGÀ obtained iu the FS5OLP filter was extracted from the Thible Legaey Archive.," In order to measure the distance of the galaxy nucleus orthogonal to the slit, an ACS image of 4486A obtained in the F850LP filter was extracted from the Hubble Legacy Archive."
+" The distance between the star alc the nucleus was deternuned to be 2,5140.05 frou nieasuriug the position ofthe respective peaks of iuteusity.", The distance between the star and the nucleus was determined to be $2.51\pm0.05$ from measuring the position of the respective peaks of intensity.
+ According to this measurement. the light peal of the ealaxy in the slit should be vat 2.36 from the star.," According to this measurement, the light peak of the galaxy in the slit should be at $2.36$ from the star."
+ Fittine two Cassiaus to the ποτ distribution of the ealaxy, Fitting two Gaussians to the light distribution of the galaxy
+large: ~ 2 pe 27 millicarcsec) in and it terminates at the much smaller ~ 14 x 7 pe (200 pravesec) diameter disk.,large; ${\sim}$ 2 pc ( 27 milli-arcsec) in and it terminates at the much smaller ${\sim}$ 14 ${\times}$ ${^{-3}}$ pc (200 ${\mu}$ arcsec) diameter disk.
+ The physical size inferred for the inflow and the disk components causes NGC 4203 (o nol conform to the correlation between DLE size and UV luminosity established for quasars and high Iuminosity AGNs using reverberation mapping (Peterson2001. 2005).," The physical size inferred for the inflow and the disk components causes NGC 4203 to not conform to the correlation between BLR size and UV luminosity established for quasars and high luminosity AGNs using reverberation mapping \citep{Pet01,Pet93, Kas05}."
+. An extrapolation of the BLR size - luminosity of INaspietal.(2005) down to the low UV luminosity estimated for the AGN in NGC 4203. predicts a radius for the DLB. that is ~ 0.5 Ld. which is  20 times smaller than inferred for the inner radius of (he disk using line profile fitting (Section 3).," An extrapolation of the BLR size - luminosity of \cite{Kas05} down to the low UV luminosity estimated for the AGN in NGC 4203, predicts a radius for the BLR that is ${\sim}$ 0.5 l.d, which is ${\sim}$ 20 times smaller than inferred for the inner radius of the disk using line profile fitting (Section 3)."
+ Conversely. the large outer radius determined for the inflow in NGC 4203 using profile filling makes it physically larger than anv DLR measured using reverberation mapping. 4 times larger than (he previous record holder: the quasar 3C 273 (Ixaspietal.2005).," Conversely, the large outer radius determined for the inflow in NGC 4203 using profile fitting makes it physically larger than any BLR measured using reverberation mapping, 4 times larger than the previous record holder; the quasar 3C 273 \citep{Kas05}."
+". However. as noted by INaspietal.(2005).. the DLR size luminosity correlation appears to break down for low luminosity AGNs. which. with L(2500 A) = 4.4 x 10!"" erg/s (Maoz.2007).. would include NGC 4203."," However, as noted by \cite{Kas05}, the BLR size – luminosity correlation appears to break down for low luminosity AGNs, which, with L(2500 ${\textrm \AA}$ ) = 4.4 ${\times}$ ${^{40}}$ erg/s \citep{Mao07}, would include NGC 4203."
+ Of course. the BLR size luminosity relationship of IXaspietal.(2005) is defined by quasars ancl AGNs that are orders of magnitude more luminous (han NGC 4203.," Of course, the BLR size – luminosity relationship of \cite{Kas05} is defined by quasars and AGNs that are orders of magnitude more luminous than NGC 4203."
+ Thus. the very large discrepancy arising [rom the comparison strongly suggests that the BLAIR in NGC! 4203 is not simply that of a scaled down quasar.," Thus, the very large discrepancy arising from the comparison strongly suggests that the BLR in NGC 4203 is not simply that of a scaled down quasar."
+ Estimating the mass of the BIT in NGC 4203 using the so called ‘virial method’ leads to a value that is substantially lower than the mass estimated from the stellar velocity dispersion bv Lewis&Eracleous(2006)., Estimating the mass of the BH in NGC 4203 using the so called `virial method' leads to a value that is substantially lower than the mass estimated from the stellar velocity dispersion by \cite{Lew06}.
+. For example. the formalism of Greene&IIo.(2005).. which uses the EWIIMand luminosity of the broadIIo. emission line. underestimates the mass of the DII in NGC 4203 by a factor of 348 using the broad La emission line observed in 1999 and a [actor of 51 using the one observed in 2010.," For example, the formalism of \cite{Gre05}, which uses the FWHM luminosity of the broad${\alpha}$ emission line, underestimates the mass of the BH in NGC 4203 by a factor of 348 using the broad ${\alpha}$ emission line observed in 1999 and a factor of 51 using the one observed in 2010."
+ This clichotomy is regarded. as further evidence that the BLR in NGC 4203 is very different from the BLR in more huninous AGNs., This dichotomy is regarded as further evidence that the BLR in NGC 4203 is very different from the BLR in more luminous AGNs.
+ Following Maoz (2007).. the ionizing continuum generated by the AGN in NGC 4203," Following \cite{Mao07}, , the ionizing continuum generated by the AGN in NGC 4203"
+of stellar radius.,of stellar radius.
+ The scatter in the ratio of (Av) computed over different intervals symmetric around vmax decreases with increasing stellar radii and that the ratios converge to 1.0., The scatter in the ratio of $\meandnu$ computed over different intervals symmetric around $\nu_{\rm max}$ decreases with increasing stellar radii and that the ratios converge to 1.0.
+ From the ratios of (Av) we find that the scatter at low radii (<8 Ro) reduces for masses below roughly 1.5 Mo., From the ratios of $\meandnu$ we find that the scatter at low radii $\lesssim 8$ $_{\odot}$ ) reduces for masses below roughly 1.5 $_{\odot}$.
+" On the other hand no significant change in the scatter is present due to different metallicities or due to the inclusion of diffusion, different atmospheres or mixing-length parameters in the models (see bottom panels of Fig. 3))."," On the other hand no significant change in the scatter is present due to different metallicities or due to the inclusion of diffusion, different atmospheres or mixing-length parameters in the models (see bottom panels of Fig. \ref{ratiofit}) )."
+ Note that the relative change in the values of (Av) for stars with radii S8 Ro as a function of the frequency range is of the same order as the precision as with which (Av) can be determined from current state-of-the-art data., Note that the relative change in the values of $\meandnu$ for stars with radii $\lesssim 8$ $_{\odot}$ as a function of the frequency range is of the same order as the precision as with which $\meandnu$ can be determined from current state-of-the-art data.
+" ? suggested two possible reasons for the decreased influence of the frequency range on (Av) for red giants: 1) the trend in Ar over a typical frequency range is approximately linear, and/or 2) Av changes relatively slowly with frequency."," \citet{hekker2011comp} suggested two possible reasons for the decreased influence of the frequency range on $\meandnu$ for red giants: 1) the trend in $\Delta \nu$ over a typical frequency range is approximately linear, and/or 2) $\Delta \nu$ changes relatively slowly with frequency."
+" Other reasons might be that in red giants fewer modes are observed, and these modes have lower radial orders than in less evolved stars."," Other reasons might be that in red giants fewer modes are observed, and these modes have lower radial orders than in less evolved stars."
+ We test these suggestions by fitting a linear polynomial through the values of Av over a range Vmax+A(Av)scanng and investigate the slope of the fit and the standard deviation of the values around the fit., We test these suggestions by fitting a linear polynomial through the values of $\Delta \nu$ over a range $\nu_{\rm max} \pm 4(\Delta \nu)_{\rm scaling}$ and investigate the slope of the fit and the standard deviation of the values around the fit.
+ These results are shown in Fig., These results are shown in Fig.
+" 4 and show that the slope or linear trend of the variation in Av as a function of frequency for stars with lower Ay, i.e., stars with larger radii, is larger than for stars with higher values of Av, with little sensitivity to the mass, metallicity, diffusion, atmosphere and parameter."," \ref{stddev} and show that the slope or linear trend of the variation in $\Delta \nu$ as a function of frequency for stars with lower $\Delta \nu$, i.e., stars with larger radii, is larger than for stars with higher values of $\Delta \nu$, with little sensitivity to the mass, metallicity, diffusion, atmosphere and parameter."
+" For higher mass models there is some variation visible in the trend, which coincides with the ‘hook’ in the H-R diagram (Fig. 1)),"," For higher mass models there is some variation visible in the trend, which coincides with the `hook' in the H-R diagram (Fig. \ref{hrd}) ),"
+" i.e., the contraction during the turn-off of the main sequence."," i.e., the contraction during the turn-off of the main sequence."
+ The scatter around the fit is however much lower for stars with Av below ~10 uHz than for stars with larger values of Av., The scatter around the fit is however much lower for stars with $\Delta \nu$ below $\sim$ 10 $\mu$ Hz than for stars with larger values of $\Delta \nu$ .
+" Thus we indeed conclude that stars with lower Av (<10 Hz), i.e., larger radii (>8 Ro), have a predominantly linear dependence on frequency, while for stars with larger Av (=10 μΗΖ), Le., smaller radii (S8 Ro), the trend is shallower and the scatter in Av increases."," Thus we indeed conclude that stars with lower $\Delta \nu$ $\lesssim 10~\mu$ Hz), i.e., larger radii $\gtrsim 8$ $_{\odot}$ ), have a predominantly linear dependence on frequency, while for stars with larger $\Delta \nu$ $\gtrsim 10~\mu$ Hz), i.e., smaller radii $\lesssim 8$ $_{\odot}$ ), the trend is shallower and the scatter in $\Delta \nu$ increases."
+" The amount of scatter does not seem to depend critically on the inclusion of diffusion, different atmosphere or mixing-length parameter in the models, while it increases with mass and with decreasing metallicity."," The amount of scatter does not seem to depend critically on the inclusion of diffusion, different atmosphere or mixing-length parameter in the models, while it increases with mass and with decreasing metallicity."
+" Acoustic glitches, i.e., regions of sharp-structure variation in the stellar interior, are known to cause a modulation of Av with frequency."," Acoustic glitches, i.e., regions of sharp-structure variation in the stellar interior, are known to cause a modulation of $\Delta \nu$ with frequency."
+" For the Sun and other main-sequence stars both the Helium second-ionization zone (He II zone) and the base of the convection zone have been shown to contribute to the modulation of the frequencies (e.g.??,andreferencesinthe introduction).."," For the Sun and other main-sequence stars both the Helium second-ionization zone (He II zone) and the base of the convection zone have been shown to contribute to the modulation of the frequencies \citep[e.g.][and references in the introduction]{monteiro2005,verner2006}."
+" Generally, the signal of the base of the convection zone is weaker than the signal from the He II zone, due to its location deeper in the star."," Generally, the signal of the base of the convection zone is weaker than the signal from the He II zone, due to its location deeper in the star."
+" Additionally, ? already showed that for ared giant only the He II zone caused a measurable modulation."," Additionally, \citet{miglio2010} already showed that for ared giant only the He II zone caused a measurable modulation."
+Tn its version in 11 dimensions. supergravity will flud new iuportauce in the late 1990s. iu counection with string theory.,"In its version in 11 dimensions, supergravity will find new importance in the late 1990s, in connection with string theory."
+ Wieh derivative corrections will also reappear. in the low cnerev limit of string theory.," High derivative corrections will also reappear, in the low energy limit of string theory."
+ Tartle aud Wawking [55]. introduce the notion of the “wave fiction of the universe” and the “uo-boundars” boundary condition for the Παππιο inteeral. opening up a new intuition on quautunm eravity and quantum cosimologv.," Hartle and Hawking \cite{hartlehawking} introduce the notion of the “wave function of the universe"" and the “no-boundary"" boundary condition for the Hawking integral, opening up a new intuition on quantum gravity and quantum cosmology."
+ But the Euclidean iutegral does not provide a way of computing eeuuine field theoretical quautities in quantum eravitv better than the Wheeler-DeWitt equation. aud the atinosphere at the middle of the eiglities is again rather eloomy.," But the Euclidean integral does not provide a way of computing genuine field theoretical quantities in quantum gravity better than the Wheeler-DeWitt equation, and the atmosphere at the middle of the eighties is again rather gloomy."
+ Ou the other hand. Jim IHartle [56] develops the idea of a sumi over histories formulation of GR iuto a full fledged extension of quantum mechanics to the general covariant setting.," On the other hand, Jim Hartle \cite{hartle} develops the idea of a sum over histories formulation of GR into a full fledged extension of quantum mechanics to the general covariant setting."
+ The idea will later be developed aud formalized by Chris Isham [57].., The idea will later be developed and formalized by Chris Isham \cite{Isham}.
+ Sorkin introduces his poset approach to quanti gravity. [5L].., Sorkin introduces his poset approach to quantum gravity \cite{sorkin}.
+" Creen and Sclavarz realize that strings might describe ""our universe"" [58}."," Green and Schwarz realize that strings might describe “our universe"" \cite{green}."
+. Excitement starts to build up around stiug theory. in connection with the unexpected anomaly cancellation aud the discovery of the heterotic string 1501.," Excitement starts to build up around string theory, in connection with the unexpected anomaly cancellation and the discovery of the heterotic string \cite{hetero}."
+ The relation between the teu dimensional superstiines theory aud four dimensional low eueregv physics is studied in teris of coupactification on Calaby-Yan manifolds [60]. aud orbifolds., The relation between the ten dimensional superstrings theory and four dimensional low energy physics is studied in terms of compactification on Calaby-Yau manifolds \cite{calabi} and orbifolds.
+ The dynamics of the choice of the vacumu remains unclear. but the compactification leads to Ld chiral niodoels reseiibliug low cucrey plivsics.," The dynamics of the choice of the vacuum remains unclear, but the compactification leads to 4d chiral models resembling low energy physics."
+ Belavin. Polvakov and Zamolodchikoy publish their analysis of couformal field theory |61]..," Belavin, Polyakov and Zamolodchikov publish their analysis of conformal field theory \cite{conformal}."
+ Coroff an Saguotti [62) finally compute the two loop divergences of pure GR. definitely nailing the corpse of pure GR perturbative quanti field theory iuto its coffer: the divergent terni is," Goroff an Sagnotti \cite{sagnotti} finally compute the two loop divergences of pure GR, definitely nailing the corpse of pure GR perturbative quantum field theory into its coffer: the divergent term is"
+and Penz&Micela(2008).. who might have confused the star with a nearby object.,"and \cite{pen08a}, who might have confused the star with a nearby object."
+ Extracted spectra were fit using standard procedures with coronal models of | to 3 temperature components (seee.g..Sanz-Foreadaetal.. 2003b).," Extracted spectra were fit using standard procedures with coronal models of 1 to 3 temperature components \citep[see
+ e.g.,][]{san03a}."
+. The actual model used in the fit has little influence on the calculated X-ray (0.12-2.48 keV or 5-103 A)) flux shown in Table l.., The actual model used in the fit has little influence on the calculated X-ray (0.12–2.48 keV or 5–103 ) flux shown in Table \ref{tabfluxes}.
+ More details on the data reduction and treatment will be given in Sanz-Forcada et al., More details on the data reduction and treatment will be given in Sanz-Forcada et al.
+ 2010 (in prep.)., 2010 (in prep.).
+ Measurements with S/N<3 were considered as upper limits., Measurements with $<$ 3 were considered as upper limits.
+ We complemented the sample with lower spatial resolution ROSAT measurements. excluding detections with low statistics (S/N«3). and further marking as upper limits the objects with suspected X-ray bright companions.," We complemented the sample with lower spatial resolution ROSAT measurements, excluding detections with low statistics $<$ 3), and further marking as upper limits the objects with suspected X-ray bright companions."
+ The sample of 75 exoplanets including XMM-Newton. Chandra. and ROSAT detections. have been compared with the whole exoplanet database (417 objects to date) to check whether our sample Is representative of the known exoplanets.," The sample of 75 exoplanets including XMM-Newton, Chandra, and ROSAT detections, have been compared with the whole exoplanet database (417 objects to date) to check whether our sample is representative of the known exoplanets."
+ We applied the Kolmogorov-Smirnov. test to compare both samples: and they represent the same distributions in mass probability) or period probability) in single variable tests., We applied the Kolmogorov-Smirnov test to compare both samples: and they represent the same distributions in mass probability) or period probability) in single variable tests.
+ The stellar X-ray flux was converted into X-ray luminosity. Lx. using the distances listed in Table 1..," The stellar X-ray flux was converted into X-ray luminosity, $L_{\rm X}$ , using the distances listed in Table \ref{tabfluxes}."
+ Also listed are some physical properties of the hosted exoplanets. collected from the Exoplanets Database (Schneider 1995. http://exoplanet.eu/).," Also listed are some physical properties of the hosted exoplanets, collected from the Exoplanets Database (Schneider 1995, http://exoplanet.eu/)."
+" The X-ray flux received at the orbit of the planet is then given by Fx=Lx(A17. where «, is the semimajor axis."," The X-ray flux received at the orbit of the planet is then given by $F_{\rm X}=L_{\rm X}/(4 \pi a_{\rm p}^2)$, where $a_{\rm p}$ is the semimajor axis."
+ The mass- rate from[/ atmospheric losses with 6=K| produced by X-rays simplifies to Table | includes this ecalewlation for a density of p=l0gecem™., The mass-loss rate from atmospheric losses with $\beta=K=1$ produced by X-rays simplifies to Table \ref{tabfluxes} includes this calculation for a density of $\rho$ $^{-3}$.
+ As reference Jupiter. HD 209458b. and HD 189733b have p=l.24. 0.37 and 0.95 ggeem™ respectively.," As reference Jupiter, HD 209458b, and HD 189733b have $\rho$ =1.24, 0.37 and 0.95 $^{-3}$ respectively."
+" The distribution of Fx against the planet mass. M,sin/. is plotted in Fig. 1.."," The distribution of $F_{\rm X}$ against the planet mass, $M_{\rm p} \sin i$, is plotted in Fig. \ref{masses}. ."
+ There is a separation that seems to be related to mass., There is a separation that seems to be related to mass.
+" We plotted a line that roughly follows this separation: logFx=3-0.5(M,sin£D). with Mj in Jovian masses and Fx in CGS units."," We plotted a line that roughly follows this separation: $\log F_{\rm X}=3 - 0.5 (M_{\rm p} \sin
+i)$, with $M_{\rm p}$ in Jovian masses and $F_{\rm X}$ in CGS units."
+ This line ts not based on any previous assumption or physical law., This line is not based on any previous assumption or physical law.
+ We also include the Solar System planets in the diagram by using current emission of the solar corona. which ranges 26€logLx<27.7 (Orlandoetal..2001).. with the vertical segments indicating the variations over the solar cycle.," We also include the Solar System planets in the diagram by using current emission of the solar corona, which ranges $26\leq\log L_{\rm X}\leq27.7$ \citep{orl01}, with the vertical segments indicating the variations over the solar cycle."
+ In this context we can compare the radiation arriving at the Earth when life first appeared about ~3.5 Gyr ago (seeCnossenetal..2007.andreferencestherein). and at an earlier stage. to see whether coronal erosion could have affected the Earth at that time.," In this context we can compare the radiation arriving at the Earth when life first appeared about $\sim$ 3.5 Gyr ago \citep[see][and references therein]{cno07}
+ and at an earlier stage, to see whether coronal erosion could have affected the Earth at that time."
+ We use two stars considered proxies of the Sun at an early age (Ribasetal..2005).. « Cet (~1 Gyr. logLy= 28.89) and EK Dra (40.1 Gyr. logὃν=30.06).," We use two stars considered proxies of the Sun at an early age \citep{rib05}, $\kappa$ Cet $\sim$ 1 Gyr, $\log L_{\rm X}=28.89$ ) and EK Dra $\sim$ 0.1 Gyr, $\log
+L_{\rm X}=30.06$."
+.. Lines indicating their Fx received at | a.u., Lines indicating their $F_{\rm X}$ received at 1 a.u.
+ are plotted in Fig., are plotted in Fig.
+ 1. to mimic the flux at the Earth's orbit in the past., \ref{masses} to mimic the flux at the Earth's orbit in the past.
+ Since the effects of erosion accumulate over the planet's lifetime. we alculated the integrated X-ray flux that has arrived on the planet orbit between the age of 20 Myr. when most protoplanetary disks would have dissipated. and the present day.," Since the effects of erosion accumulate over the planet's lifetime, we alculated the integrated X-ray flux that has arrived on the planet orbit between the age of 20 Myr, when most protoplanetary disks would have dissipated, and the present day."
+ We need to know the stellar age and the X-rayluminosity evolution with time for each star in the sample., We need to know the stellar age and the X-rayluminosity evolution with time for each star in the sample.
+ We canestimate both following Garcéss et al. (, We canestimate both following Garcéss et al. (
+2010. in prep). who relate the,"2010, in prep), who relate the"
+in Y CVn.,in Y CVn.
+ Barnhaumetal.(1991) report the possible detection of Te in two J-stars. EU And and BAL Gem (star studied here).," \citet{bar91} report the possible detection of Tc in two J-stars, EU And and BM Gem (star studied here)."
+ Their argument is based on the presence of a strong absorption at AGOSS which is partially due to the A6085.22 Te I line., Their argument is based on the presence of a strong absorption at $\lambda 6085$ which is partially due to the $\lambda 6085.22$ Tc I line.
+ From these authors. the fact that the A60835 absorption appears which similar intensity only in those carbon stars where Te has been detected unambiguously using the blue lines. supports the identification of the feature al A6085.22 as Te.," From these authors, the fact that the $\lambda 6085$ absorption appears which similar intensity only in those carbon stars where Tc has been detected unambiguously using the blue lines, supports the identification of the feature at $\lambda 6085.22$ as Tc."
+ Our quantitative analvsis of (he A5924 Te feature in BAL Gem is. however. compatible with a non-detection.," Our quantitative analysis of the $\lambda 5924$ Tc feature in BM Gem is, however, compatible with a non-detection."
+ We note that the A5924 Tc feature is a factor ~2 more intense than the A6085 one (Garstang19581).. therefore the presence of Tc in DM Gem in any measurable amount should have appeared in our analvsis.," We note that the $\lambda 5924$ Tc feature is a factor $\sim 2$ more intense than the $\lambda 6085$ one \citep{gar81}, therefore the presence of Tc in BM Gem in any measurable amount should have appeared in our analysis."
+ Furthermore. the A6085 Tc feature is strongly blended with a Ti I line of moderate intensity (4=1.05 eV. log ef=—1.35) and some CN and C absorptions. which necessarily requires a spectral svnthesis analvsis to confirm the detection.," Furthermore, the $\lambda 6085$ Tc feature is strongly blended with a Ti I line of moderate intensity $\chi=1.05$ eV, log $=-1.35$ ) and some CN and $_2$ absorptions, which necessarily requires a spectral synthesis analysis to confirm the detection."
+" Leaving apart the possible presence of Te in DM Gem for a further and accurate analvsis and the upper limits set for WZ Cas and WX Cvg. possible SC-lype stus. we can conclude that most of J-stars do not show Το,"," Leaving apart the possible presence of Tc in BM Gem for a further and accurate analysis and the upper limits set for WZ Cas and WX Cyg, possible SC-type stars, we can conclude that most of J-stars do not show Tc."
+ The Rb abundance is à monitor of the neutron densitv al which the s-process operates in AGB stars., The Rb abundance is a monitor of the neutron density at which the s-process operates in AGB stars.
+ Therefore the derivation of Rb abundances in these stars is extremely important. specifically the abundance ratios between Rb and its neighbors in the periodic table (Zr. Sr).," Therefore the derivation of Rb abundances in these stars is extremely important, specifically the abundance ratios between Rb and its neighbors in the periodic table (Zr, Sr)."
+ We have used the resonance line at A7800.23 to derive Rh abundances., We have used the resonance line at $\lambda 7800.23$ to derive Rb abundances.
+ The other accessible line αἱ A7947 is much weaker ancl very crowded with CN lines in cool stars.," The other accessible line at $\lambda
+7947$ is much weaker and very crowded with CN lines in cool stars."
+ Nevertheless. the resonance line is also blended. and so Rb abundances have to," Nevertheless, the resonance line is also blended, and so Rb abundances have to"
+Finally we cousider the effect of removing the immer disc reeions. keeping both the dise aud secondary irradiated.,"Finally we consider the effect of removing the inner disc regions, keeping both the disc and secondary irradiated."
+ Apar from iucreasing the delav between the ouset of an outburst in the disc aud accretion outo the white dw. a large inner disc radius has a stabilizing effect ou the disc itself. bv preventing nmside-out outbursts.," Apart from increasing the delay between the onset of an outburst in the disc and accretion onto the white dwarf, a large inner disc radius has a stabilizing effect on the disc itself, by preventing inside-out outbursts."
+ This effect is quite noticeable iu the case where the white dwarf surface temperature is hieh: if the iuner disc radius rg Is large enough. the unstable transition region between the stable reeion heated above hydrogen ionization teniperature aud the cooler. quiescent external part docs not exist.," This effect is quite noticeable in the case where the white dwarf surface temperature is high: if the inner disc radius $r_{\rm in}$ is large enough, the unstable transition region between the stable region heated above hydrogen ionization temperature and the cooler, quiescent external part does not exist."
+ This suppresses the bounces after a longer outburs Jas can be seen in figure 7..," This suppresses the bounces after a longer outburst, as can be seen in figure \ref{fig:rin}."
+" For a eiven mass transfer rate. there is a critical value of the iuner radius above which the dise is stable: wheu ry, approaches this value. the recurrence rate goes to infinity."," For a given mass transfer rate, there is a critical value of the inner radius above which the disc is stable; when $r_{\rm in}$ approaches this value, the recurrence rate goes to infinity."
+ Arbitrarily lone recurrence rates could therefore be expected. but only at the expeuse of very fine tuniue: in the normal case where the mass accretion rate onto the white dwarf is uceligible iu quiescence as compared with the mass trausfer rate from the secondary. the reasoning used by (Siiaxk 1993)) applies.," Arbitrarily long recurrence rates could therefore be expected, but only at the expense of very fine tuning; in the normal case where the mass accretion rate onto the white dwarf is negligible in quiescence as compared with the mass transfer rate from the secondary, the reasoning used by (Smak \cite{s93}) ) applies."
+ The recurrence nne tye IS equal to: where fis the ratio of the amount of mass trausterred during an outburst AAL and the maxima possible disc mass Mog. Obtained assuming that the surface density is everywhere the critical surface density.," The recurrence time $t_{\rm
+rec}$ is equal to: where $f$ is the ratio of the amount of mass transferred during an outburst $\Delta M$ and the maximum possible disc mass $M_{\rm crit}$, obtained assuming that the surface density is everywhere the critical surface density."
+ Numerical mocels show that the surface density is uot very far from its critical value even at large radii. aud that the amount of inass trausterred during a normal outburst is typically of the otal disc mass.," Numerical models show that the surface density is not very far from its critical value even at large radii, and that the amount of mass transferred during a normal outburst is typically of the total disc mass."
+" Therefore. f is not a very σπα] paraineter that could freely vary. and large changes in fj Cannot result from variatious in ry, alouc."," Therefore, $f$ is not a very small parameter that could freely vary, and large changes in $t_{\rm rec}$ cannot result from variations in $r_{\rm in}$ alone."
+ A simular situation is encouutered in the case of soft N-rav transients (Alenou ot al. 1999))., A similar situation is encountered in the case of soft X-ray transients (Menou et al. \cite{mhln99}) ).
+ We have shown that αμα types of lieht curves ca- be produced ly numerical iuodels that include the ilhuumatiou of both the secondary and the accretio- disc. thereby explaining a great variety of observed helt curves.," We have shown that many types of light curves can be produced by numerical models that include the illumination of both the secondary and the accretion disc, thereby explaining a great variety of observed light curves."
+ These effects account for phenomena such as post-outburst rebrighteniug (e.g. Cuc)). long outbursts or example). or SU UMa systems wit[um extremely short supercvcles.," These effects account for phenomena such as post-outburst rebrightening (e.g. ), long outbursts for example), or SU UMa systems with extremely short supercycles."
+ Iu order to explore further these possibilities. one would need to determune from observatious the mass transfer rate from the secondary as a fiction of the mass accretion rate onto the white dwarf. with a better accuracy than it is available now.," In order to explore further these possibilities, one would need to determine from observations the mass transfer rate from the secondary as a function of the mass accretion rate onto the white dwarf, with a better accuracy than it is available now."
+ Despite the fact that our approximations are very crude. in particular the one concerning the response of the secondary to dlunuimation. we cau nevertheless draw a uuuber ofconclusions.," Despite the fact that our approximations are very crude, in particular the one concerning the response of the secondary to illumination, we can nevertheless draw a number of conclusions."
+ First. the illunination of the disc is inportaut ouly if the white chwart is relatively massive. so that it can have a high temperature without contributing too much to the light emitted by the system iu outburst. aud that he efficiency of accretion is high.," First, the illumination of the disc is important only if the white dwarf is relatively massive, so that it can have a high temperature without contributing too much to the light emitted by the system in outburst, and that the efficiency of accretion is high."
+ Rebrightcnines also require à not to be too low in quiescence., Rebrightenings also require $\alpha$ not to be too low in quiescence.
+ The fact that we can reproduce au alternance of normal aud loug outbursts when the ilhuuimation of the secondary is meluded does not of course Imply that the thermaltical iustability model for SU UMa is incorrect: a tidal instability is uost probably required to account for the superhump phenomenon., The fact that we can reproduce an alternance of normal and long outbursts when the illumination of the secondary is included does not of course imply that the thermal-tidal instability model for SU UMa is incorrect; a tidal instability is most probably required to account for the superhump phenomenon.
+ The questio- of the precise role of this iustabilitv is. however. still open aud our results raise ποιο doubt on the validity of the parameters derived when fittine he observations. in particular for svstenis wving a very short supercyele.," The question of the precise role of this instability is, however, still open and our results raise some doubt on the validity of the parameters derived when fitting the observations, in particular for systems having a very short supercycle."
+ This also means that the deteruinatiou of the viscosity roni the modeling of Πο curves is a far more difficult ask than previously estimated., This also means that the determination of the viscosity from the modeling of light curves is a far more difficult task than previously estimated.
+ We obviously need some xogress iu the determination of the tidal torque: we also weed to know how the secondary responds to illunination., We obviously need some progress in the determination of the tidal torque; we also need to know how the secondary responds to illumination.
+" We finally should include 2D effects iu our models,", We finally should include 2D effects in our models.
+ First jcause the orbits iu the outer disc ave far from being circular. aud secoud because the presence of a lot spot whose temperature can be of order of 10.000 I& could iu xinciple significantly alter the stability properties of the outer disc.," First because the orbits in the outer disc are far from being circular, and second because the presence of a hot spot whose temperature can be of order of 10,000 K could in principle significantly alter the stability properties of the outer disc."
+exposure times and S/N near 6700 are listed for each observing run separately.,exposure times and S/N near 6700 are listed for each observing run separately.
+ For MSO. the star uumber and photometry are ourown?.. with E(B—V)=0.17 magnitudes assumed Guetter1976:Reed.Hesser.&Shawl 1983).," For M80, the star number and photometry are our, with $B~-~V)~=~0.17$ magnitudes assumed \citep{KG76,RHS88}."
+. Despite the rather large list of stars that we observe. ouly 21 mace it past the cuttiug room floor.," Despite the rather large list of stars that we observed, only 21 made it past the cutting room floor."
+ The main obstacle to obtaining more spectra was the need for both high S/N aud high resolution. which was aggravated by varying observing conditions and intermittent mechanical difficulties during the comuissioning of Hydra.," The main obstacle to obtaining more spectra was the need for both high S/N and high resolution, which was aggravated by varying observing conditions and intermittent mechanical difficulties during the commissioning of Hydra."
+ The data from both epochs were processed in similar fashion using the reductionroutines i IRAF (Tody1986).. beginning with the usual overscan aud bias corrections.," The data from both epochs were processed in similar fashion using the reductionroutines in IRAF \citep{IRAF}, beginning with the usual overscan and bias corrections."
+ We mace our Lat-liel« image by combining several daytime sky flats that were exposed by usiug diffusing filter in place ol the echelle filter. removing the spectral shape in the x aucl ν directions by curve-fittiug. aux ioxcar filtering the remaiuder. which lef s with a smooth milky flat that contained only variatious.," We made our flat-field image by combining several daytime sky flats that were exposed by using diffusing filter in place of the echelle filter, removing the spectral shape in the x and y directions by curve-fitting, and boxcar filtering the remainder, which left us with a smooth milky flat that contained only pixel-to-pixel variations."
+ The data images were then divided by this flat., The data images were then divided by this flat.
+ The object spectra were extractec ining the IRAF task APALL with variance weighting. background subtraction. and cosmic-ray Cleaning parameters turned ou.," The object spectra were extracted using the IRAF task APALL with variance weighting, background subtraction, and cosmic-ray cleaning parameters turned on."
+ The fiber-to-fiber respouse [unctious were ‘aleulatecl using the ISRESPID task. where the individual fiber respouses in an averaged quartz flat were determined with reference to an averaged twilight sky flat.," The fiber-to-fiber response functions were calculated using the MSRESP1D task, where the individual fiber responses in an averaged quartz flat were determined with reference to an averaged twilight sky flat."
+ To remove fiber-to-liber variations as well as most olthe instrument prolile. {οσο respouses were theu divided through the extracted spectra to create uormalized spectra.," To remove fiber-to-fiber variations as well as most of the instrument profile, these responses were then divided through the extracted spectra to create normalized spectra."
+ To determine the dispersion solutious for both epochs. we used etalon exposures bracketed around our object exposures throughout the uight.," To determine the dispersion solutions for both epochs, we used etalon exposures bracketed around our object exposures throughout the night."
+" These wielittime etalon spectra were calibrated with ""master"" etalon spectra that were acquired early during the first uiglit of each ruu aud were themselves calibrated from a Thr lamp for the first epoch. and a HeNeAr lamp for the second epoch."," These nighttime etalon spectra were calibrated with “master” etalon spectra that were acquired early during the first night of each run and were themselves calibrated from a ThAr lamp for the first epoch, and a HeNeAr lamp for the second epoch."
+ The dispersion solution for a eas-lilled lamp was applied to the extracted master etalou spectra. which were then averaged together to [orm oue spectrum.," The dispersion solution for a gas-filled lamp was applied to the extracted master etalon spectra, which were then averaged together to form one spectrum."
+ The individual lines iu the averaged master etalon spectrum were then measured aud used to calibrate the individual spectra [rom the master etalon image so that every fiber used tlie same set of lines in its clispersion solution., The individual lines in the averaged master etalon spectrum were then measured and used to calibrate the individual spectra from the master etalon image so that every fiber used the same set of lines in its dispersion solution.
+ These master etalou spectra were then used to calibrate the nighttiime etalou spectra., These master etalon spectra were then used to calibrate the nighttime etalon spectra.
+ We checked the quality of our etalou solutious by comparing ones taken ou two different niglits aud fouud tliat both fiber-to-liber and night-to-nieht dispersion solutions remained cousistent to less than 0.002A., We checked the quality of our etalon solutions by comparing ones taken on two different nights and found that both fiber-to-fiber and night-to-night dispersion solutions remained consistent to less than 0.002.
+ We prefer using the etalon because of its superior precision., We prefer using the etalon because of its superior precision.
+ For example. during the second epoch the HeNeAr exposure only contained 17 usable lines over ~300A. with most lines on the red side of the spectrum. while the etalon coutaiued 95 strong uublended lines uuilormly spaced over the spectrum.," For example, during the second epoch the HeNeAr exposure only contained 17 usable lines over 300, with most lines on the red side of the spectrum, while the etalon contained 95 strong unblended lines uniformly spaced over the spectrum."
+ The typical rms for the HeNeAr bunp was arouud 0.03 A while the etalous were, The typical rms for the HeNeAr lamp was around 0.03 while the etalons were
+sole aspects of the shear patterns while considering the function. which cau be witteu iu terms of the complex coordinate correlations as The £4 patterns are depicted on Fie. X.,"some aspects of the shear patterns while considering the function, which can be written in terms of the complex coordinate correlations as The $\xi_{3}$ patterns are depicted on Fig. \ref{xi3pattern}."
+ Thev are very simular fo what has been observed in umuerical siuulations., They are very similar to what has been observed in numerical simulations.
+" This part of the components does not however encompass the signal coming frou, equilateral configurations.", This part of the components does not however encompass the signal coming from equilateral configurations.
+ The extension of the uniform and strong pattern in between the two points depends again ou the halo profile., The extension of the uniform and strong pattern in between the two points depends again on the halo profile.
+ The use of complex coordinates proved fruitful in conrputations of shear correlation fictions., The use of complex coordinates proved fruitful in computations of shear correlation functions.
+ It has been possible to obtain the shear three-poimt correlation functions i manageable forms for the onc-halo model with a power law profile., It has been possible to obtain the shear three-point correlation functions in manageable forms for the one-halo model with a power law profile.
+ In the context of this model those results can be used to imfer general properties for tle shear correlation patterus. as in Denabed Scoccimarro (2005): a firs step toward the construction of robust methods for the measurement of the shear three-point fiction.," In the context of this model those results can be used to infer general properties for the shear correlation patterns, as in Benabed Scoccimarro (2005); a first step toward the construction of robust methods for the measurement of the shear three-point function."
+ The theoretical framework preseuted in this paper is actually very general and can be used for alternative description of the density correlation properties as such those used iu the quasilinear or the noulinear reeiue., The theoretical framework presented in this paper is actually very general and can be used for alternative description of the density correlation properties as such those used in the quasilinear or the nonlinear regime.
+ Tudeed the mathematical relation between the projected potential correlation fictions aud those of the projected density is rather simple aud could be imteerated out for suuple openiug the way for explicit computations of the shear correlation (fictions m more diverse cases., Indeed the mathematical relation between the projected potential correlation functions and those of the projected density is rather simple and could be integrated out for simple opening the way for explicit computations of the shear correlation functions in more diverse cases.
+ 2pc The caleulatious presented in this paper would not xi3-norm2-eps-converted-topat been Possible without the formula initially excavated by Roma Scoccinuuro., 2pc The calculations presented in this paper would not have been possible without the formula initially excavated by Románn Scoccimarro.
+an extrapolation of the accepted best-fit laxge-INDO spectrum at the time (?)..,an extrapolation of the accepted best-fit large-KBO spectrum at the time \citep{trujillo01}.
+ While this observed decrement of more than an order of magnitude in the number of small NMBOs clearly indicates a break between 45 and 100 kin. (he exact break position aud slope below the break may well be refined by future data on small INBOs.," While this observed decrement of more than an order of magnitude in the number of small KBOs clearly indicates a break between 45 and 100 km, the exact break position and slope below the break may well be refined by future data on small KBOs."
+ Still. the results of ? are inconsistent wilh the previously expected small-end spectrum Nox27°. org =3.5. al better than confidence.," Still, the results of \cite{bernstein03} are inconsistent with the previously expected small-end spectrum $N\propto R^{-2.5}$, or $q=3.5$, at better than confidence."
+ This paper describes a simple sell-consistent analvtic calculation of the break location and the slope below the break., This paper describes a simple self-consistent analytic calculation of the break location and the slope below the break.
+ Note that using the V(r)xr! size spectrum obtained by for large KBOs. we can estimate (he size below which collisions between equal size boclies should be [frequent to be ~ 1 kmwell below the observed. break location.," Note that using the $N(r)\propto r^{-4}$ size spectrum obtained by \cite{bernstein03} for large KBOs, we can estimate the size below which collisions between equal size bodies should be frequent to be $\sim$ 1 km—well below the observed break location."
+ HIowever. this estimate needs two modifications.," However, this estimate needs two modifications."
+ First. due to the large velocity dispersion in the Nuiper belt. small bodies can shatter much larger objects.," First, due to the large velocity dispersion in the Kuiper belt, small bodies can shatter much larger objects."
+ Since there are more small than large bodies. destructive collisions will occur frequently even for objects much lareer than 1 kin.," Since there are more small than large bodies, destructive collisions will occur frequently even for objects much larger than 1 km."
+ second. when collisions are important. thev reduce the nunber of small bodies: this in tum decreases the frequency. of collisions.," Second, when collisions are important, they reduce the number of small bodies; this in turn decreases the frequency of collisions."
+ Therefore. caleulations of the effects of collisions and the size below which collisions are important must be done in a sell-consistent manner.," Therefore, calculations of the effects of collisions and the size below which collisions are important must be done in a self-consistent manner."
+ In order to [find the break location sell-consistentlyv. we first calculate the power-law slope q For a collisional population of bodies.," In order to find the break location self-consistently, we first calculate the power-law slope $q$ for a collisional population of bodies."
+ We assume a group of bodies with isotropic velocily dispersion e in which the cillerential number of bodies of radius r is given bv a power law dN(r)/drzxr.© , We assume a group of bodies with isotropic velocity dispersion $v$ in which the differential number of bodies of radius $r$ is given by a power law $dN(r)/dr \propto r^{-q}$.
+If we assume that the population is in a steady state and (hat nass is conserved in the collision process. the total mass of bodies destroved per unit time in a logarithmic interval in radius must be independent of size.," If we assume that the population is in a steady state and that mass is conserved in the collision process, the total mass of bodies destroyed per unit time in a logarithmic interval in radius must be independent of size."
+ We can use this condition to determine q., We can use this condition to determine $q$.
+ We assume that the main channel for mass destruction is the shattering of larger targets by smaller ‘bullets’ (Fig. 1))., We assume that the main channel for mass destruction is the shattering of larger `targets' by smaller `bullets' (Fig. \ref{cascade}) ).
+ Under this condition. Ilere p is the internal density. of each body and πμ} is the size of the smallest bullet which. on impact. can shatter a target of radius r. V is the volume occupied by all the bodies: their velocity dispersion and therefore their distribution within V. are assumed independent ol body size.," Under this condition, Here $\rho$ is the internal density of each body and $r_B(r)$ is the size of the smallest bullet which, on impact, can shatter a target of radius $r$ $V$ is the volume occupied by all the bodies; their velocity dispersion and therefore their distribution within $V$ are assumed independent of body size."
+ When supplemented by a relation between (he size of the bullet and that of the target. Eq.," When supplemented by a relation between the size of the bullet and that of the target, Eq."
+1. dictates the size spectiruni q.,\ref{steadystate} dictates the size spectrum $q$ .
+to be less good than the best available planet search caudiclates. so hat in each case stars closer and/or younger than the example actually exist.,"to be less good than the best available planet search candidates, so that in each case stars closer and/or younger than the example actually exist."
+ Using the very votnugest stars would also have resulted in sensitivities better than 1Mjqy.. a nass regime uot covered by the moclels used in the Figures.," Using the very youngest stars would also have resulted in sensitivities better than 1, a mass regime not covered by the \citet{bur} models used in the Figures."
+ Figures 5. ane ιο illustrate three very iiportant polnts., Figures \ref{fig:HLM1} and \ref{fig:HLM2} illustrate three very important points.
+ First. tle L' band appears to have ouly secoudary usefuluess siuce either the A baud or tie A bane alwavs olfers sensitivity to Ower-miass planets.," First, the $L'$ band appears to have only secondary usefulness since either the $H$ band or the $M$ band always offers sensitivity to lower-mass planets."
+ Secoud. Figure 6 shows tlat with a reatively minor ilcrease of 1 maguitude iu sensitivity. the AL band will be sensitive to lower-miass platels around a[st ars withiu 5 pe than cau ye detected with H band observations. even i‘the A band sensitivity luc‘eases tlie same amount.," Second, Figure \ref{fig:HLM2} shows that with a relatively minor increase of 1 magnitude in sensitivity, the $M$ band will be sensitive to lower-mass planets around all stars within 5 pc than can be detected with $H$ band observations, even if the $H$ band sensitivity increases the same amount."
+" ""Third. Figure 5 shows that the advautage of the AZ banc decreases with iucreasing distance. but hat as larger telescopes aud lounger exposures increase sensitivities to 2.5 1jag above present levels. he AL band will be superio ‘to H out to 10 pe."," Third, Figure \ref{fig:HLM1}
+ shows that the advantage of the $M$ band decreases with increasing distance, but that as larger telescopes and longer exposures increase sensitivities to 2.5 mag above present levels, the $M$ band will be superior to $H$ out to 10 pc."
+ With au increase of [1 wae. the AZ band. would surpass 4 out to 25 pe — bt tas such a large sensitivity increase would be difficult to achieve. A uid will likely reiiain the primary wavelength [or stars at 25 pe aud beyoud.," With an increase of 4 mag, the $M$ band would surpass $H$ out to 25 pc – but as such a large sensitivity increase would be difficult to achieve, $H$ band will likely remain the primary wavelength for stars at 25 pc and beyond."
+ For stars closer thau 10 pe. however. the AM bane already. ollers excellent seusitivity that has barely been exploited so ar.," For stars closer than 10 pc, however, the $M$ band already offers excellent sensitivity that has barely been exploited so far."
+ Give reasonade sensitivity Increases. AZ should become the primary baud for planet searches around stars at a ¢listauce of 10 pc or less.," Given reasonable sensitivity increases, $M$ should become the primary band for planet searches around stars at a distance of 10 pc or less."
+ Interestingly. he conclusious of Figures 5 αμα 6 are essentially independent of age: extensive calculations by Heinze(2007 showed that the relative usefuluess of different wavelengths πας only a weak depeudeuce on age. for stars at a fixed distance — aud even this weak age dependence could Change sign ¢- swlJ.C[uning [rom the iuodels of Burrowsetal.(2003) o those of Baralleetal. (2003).," Interestingly, the conclusions of Figures \ref{fig:HLM1} and \ref{fig:HLM2} are essentially independent of age: extensive calculations by \citet{thesis} showed that the relative usefulness of different wavelengths had only a weak dependence on age, for stars at a fixed distance – and even this weak age dependence could change sign on switching from the models of \citet{bur} to those of \citet{bar}. ."
+. This means that if we chauee tlie ages of the stars in Figures 5. aud 6 but leave the distances the same. the L/. AL. anc H baud curves will slide up or down but remain essentially fixed in their 'elative positious.," This means that if we change the ages of the stars in Figures \ref{fig:HLM1} and \ref{fig:HLM2}
+ but leave the distances the same, the $L'$, $M$ , and $H$ band curves will slide up or down but remain essentially fixed in their relative positions."
+ For example. given a 3 magnitude increase in sersitivity at both wavelengths. AL xuid observations will detect lower mass planets than A-bauc oues arouud a star at 10 pc. whetler he stellar age is 2 Gyr. 1 Gyr. or 100 Myr.," For example, given a 3 magnitude increase in sensitivity at both wavelengths, $M$ band observations will detect lower mass planets than $H$ -band ones around a star at 10 pc, whether the stellar age is 5 Gyr, 1 Gyr, or 100 Myr."
+ This is to be expeced. since if oue clials down the age of a given liy]nthetical star system. the (and therefore IR color) of the faintest detectable jxanets will remain about the same. though their masses will decrease.," This is to be expected, since if one dials down the age of a given hypothetical star system, the (and therefore IR color) of the faintest detectable planets will remain about the same, though their masses will decrease."
+ Again. Figures 5 and 6 apply only to backgrouud-Iimited sensitivity.," Again, Figures \ref{fig:HLM1} and \ref{fig:HLM2} apply only to background-limited sensitivity."
+ However. given the much nore favorable planet/star flux ratios in the Af band relative to H. we would expect he longer wavelength observations to remain equally competitive closer to the star.," However, given the much more favorable planet/star flux ratios in the $M$ band relative to $H$, we would expect the longer wavelength observations to remain equally competitive closer to the star."
+ Advances it AL band coronograplhiy wil likely parallel the development of H baud extreme AO systems such as GPI and SPHERE., Advances in $M$ band coronography will likely parallel the development of $H$ band extreme AO systems such as GPI and SPHERE.
+ Though at present they are surpassed iu seusitivity by H-regime observations for all but the nearest stars. the L’ andespeciallythe AZ bands hold considerable promise for the future.," Though at present they are surpassed in sensitivity by $H$ -regime observations for all but the nearest stars, the $L'$ andespeciallythe $M$ bands hold considerable promise for the future."
+ We have surveyed unusually nearby. mature star systems for extrasolar planets in the £/ and," We have surveyed unusually nearby, mature star systems for extrasolar planets in the $L'$ and"
+In this case. the data can be lit well with 3 eliteh parameters (epoch. Av. and Ar) and 3 spin parameters (v7. P. ancl 7).,"In this case, the data can be fit well with 3 glitch parameters (epoch, $\Delta{\nu}$, and $\Delta{\dot{\nu}}$ ) and 3 spin parameters $\nu$ , $\dot{\nu}$, and $\ddot{\nu}$ )."
+ Besiduals for this fit are shown in the top panel of Figure 5.., Residuals for this fit are shown in the top panel of Figure \ref{fig:resids_glitch}.
+ The best fit glitch. parameters. as determined with TEMPO. eives a glitch epoch of MJD 51335412. in agreement with the partially phase-coherent analysis.," The best fit glitch parameters, as determined with TEMPO, gives a glitch epoch of MJD $51335 \pm 12$, in agreement with the partially phase-coherent analysis."
+ We found Av/y=(1.420.2)x10.2. which was undetectable in the partially coherent analvsis. ancl 2.50 [rom the value reported by Zhangetal.(2001).," We found $\Delta{\nu}/{\nu} = (1.4\pm0.2) \times 10^{-9}$, which was undetectable in the partially coherent analysis, and $\sigma$ from the value reported by \citet{zmg+01}."
+. We also found 4! in agreement with the partially coherent analvsis. but 106. larger than that reported by Zhangοἱal.(2001). Ni/P=(0.8520.05)10 1.," We also found $\Delta{\dot{\nu}}/{\dot{\nu}} = (1.33 \pm 0.02) \times 10^{-4}$ , in agreement with the partially coherent analysis, but $\sigma$ larger than that reported by \citet{zmg+01}, $\Delta{\dot{\nu}}/{\dot{\nu}}
+= (0.85 \pm 0.05) \times 10^{-4}$ ."
+ To conchisively show that the model with the glitch is a better description of the data. we fitted additional [frequency derivatives until (he same number of parameters was included in each fit (ie. 12 parameters).," To conclusively show that the model with the glitch is a better description of the data, we fitted additional frequency derivatives until the same number of parameters was included in each fit (i.e. 12 parameters)."
+ The bottom panel of Figure 5. shows the results of this fil., The bottom panel of Figure \ref{fig:resids_glitch} shows the results of this fit.
+ Note that the RAIS residual is a [actor of 5.15 smaller for the model with the glitch. corresponding to a reduced 4? of 1.36 [or 545 degrees of ουσια. compared to a reduced 7? of 49.4 for 545 degreesoO of [reedom for the fit with no eglitch.," Note that the RMS residual is a factor of 5.15 smaller for the model with the glitch, corresponding to a reduced ${\chi}^2$ of 1.86 for 545 degrees of freedom, compared to a reduced ${\chi}^2$ of 49.4 for 545 degrees of freedom for the fit with no glitch."
+ This confirms the conclusion of our partially coherent analvsis. namely that a glitch occured as reported by (2001).," This confirms the conclusion of our partially coherent analysis, namely that a glitch occured as reported by \citet{zmg+01}."
+ The phase-counected solution with the minimum nunber of derivatives (i.e. (he number required to obtain pliase-connection. in this case. 2). and the elitch fitted. gives a value of n—2.10651z0.00001.," The phase-connected solution with the minimum number of derivatives (i.e. the number required to obtain phase-connection, in this case, 2), and the glitch fitted, gives a value of $n=2.10657 \pm 0.00001$."
+ However. since (here is clearly low-frequency noise conteminating the data (top panel of Figure 5)). the formal uncertainty ereatly underestimates the (rue uncertainty on n.," However, since there is clearly low-frequency noise contaminating the data (top panel of Figure \ref{fig:resids_glitch}) ), the formal uncertainty greatly underestimates the true uncertainty on $n$."
+ Spin parameters varv when higher order derivatives are fitted due both (o timing noise and covariances between parameters (Livingstoneetal.2005)., Spin parameters vary when higher order derivatives are fitted due both to timing noise and covariances between parameters \citep{lkgm05}.
+.. Thus. we estimate (he upper limit on the uncertainiv in 2. albeit not rigorously. [rom (he variation in the measured value as higher order derivatives are added to the fit.," Thus, we estimate the upper limit on the uncertainty in $n$, albeit not rigorously, from the variation in the measured value as higher order derivatives are added to the fit."
+ This method implies an uncertainty of 2.11c0.06. in agreement with the value obtained in the partially analvsis.," This method implies an uncertainty of $2.11 \pm 0.06$, in agreement with the value obtained in the partially phase-coherent analysis."
+ All spin and eliteh parameters [rom the fully phase-coherent analvsis are eiven in Table 1., All spin and glitch parameters from the fully phase-coherent analysis are given in Table 1.
+ Uncertainties given for parameters are (he formal uncertainites. except in ihe case of n.," Uncertainties given for parameters are the formal uncertainites, except in the case of $n$."
+ Our initial (imine analvsis was [ist performed using (hie pulsar position determined in a observation (INaaretetal.2001)., Our initial timing analysis was first performed using the pulsar position determined in a observation \citep{kma+01}.
+. The reported lo uncertainty wasoon the position., The reported $\sigma$ uncertainty wason the position.
+ Holding (his position fixed. the post-lit residuals for the fully phase-coherent analvsis showed periodic behaviour. wilh a period of 1 vear and amplitude," Holding this position fixed, the post-fit residuals for the fully phase-coherent analysis showed periodic behaviour, with a period of 1 year and amplitude"
+fluctuations actually diminish with increasing n4.,fluctuations actually diminish with increasing $n_\alpha$.
+" Accordingly, the 67, power spectrum peaks and then decreases towards lower redshift."," Accordingly, the $\delta T_b$ power spectrum peaks and then decreases towards lower redshift."
+" Of course, our assumption that Ταν remains close to its adiabatic value becomes suspect in this regime."," Of course, our assumption that $T_{\rm kin}$ remains close to its adiabatic value becomes suspect in this regime."
+" Eventually, a sufficiently strong Lyman-a intensity will not merely pump the 21-cm transition, but will appreciably heat the gas as well."," Eventually, a sufficiently strong $\alpha$ intensity will not merely pump the 21-cm transition, but will appreciably heat the gas as well."
+" In order to measure the very large scale structure, one needs high brightness sensitivity on BAO scales."," In order to measure the very large scale structure, one needs high brightness sensitivity on BAO scales."
+" At z>10, the angular scale of BAO changes little, and —10' scales are important, just as they are in the CMB."," At $z>10$, the angular scale of BAO changes little, and $\sim 10^\prime$ scales are important, just as they are in the CMB."
+" To achieve high brightness sensitivity, a filled aperture is desirable."," To achieve high brightness sensitivity, a filled aperture is desirable."
+ Telescopes such as Arecibo and FAST would be well suited.," Telescopes such as Arecibo and FAST \citep{FAST}
+ would be well suited."
+" With filled apertures of 300m and 500m respectively, their angular resolution (Smitsfor 21 cm at 2009)z~20 is 40 and 30 arc minutes, respectively."," With filled apertures of 300m and 500m respectively, their angular resolution for 21 cm at $z\sim 20$ is 40 and 30 arc minutes, respectively."
+" We focus our attention on FAST, which is expected to observe at sufficiently long wavelengths."," We focus our attention on FAST, which is expected to observe at sufficiently long wavelengths."
+" At these low frequencies, sensitivities are sky limited, with T,,,~ 3000K. It is also straightforward to observe with a focal plane array, so we reference our forecasts to a 100 pixel array, which would be a 40m dipole array."," At these low frequencies, sensitivities are sky limited, with $T_{\rm sys} \sim 3000$ K. It is also straightforward to observe with a focal plane array, so we reference our forecasts to a 100 pixel array, which would be a 40m dipole array."
+" Such focal plane array would allow primary beam illumination to compensate with frequency to make frequency independenta beams, enabling accurate foreground subtraction."," Such a focal plane array would allow primary beam illumination to compensate with frequency to make frequency independent beams, enabling accurate foreground subtraction."
+" The maximum transverse k,=0.08.", The maximum transverse $k_\perp=0.08$.
+" We further assume the sky is drift scanned, which minimizes systematic errors, and makes this experiment comparable to the recent GBT 21cm intensity detection (Changetal."," We further assume the sky is drift scanned, which minimizes systematic errors, and makes this experiment comparable to the recent GBT 21cm intensity detection \citep{Chang10}."
+" For a square array, the field of view is 10 beam width, or about 5 degrees."," For a square array, the field of view is 10 beam width, or about 5 degrees."
+" We use half the scanned area as useful, 2010)..allowing for galaxy and point source cuts."," We use half the scanned area as useful, allowing for galaxy and point source cuts."
+" At zenith for a latitude of 30 degrees, this scans 1000 square degrees."," At zenith for a latitude of 30 degrees, this scans 1000 square degrees."
+ We use an integration time of 1000 days on the sky., We use an integration time of 1000 days on the sky.
+ The exposure time per pixel is 5x106 seconds., The exposure time per pixel is $5\times 10^6$ seconds.
+" In a bandwidth of 2 MHz, this results in a pixel noise of 1 mK, well matched to the expected signal."," In a bandwidth of 2 MHz, this results in a pixel noise of 1 mK, well matched to the expected signal."
+" This map contains ~10° pixels, so one expects to measure the power on the beam scale with ~100σ."," This map contains $\sim 10^5$ pixels, so one expects to measure the power on the beam scale with $\sim 100 \sigma$."
+" The specific forecast parameters used for the noise estimate in figure 4 are: Ta,=300(v/150MHz)?7, 24 beams on the sky (the equivalent of a LOFAR station signal processor), and a diffraction limited beam."," The specific forecast parameters used for the noise estimate in figure \ref{pumpfig} are: $T_{\rm sky}=300 (\nu/150 {\rm MHz})^{-2.7}$, 24 beams on the sky (the equivalent of a LOFAR station signal processor), and a diffraction limited beam."
+" We have investigated the effect of relative motions between baryons and dark matter on the formation of the smallest galaxies at high redshift, z~20."," We have investigated the effect of relative motions between baryons and dark matter on the formation of the smallest galaxies at high redshift, $z\sim 20$."
+" The formation of galaxies in minihalos of mass MX109M, is modulated by bulk velocities between gas and dark matter, and so the clustering of these objects will contain a contribution proportional to the relative velocity two-point function."," The formation of galaxies in minihalos of mass $M\lesssim 10^6 M_\odot$ is modulated by large-scale bulk velocities between gas and dark matter, and so the clustering of these objects will contain a contribution proportional to the relative velocity two-point function."
+" The velocity power spectrum exhibits significant correlations on large scales of order 100 Mpc, with pronounced baryon acoustic oscillations."," The velocity power spectrum exhibits significant correlations on large scales of order 100 Mpc, with pronounced baryon acoustic oscillations."
+" Accordingly, the large-scale clustering of minihalos exhibits similar features, and the large-scale correlations of any observable that traces minihalos will show similar behavior."," Accordingly, the large-scale clustering of minihalos exhibits similar features, and the large-scale correlations of any observable that traces minihalos will show similar behavior."
+" We illustrated this in the previous section with a calculation of the 21 cm absorption power spectrum prior to reionization, when the kinetic temperature of intergalactic gas is much colder than the CMB temperature."," We illustrated this in the previous section with a calculation of the 21 cm absorption power spectrum prior to reionization, when the kinetic temperature of intergalactic gas is much colder than the CMB temperature."
+was investigated anc it was found that for values of the parameters which vield a power spectrum P(&) in. fair agreement with observations. the Doppler peak turns out to be low.,"was investigated and it was found that for values of the parameters which yield a power spectrum $P(k)$ in fair agreement with observations, the Doppler peak turns out to be low."
+ This is related to the effective tilt of the spectrum on very large scales., This is related to the effective tilt of the spectrum on very large scales.
+ Another possibilitv is to add a positive cosmological constant leaving the initial η&1 sspeetrum unchanged., Another possibility is to add a positive cosmological constant leaving the initial $n\simeq 1$ spectrum unchanged.
+ lt was long known that the cosmological constant is viable only if it is accompanied by cold. clark matter and. vice versa. its inclusion improves the CDM mocel a ereat deal (as emphasized c.g. in Ixofman Starobinsky (19852)Ehis model remains viable after the detection of the CMD anisotropies on Large angular scales by CODIS. (Ixofman. CGnedin Baheall 1993). and it is now perhaps the most promising CDM variant (Bagla. Pacimanabhan Navlikar 1995: Ostriker Steinhart 1995).," It was long known that the cosmological constant is viable only if it is accompanied by cold dark matter and, vice versa, its inclusion improves the CDM model a great deal (as emphasized e.g. in Kofman Starobinsky (1985)).This model remains viable after the detection of the CMB anisotropies on large angular scales by COBE (Kofman, Gnedin Bahcall 1993), and it is now perhaps the most promising CDM variant (Bagla, Padmanabhan Narlikar 1995; Ostriker Steinhardt 1995)."
+ As well known. one important motivation for a positive cosmological constant A is that it provides the possibility to accomodate both a high Llubble “constant”. 5>0.6. and a sullicienthy old universe. fyc9Ll Gyes.," As well known, one important motivation for a positive cosmological constant $\Lambda$ is that it provides the possibility to accomodate both a high Hubble “constant”, $h>0.6$, and a sufficiently old universe, $t_0>11$ Gyrs."
+ Also. the barvon fraction in clusters seems to imply Ox0.55 (Q=|Qy sstands for the total matter density including CDAL ancl barvons).," Also, the baryon fraction in clusters seems to imply $\Omega\le 0.55$ $\Omega=1-\Omega_\Lambda$ stands for the total matter density including CDM and baryons)."
+ The most recent strong argument in favour of 2«1 (and Qe(0.20.4)) {follows from the evolution of rich galaxy clusters (Baheall. Fan Cen 1997. sce also Fan. Dahcall Cen 1997).," The most recent strong argument in favour of $\Omega< 1$ (and $\Omega \simeq (0.2-0.4))$ follows from the evolution of rich galaxy clusters (Bahcall, Fan Cen 1997, see also Fan, Bahcall Cen 1997)."
+ Up to now. these two possibilities were considered. as mutually exclusive.," Up to now, these two possibilities were considered as mutually exclusive."
+ Now we want to unite them and compare he BSL CDM model including a cosmological constant with the observational data., Now we want to unite them and compare the BSI CDM model including a cosmological constant with the observational data.
+ The reasons for this are the ollowing., The reasons for this are the following.
+ First. it can enlarge. the allowed: cosmological parameters window.," First, it can enlarge the allowed cosmological parameters window."
+ Second. the possibility exists that the initial power spectrum of scalar (density) perturbations in he Universe is not scale [free but has instead. some non-rivial structure near &=0.05Alpe|.," Second, the possibility exists that the initial power spectrum of scalar (density) perturbations in the Universe is not scale free but has instead some non-trivial structure near $k=0.05 \,h\,\,{\mathrm Mpc^{-1}}$."
+ In fact observations may point to such a feature: the analysis of the threc-dimensional distribution of rich Abell galaxy clusters located in superclusters. performed in Einasto (L99Ta). pleads for an unexpected spatial quasi-periocicity of the data (see also Einasto 1997b. 19970).," In fact observations may point to such a feature: the analysis of the three-dimensional distribution of rich Abell galaxy clusters located in superclusters, performed in Einasto (1997a), pleads for an unexpected spatial quasi-periodicity of the data (see also Einasto 1997b, 1997c)."
+ Also. the spatial distribution of all Abell clusters of galaxies has a well-marked peak in the power spectrum at &c0.05fAlpe! (Einasto 1997a).," Also, the spatial distribution of all Abell clusters of galaxies has a well-marked peak in the power spectrum at $k\simeq 0.05\,h\,\,{\mathrm Mpc^{-1}}$ (Einasto 1997a)."
+ Phe Fourier power spectrum of the spatial distribution of APAL galaxies also has a feature on the same scale. though of a slightly clifferent form (Caztanaga Baugh. 1997).," The Fourier power spectrum of the spatial distribution of APM galaxies also has a feature on the same scale, though of a slightly different form (Caztanaga Baugh, 1997)."
+ Note that the natural attempt to explain this feature by Sakharov oscillations produced by the barvon acdmixture to CDM cloes not work (Atrio-DarandelaaL. 1997. EisensteinaL. 1997).," Note that the natural attempt to explain this feature by Sakharov oscillations produced by the baryon admixture to CDM does not work (Atrio-Barandela, 1997, Eisenstein, 1997)."
+ So this feature. if confirmed by future improved large scale structure observations. shoul be ascribed to the initial perturbation spectrum itself.," So this feature, if confirmed by future improved large scale structure observations, should be ascribed to the initial perturbation spectrum itself."
+ Therefore. we need an initial spectrum which has a non-trivial structure around: some scale. (preferably. with a bump) and has essentially no tilt at larger and. smaller scales.," Therefore, we need an initial spectrum which has a non-trivial structure around some scale (preferably, with a bump) and has essentially no tilt at larger and smaller scales."
+ The latter condition is necessary in order to have sullicienthy carly galaxy ancl quasar formation., The latter condition is necessary in order to have sufficiently early galaxy and quasar formation.
+ On. the other hand. this spectrum should be derivable from some first. principles (e.g.. it could. be generated. in a concrete inflationary model).," On the other hand, this spectrum should be derivable from some first principles (e.g., it could be generated in a concrete inflationary model)."
+ Such a spectrum naturally arises in a well-defined and rather generic (though idealized. of course) inflationary model where the inflaton potential V(4) hhas a local steplike feature in the first. derivative.," Such a spectrum naturally arises in a well-defined and rather generic (though idealized, of course) inflationary model where the inflaton potential $V(\varphi)$ has a local steplike feature in the first derivative."
+ An cxact analytical expression for the scalar. (densitv) perturbations ecncrated in this model was found in Starobinsky (1992)., An exact analytical expression for the scalar (density) perturbations generated in this model was found in Starobinsky (1992).
+ Lt has a universal shape depending on only one parameter p., It has a universal shape depending on only one parameter $p$.
+ Actually. it seems to be the only example of a perturbation spectrum with the desired: properties. lor which a closed analytical formi exists.," Actually, it seems to be the only example of a perturbation spectrum with the desired properties, for which a closed analytical form exists."
+ Thus. we suppose that the inflaton potential V(s) hhas a rapid change of slope in a neighborhood Ag oof yu: The resulting adiabatie perturbation spectrum is non- around the point Ay=e(to)H(lu). fy bbeing the time at which 4=yo while /f=ea is the Hubble parameter.," Thus, we suppose that the inflaton potential $V(\varphi)$ has a rapid change of slope in a neighborhood $\Delta \varphi$ of $\varphi_0$: The resulting adiabatic perturbation spectrum is non-flat around the point $k_0=a(t_0) H(t_0)$, $t_0$ being the time at which $\varphi=\varphi_0$ while $H\equiv \dot{a}/a$ is the Hubble parameter."
+ One can show (Starohinsky 1992) that if the width Ag oof the singularity is small enough. namely. Ay{1(fo)7min(.lai Aj). then the adiabatic perturbation spectrum has maximal deviation from fatness. and acequires a universal form that can be derived. analytically: where ® iis the (peculiar) gravitational potential.," One can show (Starobinsky 1992) that if the width $\Delta \varphi$ of the singularity is small enough, namely, $\Delta\varphi\,\,H(t_0)^2\ll\min(A_+,|A_+-A_-|)$ , then the adiabatic perturbation spectrum has maximal deviation from flatness, and acquires a universal form that can be derived analytically: where $\Phi$ is the (peculiar) gravitational potential."
+ This expression. plotted in Fig. L..," This expression, plotted in Fig. \ref{figPHI},"
+ depends (besides the overal normalization) on two parameters p aand. Ay., depends (besides the overall normalization) on two parameters $p$ and $k_0$.
+ Ehe shape of 10 spectrum does not depend on Au. Ay oonly. determines 1 location of the step.," The shape of the spectrum does not depend on $k_0$, $k_0$ only determines the location of the step."
+ For pz1. the spectrum has a Ila upper plateau on larger scales. even with a small bump. a v sharp decrease on smaller scales. with large oscillations rough.," For $p>1$, the spectrum has a flat upper plateau on larger scales, even with a small bump, and a sharp decrease on smaller scales, with large oscillations though."
+ For p«1 this picture is inverted., For $p<1$ this picture is inverted.
+ The ratio of over between the plateaux. equals. p. and for p=1 we just recover the (flat) scale-invariant Harrison-Zeldovich μα»ectrum.," The ratio of power between the plateaux equals $p^2$, and for $p=1$ we just recover the (flat) scale-invariant Harrison-Zel'dovich spectrum."
+ Note that this spectrum cannot be obtained in 1o slow-roll approximation (even with any finite number of acliabatic corrections to it)., Note that this spectrum cannot be obtained in the slow-roll approximation (even with any finite number of adiabatic corrections to it).
+ In this model it is still possible o fix freely the amount of primordial gravitational waves (CW's) for given p and normalization ancl we consider here he model with no GW's at all., In this model it is still possible to fix freely the amount of primordial gravitational waves (GW's) for given $p$ and normalization and we consider here the model with no GW's at all.
+ Without the inclusion of a cosmological constant. we would be forced to consider the case po»1 oonlv. in order to increase power on large scales.," Without the inclusion of a cosmological constant, we would be forced to consider the case $p>1$ only, in order to increase power on large scales."
+ Since AoO aalreacky procluces a desired excess of larec-scale, Since $\Lambda>0$ already produces a desired excess of large-scale
+for EOS L shows an unusually large relative error of 455€ in 53.,for EOS L shows an unusually large relative error of $45\%$ in $S_3$.
+ This. again. is related to the compactness of the various models and to the value of the multipole moments for ¢=0.," This, again, is related to the compactness of the various models and to the value of the multipole moments for $a=0$."
+ As a second test of the accuracy of the analytic solution for rapidly rotating neutron stars. we performed a direct comparison of specific metric components for several representative models. using all EOSs in our sample.," As a second test of the accuracy of the analytic solution for rapidly rotating neutron stars, we performed a direct comparison of specific metric components for several representative models, using all EOSs in our sample."
+ Here. we focus on the most rapidly rotating model of the maximum mass sequence with EOS FPS. since the other eases we examined showed similar behavior.," Here, we focus on the most rapidly rotating model of the maximum mass sequence with EOS FPS, since the other cases we examined showed similar behavior."
+ For this model we computed the metric components 2j. £o and go on the equatorial plane and along the symmetry axis using the analytic metric and the Kerr metric.," For this model we computed the metric components $g_{tt}$, $g_{t\phi}$ and $g_{\phi\phi}$ on the equatorial plane and along the symmetry axis using the analytic metric and the Kerr metric."
+ Then we compared the relative error of both metrics with respect to the corresponding components of the numerical metric., Then we compared the relative error of both metrics with respect to the corresponding components of the numerical metric.
+ Fig., Fig.
+" 5 shows the relative error of the ¢,,-component of the analytic metric and of the Kerr metric in the equatorial plane.", \ref{nu_equat} shows the relative error of the $g_{tt}$ -component of the analytic metric and of the Kerr metric in the equatorial plane.
+ For the analytic metric the error is only 0.3% at the surface of the star (located. at p= 10.6). and decreases monotonically with increasing distance. becoming of order 10 near intiniuw," For the analytic metric the error is only $0.3 \% $ at the surface of the star (located at $\tilde \rho=10.6$ ), and decreases monotonically with increasing distance, becoming of order $10^{-6}$ near infinity."
+ In comparison. the relative difference between the Kerr metric and the numerical metric is 1.3% at the equator (i.e. four times larger than the error in the analytic metric).," In comparison, the relative difference between the Kerr metric and the numerical metric is $1.3\%$ at the equator (i.e., four times larger than the error in the analytic metric)."
+ The relative difference between the Kerr metric and the numerical metric also decreases with increasing distance. as expected. and for distances larger than about 200 equatorial radii. the difference between the analytic solution and the Kerr solution is negligible.," The relative difference between the Kerr metric and the numerical metric also decreases with increasing distance, as expected, and for distances larger than about 200 equatorial radii, the difference between the analytic solution and the Kerr solution is negligible."
+ In other words. at such a distance the error in the analytic. solution is dominated by the Kerr contribution at first order in the rotational parameter. while the effects of the higher-order multipole moments Q and 53 have become unimportant.," In other words, at such a distance the error in the analytic solution is dominated by the Kerr contribution at first order in the rotational parameter, while the effects of the higher-order multipole moments $Q$ and $S_3$ have become unimportant."
+ The corresponding figure forthe relative error in 299 on the equatorial plane is nearly identical to Fig., The corresponding figure forthe relative error in $g_{\phi \phi}$ on the equatorial plane is nearly identical to Fig.
+ 5 for οµ., \ref{nu_equat} for $g_{t t}$.
+" When we consider 2,9 in the equatorial plane. the relative error at the surface is 1.3% for the analytic metric and 5.3% for the Kerr metric."," When we consider $g_{t \phi}$ in the equatorial plane, the relative error at the surface is $1.3 \%$ for the analytic metric and $5.3\%$ for the Kerr metric."
+" This larger error for 2,4 should be expected: this metric component vanishes in the nonrotating limit. so it is more sensitive to contributions by the higher-order multipole moments Q and Ss than the metric components ὁμ and eoo."," This larger error for $g_{t \phi}$ should be expected: this metric component vanishes in the nonrotating limit, so it is more sensitive to contributions by the higher-order multipole moments $Q$ and $S_3$ than the metric components $g_{tt}$ and $g_{\phi \phi}$ ."
+ In order to compare the metric components on the symmetry UNIS. we first need to integrate the Cauchy-Riemann conditions (75)) and obtain the coordinate = in terms of the coordinate. <.," In order to compare the metric components on the symmetry axis, we first need to integrate the Cauchy-Riemann conditions \ref{CauRie}) ) and obtain the coordinate $\tilde z$ in terms of the coordinate $z$."
+ This ean be done easily. once the numerical solution for the metric function B is obtained.," This can be done easily, once the numerical solution for the metric function $B$ is obtained."
+ Fig., Fig.
+" 6. shows the relative error in ¢,, for the analytic solution and the Kerr solution along the symmetry axis.", \ref{nu_offaxis} shows the relative error in $g_{tt}$ for the analytic solution and the Kerr solution along the symmetry axis.
+ The location of the surface (as determined from the numerical solution) is at z=6.05., The location of the surface (as determined from the numerical solution) is at $\tilde z=6.05$.
+" At the surface. the relative error for the analytic solution is Το, while it is 155€ for the Kerr metric."," At the surface, the relative error for the analytic solution is $7\%$, while it is $15\%$ for the Kerr metric."
+ Thus. we see that the effect of a large quadrupole moment Q shows up oedominantly in the metric components along the symmetry axis. while in the equatorial plane this effect is very small.," Thus, we see that the effect of a large quadrupole moment $Q$ shows up predominantly in the metric components along the symmetry axis, while in the equatorial plane this effect is very small."
+ The reason for his difference is that a rapidly rotating star becomes very oblate. so that the stellar surface on the symmetry axis is located deeper in the gravitational potential well than the surface in the equatorial vane.," The reason for this difference is that a rapidly rotating star becomes very oblate, so that the stellar surface on the symmetry axis is located deeper in the gravitational potential well than the surface in the equatorial plane."
+ The specific example shown in the above figures has polar to equatorial axes ratio of 0.6. thus the equatorial radius is roughly wice as large as the polar radius.," The specific example shown in the above figures has polar to equatorial axes ratio of $0.6$, thus the equatorial radius is roughly twice as large as the polar radius."
+" The analytic value of 9, on qe surface in the equatorial plane is 0.59. while itis 0.44 on je surface on the symmetry axis (the asymptotic value at large distances is Ly."," The analytic value of $g_{tt}$ on the surface in the equatorial plane is $-0.59$, while it is $-0.44$ on the surface on the symmetry axis (the asymptotic value at large distances is $-1$ )."
+" Gravity is stronger on the polar surface. and this justifies a larger relative error in ο), there."," Gravity is stronger on the polar surface, and this justifies a larger relative error in $g_{tt}$ there."
+" At about 3 polar radii. qe relative error in ¢,, along the symmetry axis decreases to the Πές level for both the analytic and Kerr solutions."," At about 3 polar radii, the relative error in $g_{tt}$ along the symmetry axis decreases to the $1\%$ level for both the analytic and Kerr solutions."
+ The above direct comparison of metric components shows that je analytic metric is a good approximation to the numerical one (Cor. at least. a much better approximation than the Kerr metric) in the equatorial plane. where one expects particle orbits to be astrophysically more relevant.," The above direct comparison of metric components shows that the analytic metric is a good approximation to the numerical one (or, at least, a much better approximation than the Kerr metric) in the equatorial plane, where one expects particle orbits to be astrophysically more relevant."
+ For gravitational-wave extraction in numerical relativity. the larger inaccuracies near the polar surface could influence the waveforms.," For gravitational-wave extraction in numerical relativity, the larger inaccuracies near the polar surface could influence the waveforms."
+ In order to minimize this effect. the extraction should be done as far as possible from the surface of the star.," In order to minimize this effect, the extraction should be done as far as possible from the surface of the star."
+ In any case. the analytic metric is everywhere more accurate than the Kerr metric.," In any case, the analytic metric is everywhere more accurate than the Kerr metric."
+ Thus a perturbative wave-extraction scheme. built with the analytic metric as a background. should yield more accurate Waveforms than those obtained with techniques available at present (which are based on a perturbative extraction of waveforms around a Sehwarzschild or Kerr background).," Thus a perturbative wave-extraction scheme, built with the analytic metric as a background, should yield more accurate waveforms than those obtained with techniques available at present (which are based on a perturbative extraction of waveforms around a Schwarzschild or Kerr background)."
+ Tt is well known that not all orbits around relativistic stars are stable., It is well known that not all orbits around relativistic stars are stable.
+ For nonrotating stars. the ISCO is located at a circumferential radius of Αι= 6M.," For nonrotating stars, the ISCO is located at a circumferential radius of $R_{ISCO}=6M$ ."
+ Depending on the EOS and the mass of thestar. the ISCO can be located outside the stellar surface.," Depending on the EOS and the mass of thestar, the ISCO can be located outside the stellar surface."
+ Rotation introduces a preferred direction in the $ coordinate. so ISCOs," Rotation introduces a preferred direction in the $\phi$ coordinate, so ISCOs"
+As a result. the halo becomes triaxial in shape.,"As a result, the halo becomes triaxial in shape."
+ This phenomenon has been examined in detail by DejougheMerritt(1988) using Nbody simiulatious. by (1991) using linear analysis. and by Tissetal.(1999) and Alac\Gillanctal.(2006) for the specific case of isolated dark matter halo collapse.," This phenomenon has been examined in detail by \citet{Dejonghe88} using N–body simulations, by \citet{
+Weinberg91} using linear analysis, and by \citet{Huss99} and \citet{MacMillan06} for the specific case of isolated dark matter halo collapse."
+ These eroups fiud that he onset of the ROI correspouds to a flattening of the central density cusp of the halo. aud may be responsible or the double power-law shape of the halo.," These groups find that the onset of the ROI corresponds to a flattening of the central density cusp of the halo, and may be responsible for the double power-law shape of the halo."
+ The ROI is reviewed in Mexitt (1999: 86.2)., The ROI is reviewed in Merritt (1999; 6.2).
+ Seni-anualvtie models have also been used to examine he spherical collapse of dark matter halos. begimnius with Camu&Cott(1972).," \nocite{Merritt99}
+ Semi-analytic models have also been used to examine the spherical collapse of dark matter halos, beginning with \citet{Gunn72}."
+. These models (Cott1975:Gunn1977:Bertschinger1985) inchide only radial notions and produce single power-law density profiles.," These models \citep{Gott75, Gunn77, Bertschinger85} include only radial motions and produce single power-law density profiles."
+ When nouradial motions are included. halo density xofiles range from power-luvs (Ryden&Coun1987) to NEWlike (Motels2002:LeDelliou&IIeuriksenAscasibaretal. 2001).," When non–radial motions are included, halo density profiles range from power-laws \citep{Ryden87} to NFW–like \citep{Hiotelis02,LeDelliou03,Ascasibar04}."
+. Receuth. however. Barnes(2005). extended the senmi-aualvtic model of Williamsetal.(2001) to include a plivsical representation of the ROI.," Recently, however, \citet{Barnes05} extended the semi-analytic model of \citet{Williams04} to include a physical representation of the ROI."
+ The resulting halos have a double power-law density slope simular to those of Nbody simulations., The resulting halos have a double power-law density slope similar to those of N–body simulations.
+ Darnesetal.(2005) concluded. that the density scale leneth and the anisotropy radius are correlated. aud that the ROT is directly responsible for the shape of the density profile.," \citet{Barnes05} concluded that the density scale length and the anisotropy radius are correlated, and that the ROI is directly responsible for the shape of the density profile."
+ While Barnesctal. (2005)s work is suggestive. the complexities of a ανασα. instability are not casily captured bv analytic or scianalytic techniques.," While \citet{Barnes05}' 's work is suggestive, the complexities of a dynamical instability are not easily captured by analytic or semi–analytic techniques."
+ We therefore turn to studving the link between the ROI and halo structure by using high resolution Nbody simulations., We therefore turn to studying the link between the ROI and halo structure by using high resolution N–body simulations.
+ Iu this paper. we exanuue in detail the effect of random motions on the onset of the ROI aud the final structure of collapsed dark iatter halos. aud attempt to verity the relation between the scale leneth aud the anisotropy radius reported by Barnesctal.(2005).," In this paper, we examine in detail the effect of random motions on the onset of the ROI and the final structure of collapsed dark matter halos, and attempt to verify the relation between the scale length and the anisotropy radius reported by \citet{Barnes05}."
+. We analyze a set of Nbody simulations of isolated. collapsing dark matter halos with a variety of initial velocity dispersious to study the evolution of halo properties.," We analyze a set of N–body simulations of isolated, collapsing dark matter halos with a variety of initial velocity dispersions to study the evolution of halo properties."
+ Considering a ranee of velocity dispersions allows us to supress the ROI in some cases. helpiug us to isolate its physical effects.," Considering a range of velocity dispersions allows us to supress the ROI in some cases, helping us to isolate its physical effects."
+ Tn 822 we describe our sinnulatious aud show that thev are robust to resolution and softeuiuse effects., In 2 we describe our simulations and show that they are robust to resolution and softening effects.
+ We describe the properties of the halos iu 833. including the shape aud anisotropy evolution. the effects of velocity dispersion. and the scale-leneth anisotropy radius relation.," We describe the properties of the halos in 3, including the shape and anisotropy evolution, the effects of velocity dispersion, and the scale-length – anisotropy radius relation."
+ We sumuuuuizeourresults =in &ll., We summarize our results in 4.
+ We performed simulationsa. using PIKRDGRAV (Stadcl2001:Wadsleyctal. 2001).. a parallel IND Tree eravity solver.," We performed simulations using PKDGRAV \citep{Stadel01,Wadsley04}, a parallel KD Tree gravity solver."
+ The initial svstem is an isolated. spherical halo with a Gaussian radial density profile. mimicking au overdensity in the early universe.," The initial system is an isolated, spherical halo with a Gaussian radial density profile, mimicking an overdensity in the early universe."
+ All particles are given IIubble flow velocities. such that the halo is initially expanding.," All particles are given Hubble flow velocities, such that the halo is initially expanding."
+ We evolve the svstem in plysical coordinates for 12.5 Gyr. which was set to allow of the mass of the halo to collapse by the preseut time: this corresponds to a starting redshift of ~12 according to the most receut results from WALAP (Spereeletal.2007).," We evolve the system in physical coordinates for 12.8 Gyr, which was set to allow of the mass of the halo to collapse by the present time; this corresponds to a starting redshift of $\sim12$ according to the most recent results from WMAP \citep{WMAP3}."
+. We ran several simulations with comparable deusitv profiles but with a rauge of initial velocity dispersions. allowing us to estimate the threshold at which the ROI will occur.," We ran several simulations with comparable density profiles but with a range of initial velocity dispersions, allowing us to estimate the threshold at which the ROI will occur."
+ In principle. a dynamically παλιο system should resist the ROI (Merritt&Aeuilar1985).. since anv perturbations will be washed out bv the tangential velocity. dispersious.," In principle, a dynamically “warmer” system should resist the ROI \citep{Merritt85}, since any perturbations will be washed out by the tangential velocity dispersions."
+ We assigned initial random velocities (0) to the particles assuniusg a Gaussian distributionwith a iieau of zero: these uunuboers are then added to the existing Ifubble flow velocities., We assigned initial random velocities $\sigma$ ) to the particles assuming a Gaussian distributionwith a mean of zero; these numbers are then added to the existing Hubble flow velocities.
+ The wuplitudes of these initial velocity dispersions were 1σ. 2a. aud 30. where o6 d: of the circular velocity at the virial radius of the final svstem.," The amplitudes of these initial velocity dispersions were $1\sigma$, $2\sigma$, and $3\sigma$, where $\sigma$ is of the circular velocity at the virial radius of the final system."
+ Simulations with a higher velocity dispersion are too wari to collapse at all., Simulations with a higher velocity dispersion are too warm to collapse at all.
+ We also carried out simmlations with no velocity dispersion. at both standard and high resolution.," We also carried out simulations with no velocity dispersion, at both standard and high resolution."
+ Finally. we attempted to isolate the iuportauce of tangeutial velocity dispersions bv carrving out a simulation witli a pure radial initial velocity dispersion of 3e.," Finally, we attempted to isolate the importance of tangential velocity dispersions by carrying out a simulation with a pure radial initial velocity dispersion of $3\sigma$."
+ Table 1 lists the initial value of the ratio 277W. (see refsec:elobalanis)) for cach simulation for a more intuitive erasp of the value of cach σ., Table 1 lists the initial value of the ratio $-2T/W$ (see \\ref{sec:globalanis}) ) for each simulation for a more intuitive grasp of the value of each $\sigma$.
+ Tn addition to various velocity dispersions. we ran siuulations with a range of softeniuss and particle resolutions.," In addition to various velocity dispersions, we ran simulations with a range of softenings and particle resolutions."
+ Our highest resolution simulation has over 2.8 million. particles. while the simulations testing the effects of velocity dispersion have about one fifth of that nuuber.," Our highest resolution simulation has over 2.8 million particles, while the simulations testing the effects of velocity dispersion have about one fifth of that number."
+ These staucdard-resolution simmlatious are comparable i nuniber of particles to MacMillanetal.(2006).. but are much larger than the simulatious iu IIussetal.(1999)... which have —10.000 particles within the virial radius of cach halo.," These standard-resolution simulations are comparable in number of particles to \citet{MacMillan06}, but are much larger than the simulations in \citet{Huss99}, which have $\sim$ 10,000 particles within the virial radius of each halo."
+ The softening. e. was sot to 0.051 times royy of the final collapsed. halo. aud is the same for all sinmulatiouns.," The softening, $\epsilon$, was set to 0.054 times $r_{200}$ of the final collapsed halo, and is the same for all simulations."
+ Time steps were sot to an accuracy criterion proportional to νεα where jj=0.2 is the thuecstep criterion and e is the acceleration of a particle (Wadsleyetal.2001)., Time steps were set to an accuracy criterion proportional to $\eta\sqrt{\epsilon/a}$ where $\eta = 0.2$ is the timestep criterion and $a$ is the acceleration of a particle \citep{Wadsley04}.
+.. We use a force accuracy criterion of 0.55., We use a force accuracy criterion of 0.55.
+" Α παπα of all simulations is in Table 1,", A summary of all simulations is in Table 1.
+ Figure d. shows how the structure of our simulated halos are affected by particle nuniber (also referred to in the text as resolution)., Figure \ref{fig:restest} shows how the structure of our simulated halos are affected by particle number (also referred to in the text as resolution).
+ The top pancls show the density profile for the fully evolved. 060 halo at standard resolution (left) aud at high resolution (right).," The top panels show the density profile for the fully evolved, $0\sigma$ halo at standard resolution (left) and at high resolution (right)."
+ Each profile is ft with a NEW profile (purple dotted linc) and a Navarroetal.(2001) profile (also kuowun as au Einasto profile) (red dashed line)., Each profile is fit with a NFW profile (purple dotted line) and a \citet{Navarro04} profile (also known as an Einasto profile) (red dashed line).
+ The profiles for the two runs are simular over a large range in radius., The profiles for the two runs are similar over a large range in radius.
+ However. as can be seen by the residuals from the fits. the standard resolution simulation does not adequately model the inner core of the halo (r<0.05rogo. where we define royy as the virial radius in which the enclosed density is 200 times the critical deusitv of the universe.," However, as can be seen by the residuals from the fits, the standard resolution simulation does not adequately model the inner core of the halo $r < 0.05r_{200}$, where we define $r_{200}$ as the virial radius in which the enclosed density is 200 times the critical density of the universe."
+ There is also a differeuce iu the innerinost cores when looking at the +=0 axis ratio profiles (Fie. 1..," There is also a difference in the innermost cores when looking at the $z=0$ axis ratio profiles (Fig. \ref{fig:restest},"
+ muddle paucls)., middle panels).
+ The lig[um run (right) shows a somewhat more triaxial core. particularly at the ceuter.," The high--resolution run (right) shows a somewhat more triaxial core, particularly at the center."
+ However. the anisotropy and phasespace density profiles show no siguificaut differeuces between the two smulatious (Fig. 1.. ," However, the anisotropy and phase–space density profiles show no significant differences between the two simulations (Fig. \ref{fig:restest}, ,"
+bottom panels. left and right. respectively).," bottom panels, left and right, respectively)."
+ We analyze these profiles in more detail in 833., We analyze these profiles in more detail in 3.
+ The bulk of the difference between the two resolutions is apparent oulv within ~ rogo., The bulk of the difference between the two resolutions is apparent only within $\sim0.05 \times r_{200}$ .
+ Outside of this radius our mecdimmresolution simulations are robust to resolution effects. while iuside this radius the profiles max be less certain. aud biased," Outside of this radius our medium–resolution simulations are robust to resolution effects, while inside this radius the profiles may be less certain, and biased"
+galaxies. we compare their colours z FSIAN1.92.00 with the svnthesized galaxv colours in. bPukugita. Shimasaku Lchilkawa (1995).,"galaxies, we compare their colours $\approx$ 1.9--2.0) with the synthesized galaxy colours in Fukugita, Shimasaku Ichikawa (1995)."
+ For the cüllerent galaxy tvpes approximate photometric redshifts of 0.3. (02). 0.4 (80). 0.4 (Sab). 0.7 (She) and 0.9 (Sed) are found.," For the different galaxy types approximate photometric redshifts of 0.3 (E), 0.4 (S0), 0.4 (Sab), 0.7 (Sbc) and 0.9 (Scd) are found."
+ GI and G2 could herefore. be carly type galaxies. (including Sab) at low redshift (0.3.0.4) or late type galaxies at high redshift (0.71.9)., G1 and G2 could therefore be early type galaxies (including Sab) at low redshift (0.3–0.4) or late type galaxies at high redshift (0.7--0.9).
+" The integrated luminosities of Gl and G2 in 2 band or the different twpes of galaxies are: log),(Ly.)2S4Ploe(hay) (E). S.S.21og(hzo) (SO and Sab). 9.621og(hzu) (She) and OS—2log(hzu) (Sed) for Gl (lg=50-hz, km + 1). where we used the photometric redshifts found above and the £2 FSI4W colours and. Ix. corrections. from Fukugita ct al. ("," The integrated luminosities of G1 and G2 in $B$ band for the different types of galaxies are: $\log_{10}(L_{{\rm B}\odot})
+\approx 8.4-2\log(h_{50})$ (E), $8.8-2\log(h_{50})$ (S0 and Sab), $9.6-2\log(h_{50})$ (Sbc) and $9.8-2\log(h_{50})$ (Scd) for G1 ${\rm
+H}_{0}=50\cdot h_{50}$ km $^{-1}$ $^{-1}$ ), where we used the photometric redshifts found above and the $B$ –F814W colours and K corrections from Fukugita et al. ("
+1995).,1995).
+ For G2 these values are 0.4 lower., For G2 these values are 0.4 lower.
+ Ata redshift 2.=0.3. £=22.5 (G1) corresponds to an absolute magnitude M;~19Sloelhso).," At a redshift $z_{\rm d}=0.3$, $I$ =22.5 (G1) corresponds to an absolute magnitude $M_I \sim-19-5\log(h_{50})$."
+ For an E and SO type galaxy this would mean it is 4 mag underluminous compared to I2 and SO type galaxies in the Hubble Deep Field (Alobasher et al., For an E and S0 type galaxy this would mean it is $\sim$ 4 mag underluminous compared to E and S0 type galaxies in the Hubble Deep Field (Mobasher et al.
+ 1996)., 1996).
+ Placing the galaxy at higher redshifts would make the V£ colours of Gl anc C2 inconsistent with those of I or SO tvpe galaxies (Fukugita et al., Placing the galaxy at higher redshifts would make the $V-I$ colours of G1 and G2 inconsistent with those of E or S0 type galaxies (Fukugita et al.
+ 1995)., 1995).
+ Phe absolute Z magnitude is consistent however with somewhat later tvpe spiral galaxies at higher redshifts (σαX 0.5)., The absolute $I$ magnitude is consistent however with somewhat later type spiral galaxies at higher redshifts $z_{\rm d}\ga0.5$ ).
+ This would also explain why the velocity dispersions of Gl and C2 appears significantly smaller than those expected for L; E and S0 type galaxies (e.g. Ixochanek 1905. 1994).," This would also explain why the velocity dispersions of G1 and G2 appears significantly smaller than those expected for $L_*$ E and S0 type galaxies (e.g. Kochanek 1993, 1994)."
+ Using the velocity. dispersions listed in Fable 4 (mocel ID. the mass-to-light. ratios (using the mass inside the Einstein radius. the photometric lens redshifts and the total," Using the velocity dispersions listed in Table 4 (model II), the mass-to-light ratios (using the mass inside the Einstein radius, the photometric lens redshifts and the total"
+(qan) it closely relies on physical processes in the solar atmosphere.,(iii) it closely relies on physical processes in the solar atmosphere.
+ The signal in each pixel is a mixture of emission from different plasmas along the LOS and from unresolved fine structures., The signal in each pixel is a mixture of emission from different plasmas along the LOS and from unresolved fine structures.
+ We have decomposed them by a statistical method based on the assumption that statistically different components can be separated in the radiance distribution., We have decomposed them by a statistical method based on the assumption that statistically different components can be separated in the radiance distribution.
+ In our new method. we differentiate among different classes of brightness in the radiance distribution instead of simply averaging the brightness over all pixels.," In our new method, we differentiate among different classes of brightness in the radiance distribution instead of simply averaging the brightness over all pixels."
+ We assume that individual bins of different brightness do behave differently., We assume that individual bins of different brightness do behave differently.
+ In. particular. we make use of the fact that brighter pixels have a tendency to appear redshifted in many emission lines.," In particular, we make use of the fact that brighter pixels have a tendency to appear redshifted in many emission lines."
+ Such a brightness relationship was also noted by Dammaschetal.(2008) in a different data set. and must consequently have an imprint on the network contrast.," Such a redshift-to-brightness relationship was also noted by \citet{Dammasch08} in a different data set, and must consequently have an imprint on the network contrast."
+ This is the core of our method: the strong redshift observed in the network contrast as compared to the position of the line itself., This is the core of our method: the strong redshift observed in the network contrast as compared to the position of the line itself.
+ For a collection of prominent. blend-free emission lines. we determined the network contrast (ratio of of the brighter pixels às compared to of the dim pixels) and comparec the contrast curve to the emission line itself.," For a collection of prominent, blend-free emission lines, we determined the network contrast (ratio of of the brighter pixels as compared to of the dim pixels) and compared the contrast curve to the emission line itself."
+ A rather crude and empirical method was used in the SDA to display the network contrast: the radiance of a few bright pixels out of 300 along the slit. which were thought to represent the network. was compared to the radiance of the remaining pixels.," A rather crude and empirical method was used in the SDA to display the network contrast; the radiance of a few bright pixels out of 300 along the slit, which were thought to represent the network, was compared to the radiance of the remaining pixels."
+ The pixel selection was made 1 ἃ static way for all 36 individual exposures of the data set. a procedure that may not be appropriate in view of the temporal evolution of the Sun.," The pixel selection was made in a static way for all 36 individual exposures of the data set, a procedure that may not be appropriate in view of the temporal evolution of the Sun."
+ However. the shift of the slit image caused by the wavelength scan (due to misalignment of the grating) had been compensated for by employing the standard ppixel routine.," However, the shift of the slit image caused by the wavelength scan (due to misalignment of the grating) had been compensated for by employing the standard pixel routine."
+ In the present analysis. we measure and compare the position of the line centroids (the line center determined by spectral centroiding) relative to the position. of the contrast maximum or minimum.," In the present analysis, we measure and compare the position of the line centroids (the line center determined by spectral centroiding) relative to the position of the contrast maximum or minimum."
+ We find that the centroid of the contrast profile is normally shifted by several pixels towards longer wavelengths., We find that the centroid of the contrast profile is normally shifted by several pixels towards longer wavelengths.
+ The data set in the SDA its a snapshot of 300 pixels along the slit., The data set in the SDA is a snapshot of 300 pixels along the slit.
+ Thus. the given spectral radiances still have significant uncertainties and can only approximately represent the quiet Sun.," Thus, the given spectral radiances still have significant uncertainties and can only approximately represent the quiet Sun."
+ Also. better values for the network contrast could have been achieved if a better statistical basis were available like the one we report here. where rasters are observed instead of a single exposure.," Also, better values for the network contrast could have been achieved if a better statistical basis were available like the one we report here, where rasters are observed instead of a single exposure."
+ Our data set consists of raster scans of size x in 14 different wavelength windows of = 44 covering the entire wavelength range from 670 to 1490 with only few insignificant gaps., Our data set consists of raster scans of size $\times$ in 14 different wavelength windows of $\approx$ 44 covering the entire wavelength range from 670 to 1490 with only few insignificant gaps.
+" We employed a slit of size 1x120"".", We employed a slit of size $\times$.
+ The raster increment was and the exposure time was 90 s. The rasters were obtained in the so-called *“Schmierschritt’ mode. which means that each transmitted spectrum is composed of four elementary exposures with a 22.5 s dwell time and a step increment (Wilhelmetal.. 1995).," The raster increment was and the exposure time was 90 s. The rasters were obtained in the so-called 'Schmierschritt' mode, which means that each transmitted spectrum is composed of four elementary exposures with a 22.5 s dwell time and a step increment \citep{Wilhelm95}."
+". We already compensated on board for the parasitic movement of the slit image mentioned earlier,", We already compensated on board for the parasitic movement of the slit image mentioned earlier.
+ Therefore. we can safely assume that the North-South offset between the individual rasters is negligible.," Therefore, we can safely assume that the North-South offset between the individual rasters is negligible."
+ We have compensated for the solar rotation after each raster scan., We have compensated for the solar rotation after each raster scan.
+ Therefore. we also assume that all rasters do map the same portion of the Sun in the direction.," Therefore, we also assume that all rasters do map the same portion of the Sun in the East-West direction."
+ Our data set was obtained during an observation on 5 April 2007 running from 00:5] UTC to 13:10 UTC., Our data set was obtained during an observation on 5 April 2007 running from 00:51 UTC to 13:10 UTC.
+ The initial pointing — centre of the first raster — was x=0. v=0.," The initial pointing – centre of the first raster – was $x=0$, $y=0$."
+ Fifty minutes are needed for each raster., Fifty minutes are needed for each raster.
+ Standard procedures from the SUMERsoft were applied for the data reduction., Standard procedures from the SUMERsoft were applied for the data reduction.
+ The major improvement of this observation (called. the ‘super atlas’) as compared to a normal reference spectrum is the increase of the number of pixels by more than an order of magnitude., The major improvement of this observation (called the 'super atlas') as compared to a normal reference spectrum is the increase of the number of pixels by more than an order of magnitude.
+ This data set allows us to produce monochromatic raster scans for all emission lines and all continua in. the SUMER spectral range., This data set allows us to produce monochromatic raster scans for all emission lines and all continua in the SUMER spectral range.
+ As an example. we show in Fig.," As an example, we show in Fig."
+ 2 the maps obtained simultaneously in the emission ofiv... and in the continuum around 780A.," 2 the maps obtained simultaneously in the emission of, and in the continuum around 780."
+ We have selected 19 prominent and blend-free emission lines to produce monochromatic maps., We have selected 19 prominent and blend-free emission lines to produce monochromatic maps.
+ For each raster we also produce a map of the continuum., For each raster we also produce a map of the continuum.
+ We use this continuum map. where the elements of the chromospheric network are well-structured and at instrument resolution. to rank all pixels by radiance.," We use this continuum map, where the elements of the chromospheric network are well-structured and at instrument resolution, to rank all pixels by radiance."
+ We define equally-sized bins of bright network, We define equally-sized bins of bright network
+"energv in the GRB pulse. v, is the lrequeney. at the peak of the spectiun. and. is the high energv spectral index.","energy in the GRB pulse, $\nu_p$ is the frequency at the peak of the spectrum, and $\beta$ is the high energy spectral index."
+ Approximately 10° pairs can be created per ISM electron when conditions are right. but the number can be much less in other situations.," Approximately $^3$ pairs can be created per ISM electron when conditions are right, but the number can be much less in other situations."
+ If the optical depth of the pair screen becomes 1. further pair creation is expected to be quenched ancl so the optical depth will not increase much above unity.," If the optical depth of the pair screen becomes $\sim 1$, further pair creation is expected to be quenched and so the optical depth will not increase much above unity."
+ Let (he medium in (he vicinity of the GRB explosion be stratified with clensity decreasing ἂν τν αν appropriate for a wind [rom the pre-supernova star: nir)=njtr/rg)7.," Let the medium in the vicinity of the GRB explosion be stratified with density decreasing as $r^{-2}$, as appropriate for a wind from the pre-supernova star: $n(r)= n_0 (r/r_0)^{-2}$."
+ For a mass loss rate of LO? |. which is average for a Woll-Ravet star. and a wind speed of 10* km Ll (he number density at a radius of LOY cem is about 105 7.," For a mass loss rate of $^{-5}$ $^{-1}$, which is average for a Wolf-Rayet star, and a wind speed of $^3$ km $^{-1}$, the number density at a radius of $^{15}$ cm is about $^6$ $^{-3}$."
+ The optical depth in the wind is dominated by the radius al which the GRB photons are produced (where the density is highest)., The optical depth in the wind is dominated by the radius at which the GRB photons are produced (where the density is highest).
+ In the internal shock model for GRBs this radius is a ry~lewx10! cem., In the internal shock model for GRBs this radius is a $r_1\sim {\rm few}\times 10^{14}$ cm.
+ Let us assume (hal the pair screen forms between rj and an outer radius rs., Let us assume that the pair screen forms between $r_1$ and an outer radius $r_2$.
+ Let us also imagine that. on average. each ISM electron produces 7-2 pairs.," Let us also imagine that, on average, each ISM electron produces $\eta_\pm$ pairs."
+" The optical depth of the pair screen is then Assuming a mass 27, per ISM electron. the mass of the screen is where £; is the jet opening angle."," The optical depth of the pair screen is then Assuming a mass $2m_p$ per ISM electron, the mass of the screen is where $\theta_j$ is the jet opening angle."
+ Let the pair screen move with Lorentz [actor 55 as a result of absorbing momentum from the GRB photons., Let the pair screen move with Lorentz factor $\gamma_\pm$ as a result of absorbing momentum from the GRB photons.
+ The momentum of the screen may be written in the form The momentum taken out of the GRB pulse is Leenxmin(z. D)/(052). where the factor +_~> allows forH the factH that. forH a stationary. observer. the scattered photons move at an angle of 1/55 with respect to the radial direction.," The momentum of the screen may be written in the form The momentum taken out of the GRB pulse is $E_{GRB}\times\min(\tau,1)/(c\gamma_\pm^2)$ , where the factor $\gamma_\pm^{-2}$ allows for the fact that, for a stationary observer, the scattered photons move at an angle of $1/\gamma_\pm$ with respect to the radial direction."
+" Equating (he momentum taken out of the GRB pulse to P? we obtain an estimate lor the Lorentz [actor of the pair screen: For E;,=0.5. 6;=0.2. i=10. rq45=2415/40.5 and 7€1 we lind 5=2.1."," Equating the momentum taken out of the GRB pulse to $P$ we obtain an estimate for the Lorentz factor of the pair screen: For $E_{51}=0.5$, $\theta_j=0.2$, $\eta_\pm=10$, $r_{1,15}=r_{2,15}/4=0.5$ and $\tau\le1$ we find $\gamma_\pm=2.1$."
+ The total energy scattered backward is about a factor of 13 less than the energy in the original GRB pulse because of Doppler ce-amplification., The total energy scattered backward is about a factor of 18 less than the energy in the original GRB pulse because of Doppler de-amplification.
+ However. the peak of the spectrum of the back-scattered radiation is à [actor of ~4 smaller thanthe peak of the GRD spectrum.," However, the peak of the spectrum of the back-scattered radiation is a factor of $\sim4$ smaller thanthe peak of the GRB spectrum."
+ If the, If the
+jehavior found by GP9S8 where the constant. value of the rate lor reaction (1) caused the steady drop of LiH. The final (2=0) value of LiH is —7x10IS itl au inerease of a factor of ~το 'elative to QGP98s. (Gi),behavior found by GP98 where the constant value of the rate for reaction (1) caused the steady drop of LiH. The final $z=0$ ) value of LiH is $\sim 7\times 10^{-18}$ with an increase of a factor of $\sim 70$ relative to GP98. )
+) Die to the differences in behavior showi by the cross sectious iu Figure 1. the ionic yartuer. . behaves differently at Vi," Due to the differences in behavior shown by the cross sections in Figure 1, the ionic partner, $^+$ , behaves differently at $z\lesssim 30$."
+rst. (ie iucreased elliciency of destruction (reaction 3: ‘ate d4 ) limits the slarp rise in abuickuice al z20.," First, the increased efficiency of $^+$ destruction (reaction 3; rate $d_3^+$ ) limits the sharp rise in abundance at $z \sim 20$."
+ Then. the gentle decline below zoD ds due to the efficiency of eectronie recoilijuation (rate dDT ," Then, the gentle decline below $z \sim 5$ is due to the efficiency of electronic recombination (rate $d_4^+$ )."
+he new final abundances of (solid ine) is now sxinaller tan earlier estilaes ol (05 by a factor of ~20., The new final abundances of $^+$ (solid line) is now smaller than earlier estimates of GP98 by a factor of $\sim 20$.
+ The scelarlo elmereing [rom the above caletla101s therefore indicates that in the regious of redshift below z~30. LiH. remais the more abLLClaut species compared to . but ouly by a factor of  2-7.," The scenario emerging from the above calculations therefore indicates that in the regions of redshift below $z\simeq 30$, LiH remains the more abundant species compared to $^+$, but only by a factor of $\sim 2$ –7."
+ The two molecuar abundanuces. on tie other hand. reach the largest. values lor <10. remaining both fairly sina| (LiH~10 PF. LiH~10 48) and hard to detect.," The two molecular abundances, on the other hand, reach the largest values for $z\lesssim 10$, remaining both fairly small $\sim$ $^{-17}$, $^+ \sim$ $^{-18}$ ) and hard to detect."
+ Buikliug upon recent quailtun reactive calculations (Bovino et al., Building upon recent quantum reactive calculations (Bovino et al.
+ 2009. 2010a. 2010b) 1ivolviug the chemi‘al evolution of molectles within the expected couclitious iu tlie early iullverse. the presett work has revisited the analysis of all the cyuaimical processes that a'e known to significantly contribute tothe »oductior/destruction of the lithitun-containing molecles.," 2009, 2010a, 2010b) involving the chemical evolution of $^+$ molecules within the expected conditions in the early universe, the present work has revisited the analysis of all the dynamical processes that are known to significantly contribute to the production/destruction of the lithium-containing molecules."
+ We have euniployed as many results as yossible from calculations based on ab-iuitio methocs. both [or the interaction forces aud the quautuuim dynamics. 'esortit& to estimates only for a few of e considered JJOCOesSsSes.," We have employed as many results as possible from calculations based on ab-initio methods, both for the interaction forces and the quantum dynamics, resorting to estimates only for a few of the considered processes."
+ One of the main results from the quautuLL leacive calculations is that two of e important ‘ates exhibit a teniperature dependence that was nol present iu tlie earlier estimates., One of the main results from the quantum reactive calculations is that two of the important rates exhibit a temperature dependence that was not present in the earlier estimates.
+ ΤΙe results for tle abunance of LiH incicale that this nolecule is now much more likely to lave slrvived at LE»w redshift. since its factio1al abtndaauce. albeit still fairly small. goes up by a actor of TO coiipa'ed to preVIOUS esnuates.," The results for the abundance of LiH indicate that this molecule is now much more likely to have survived at low redshift, since its fractional abundance, albeit still fairly small, goes up by a factor of $\sim 70$ compared to previous estimates."
+ However. lt jecomes sinaller by 1early the same amount for 2>300.," However, it becomes smaller by nearly the same amount for $z>300$."
+ The fractiona abuxlauce of οuly becomes significai (in the low-redshift region of 2z30 and shows a reduced value of tle relative abuncance by about one orce ‘of magnitude with respect to the earlier estinates., The fractional abundance of $^+$ only becomes significant in the low-redshift region of $z \lesssim 30$ and shows a reduced value of the relative abundance by about one order of magnitude with respect to the earlier estimates.
+ The «'omparison betwee1 the specili€ [ractioual abundauces of the two species now Indicaes that the 1eutral molecule is ikely to betjore abuidant than the ionic species aud that their rati yin the regio1 of small redshift iicleases Up O n [actor of 7 downto:~1., The comparison between the specific fractional abundances of the two species now indicates that the neutral molecule is likely to be more abundant than the ionic species and that their ratio in the region of small redshift increases up to a factor of 7 down to $z\simeq 1$.
+ Unlike prevlotus estimates that preclicted a difference at ow-2 of aboul two o‘ers of magnitude iu favor of . the new calculations iud more comparable abundaices of LiH aud .," Unlike previous estimates that predicted a difference at $z$ of about two orders of magnitude in favor of $^+$, the new calculations find more comparable abundances of LiH and $^+$."
+ Furthermore. we [iud that the neutral molectile. in spiteof a la‘ee dilution at the low redshifts. should be more amenable to experimental obse‘vation fe... Persson et al.," Furthermore, we find that the neutral molecule, in spiteof a large dilution at the low redshifts, should be more amenable to experimental observation (e.g., Persson et al."
+ 2010) than its ionic couuterpart., 2010) than its ionic counterpart.
+The possibility of forming elobular clusters from collisions between gas rich bodies was considered by Gunn (1980) ancl independently by MeCrea (1982).,The possibility of forming globular clusters from collisions between gas rich bodies was considered by Gunn (1980) and independently by McCrea (1982).
+ We revive and extend the Gunn and MeCrea analvsis here., We revive and extend the Gunn and McCrea analysis here.
+" Within each of (wo equal mass interacting galaxies consider an undisturbed column of eas of length I. cross-sectional area. A and density p, meeting its counter- part with relative velocity 2V (see Fie.1)."," Within each of two equal mass interacting galaxies consider an undisturbed column of gas of length l, cross-sectional area A and density $\rho_{o}$ meeting its counter- part with relative velocity 2V (see Fig.1)."
+ During the collision. the gas is compressed to column length A and density p.," During the collision, the gas is compressed to column length $\lambda$ and density $\rho$."
+ The shocked gas is assumed to cool to Tc10! degrees., The shocked gas is assumed to cool to $\sim10^{4}$ degrees.
+" This assumption will be justified later-suffice il to sav that the cooling (ime within the compressed. volume (7,4) must be shorter (han the collision time τρ~ //V) for the assumption to be valid.", This assumption will be justified later-suffice it to say that the cooling time within the compressed volume $\tau_{cool}$ ) must be shorter than the collision time $\tau_{coll}\sim l/V$ ) for the assumption to be valid.
+" By conservation of mass and for simplicity assuming A to remain constant equivalently in terms of the column density ancl with the relative velocity. being 2V. the relation between p, and p becomes (e.g. Spitzer (LOTS). equation 10-24) so (hal where the bracketed quantity in equation (4) represents (tlie square of the sound speed. and in equation (5) the unit of V is km/sec. T; is the temperature in units of 10! degrees and 115 the mean molecular weight of the shocked gas."," By conservation of mass and for simplicity assuming A to remain constant equivalently in terms of the column density and with the relative velocity being 2V, the relation between $\rho_{o}$ and $\rho$ becomes (e.g. Spitzer (1978), equation 10-24) so that where the bracketed quantity in equation (4) represents the square of the sound speed, and in equation (5) the unit of V is km/sec, $_{4}$ is the temperature in units of $10^{4}$ degrees and $\mu$ is the mean molecular weight of the shocked gas."
+ The quantity Εμ) occurs IrequentIy in, The quantity $_{4}/\mu$ ) occurs frequently in
+decrease with time. instead. of remaining constant.,"decrease with time, instead of remaining constant."
+ In our opinion. this poses considerable problems for the explanation of multiple populations with high helium abundances in GCs using intermediate-mass stars.," In our opinion, this poses considerable problems for the explanation of multiple populations with high helium abundances in GCs using intermediate-mass stars."
+ In the scenario. the FG stars are born in an intracluster medium previously enriched in. metals GC metallicity). while after some Myr the remaining gas is completely expelled by core-collapse SNe.," In the scenario, the FG stars are born in an intracluster medium previously enriched in metals GC metallicity), while after some Myr the remaining gas is completely expelled by core-collapse SNe."
+ The GC is then ready to start the accretion of the intermediate-mass AGB ejecta for several x10? yr., The GC is then ready to start the accretion of the intermediate-mass AGB ejecta for several $\times 10^8$ yr.
+ During this time. the GC is also accreting pristine material from the ambient interstellar medium (ISM). which is mixed with the AGB ejecta to form SG stars.," During this time, the GC is also accreting pristine material from the ambient interstellar medium (ISM), which is mixed with the AGB ejecta to form SG stars."
+ Finally. SG core-collapse SNe. and later FG type Ia SNe. begin to explode. keeping the GC gas-free. which stops the star formation process.," Finally, SG core-collapse SNe, and later FG type Ia SNe, begin to explode, keeping the GC gas-free, which stops the star formation process."
+ Even though the accretion of material from the ISM is plausible. it is unlikely that the chemical composition of this accreted material is similar to the GC itself.," Even though the accretion of material from the ISM is plausible, it is unlikely that the chemical composition of this accreted material is similar to the GC itself."
+ Moreover. if core-collapse SN explosions cleaned the intracluster medium. it must also clean the surrounding ISM or/and increase its metallicity. implying that SG stars will be formed with different metal abundances.," Moreover, if core-collapse SN explosions cleaned the intracluster medium, it must also clean the surrounding ISM or/and increase its metallicity, implying that SG stars will be formed with different metal abundances."
+ To close. we note that have very recently also argued against this scenario. pointing out in particular that if 1s not possible to form very O-poor stars through the process envisaged by ?.," To close, we note that have very recently also argued against this scenario, pointing out in particular that it is not possible to form very O-poor stars through the process envisaged by ."
+. also point out. following?.. that the scenario also fails to account for the high amount of He that is needed to explain the highly He-enriched populations that appear to be present in some GCs (see refintro)).," also point out, following, that the scenario also fails to account for the high amount of He that is needed to explain the highly He-enriched populations that appear to be present in some GCs (see \\ref{intro}) )."
+ have made a remarkable effort to obtain a homogeneous spectroscopic database for 17 GCs. where they have focused mainly on the O-Na anticorrelation.," have made a remarkable effort to obtain a homogeneous spectroscopic database for 17 GCs, where they have focused mainly on the O-Na anticorrelation."
+" These data were used to divide the stellar populations in GCs into three components: 1) The primordial population. with stars with O and Na abundances similar to field stars (1.e.. stars with [Na/Fe]«|[Na/Fe],,;,+0.3. where [Na/Fe],,;, is the lowest Na abundance detected in that GC): 11) The intermediate population. with stars with [O/Na]>-θ.9 that do not belong to the primordial population. amounting to to of the observed stars; and in) The extreme population. with stars with [O/Na]«-0.9. which are not presented in all GCs."," These data were used to divide the stellar populations in GCs into three components: i) The primordial population, with stars with O and Na abundances similar to field stars (i.e., stars with ${\rm [Na/Fe]} < {\rm [Na/Fe]}_{\rm min} + 0.3$, where ${\rm [Na/Fe]}_{\rm min}$ is the lowest Na abundance detected in that GC); ii) The intermediate population, with stars with ${\rm [O/Na]} > -0.9$ that do not belong to the primordial population, amounting to to of the observed stars; and iii) The extreme population, with stars with ${\rm [O/Na]} < -0.9$, which are not presented in all GCs."
+ Using these data and other global parameters characterizing the GCs in their sample. have proposed a scenario where GCs were formed in three different stages.," Using these data and other global parameters characterizing the GCs in their sample, have proposed a scenario where GCs were formed in three different stages."
+ First. à precursor population of stars was formed when the unborn GC (with a size of ~100 pe. and made of gas and dark matter) interacts strongly with other structures.," First, a precursor population of stars was formed when the unborn GC (with a size of $\sim 100$ pc, and made of gas and dark matter) interacts strongly with other structures."
+ Core-collapse SNe of this population enrich the remaining gas. and trigger the formation of the primordial population.," Core-collapse SNe of this population enrich the remaining gas, and trigger the formation of the primordial population."
+ Then. the gas ejected by primordial FRMS or super-AGB stars give rise to a gas cloud chemically enriched in the center. where the second generation (SG) of stars is born.," Then, the gas ejected by primordial FRMS or super-AGB stars give rise to a gas cloud chemically enriched in the center, where the second generation (SG) of stars is born."
+ Finally. SG core-collapse SNe clean the remaining gas. thus halting star formation.," Finally, SG core-collapse SNe clean the remaining gas, thus halting star formation."
+ During this time. the structure has lost all its dark matter content. almost all the precursor stars. and a large fraction of the primordial population.," During this time, the structure has lost all its dark matter content, almost all the precursor stars, and a large fraction of the primordial population."
+ While this appears to provide a promising framework. there are some points which deserve a deeper inspection.," While this appears to provide a promising framework, there are some points which deserve a deeper inspection."
+ Primordial low-mass stars. which are not observed in the actual GC because they were expelled from the initial structure. must still be present in some place.," Primordial low-mass stars, which are not observed in the actual GC because they were expelled from the initial structure, must still be present in some place."
+ Therefore. according to this scenario. a large number of metal-poor stars must be present in the field.," Therefore, according to this scenario, a large number of metal-poor stars must be present in the field."
+" However. to date there are only 174 known stars with [Fe/H]€—3. and 659 stars with[Fe/H]€-|Suda,ral2008."," However, to date there are only 174 known stars with ${\rm [Fe/H]} \le -3$, and 659 stars with${\rm [Fe/H]} \le -2$."
+ Alternatively. thei heavyatlowmetallicitiesC!?).," Alternatively, the initial mass function (IMF) must be top-heavy at low metallicities."
+.Note. inaddition. thati f thegase jectedbythe] collapseSNeisretainedintheinitialstructure. thee jectao f the primordialee collapseSNemustalsoberetained. whichwouldincreasethemetallicttyo fth ," Note, in addition, that if the gas ejected by the precursor core-collapse SNe is retained in the initial structure, the ejecta of the primordial core-collapse SNe must also be retained, which would increase the metallicity of the SG stars unless the initial structure had already lost a large fraction of its mass in a short period of time."
+have shown that. in order to reproduce the observed O- anticorrelation. as well as the suggested He spread in NGC 2808. SG stars must form from super-AGB ejectaonly. whereas in the case of M4 (NGC 6121) SG stars must be formed with a large amount of pristine gas mixed in.," have shown that, in order to reproduce the observed O-Na anticorrelation, as well as the suggested He spread in NGC 2808, SG stars must form from super-AGB ejecta, whereas in the case of M4 (NGC 6121) SG stars must be formed with a large amount of pristine gas mixed in."
+ This is very surprising. since NGC 2808 is ten times more massive than M4 (1.8x10? M). whereas it should be easier for a massive PS than for a low-mass PS to retain pristine gas.," This is very surprising, since NGC 2808 is ten times more massive than M4 $1.8\times 10^5 M_\odot $ ), whereas it should be easier for a massive PS than for a low-mass PS to retain pristine gas."
+ In this case. M4 would have lost many more stars since its formation than did NGC 2808. implying a more massive PS in the past — but not so massive as to have retained any (unobserved) trace of metals from SN explosions.," In this case, M4 would have lost many more stars since its formation than did NGC 2808, implying a more massive PS in the past – but not so massive as to have retained any (unobserved) trace of metals from SN explosions."
+ However. due to their deeper potential wells. massive PSs are expected to lose a smaller number of stars than less massive PSs.," However, due to their deeper potential wells, massive PSs are expected to lose a smaller number of stars than less massive PSs."
+ This is contrary to what would have been expected tf M4 were related with a massive PS but NGC 2808 with a less massive one. unless the M4 progenitor has been affected by stronger dynamical interactions in the course of its lifetime.," This is contrary to what would have been expected if M4 were related with a massive PS but NGC 2808 with a less massive one, unless the M4 progenitor has been affected by stronger dynamical interactions in the course of its lifetime."
+ Finally. we believe it is important to emphasize that. 1 our scenario. super-AGB and/or normal AGB stars play an important role in the formation of multiple populations.," Finally, we believe it is important to emphasize that, in our scenario, super-AGB and/or normal AGB stars play an important role in the formation of multiple populations."
+" It fact. as we shall see. while our scenario implies that these stars cannot be the progenitors of the most helium-rich populatior observed in NGC 2808 and w Cen. it does also indicate that they are likely to be the progenitors of the stars with helium abundances falling in between the ""normal"" (1.e.. and most He-rich stars."," In fact, as we shall see, while our scenario implies that these stars cannot be the progenitors of the most helium-rich population observed in NGC 2808 and $\omega$ Cen, it does also indicate that they are likely to be the progenitors of the stars with helium abundances falling in between the “normal” (i.e., non-He-enriched) and most He-rich stars."
+ Our new scenario ts divided in two stages. and is schematically described in Figs. 1.. 3..a," Our new scenario is divided in two stages, and is schematically described in Figs. \ref{FIGNonMassiveGC}, , \ref{FIGFairlyMassiveGC}, ,"
+nd 5)).,and \ref{FIGMassiveGC}) ).
+ The first stage gives the basis, The first stage gives the basis
+Or a variety of models under case Bo recombination and optically thin NLR gas.,for a variety of models under case B recombination and optically thin NLR gas.
+ Narrow line. ratios can oe reddening corrected. using the Balmer cecrement as he indicator., Narrow line ratios can be reddening corrected using the Balmer decrement as the indicator.
+ Following Daldwin. Phillips Verlevich (1981). we find logO41/]5007//1:34861]—1.20 anc ogANIIJ658S3/Ha6562]=0.145 which. as expected. place his object in the AGN zone in line diagnostic diagrams (c.g. Osterbrock 1989. fig 12.1).," Following Baldwin, Phillips Terlevich (1981) we find $\log [[OIII]5007/H\beta
+4861]=1.20$ and $\log [[NII]6583/H\alpha 6562]=0.145$ which, as expected, place this object in the AGN zone in line diagnostic diagrams (e.g. Osterbrock 1989, fig 12.1)."
+ We have further usec he measured Balmer clecrement from the narrow. lines (Ilof1l3~ 6). to estimate a eas column cdensitv of Na(NLR)~31em7. assuming standard. gas-to-dus ratio (Bohlin et al 1978).," We have further used the measured Balmer decrement from the narrow lines $H\alpha/H\beta\sim
+6$ ), to estimate a gas column density of $N_{H}({\rm NLR})\sim
+3\times10^{21}\, {\rm cm}^{-2}$, assuming standard gas-to-dust ratio (Bohlin et al 1978)."
+ Since the OLLI] emission is likely to come from wel above the torus. the intensity. of the OLLJ5007 line. is supposed. to be an orientation-independent estimator. of the total AGN power.," Since the [OIII] emission is likely to come from well above the torus, the intensity of the [OIII]5007 line is supposed to be an orientation-independent estimator of the total AGN power."
+" Following Bassani ct al (1999) anc Pappa ct al (2001). who use the interstellar reddening law bv Savage Mathis (LO79). we estimate the unreddenec ΟΠΗ5007 Hux by correcting the observed one by a [actor (dHafllxinBret, "," Following Bassani et al (1999) and Pappa et al (2001), who use the interstellar reddening law by Savage Mathis (1979), we estimate the unreddened [OIII]5007 flux by correcting the observed one by a factor $[(H\alpha/H\beta)_{\rm NLR}/3]^{2.94}$."
+The resulting OLUJ5007 Dux. is ~134.10ereem7s+ When a similar analysis is performed in. the Broad Line Region (19119). a Balmer decrement (faνι~27 is found.," The resulting [OIII]5007 flux is $\sim 1.34\times 10^{-12}\, {\rm erg}\, {\rm cm}^{-2}\, {\rm
+s}^{-1}$ When a similar analysis is performed in the Broad Line Region (BLR), a Balmer decrement $(H\alpha/H\beta)_{\rm BLR}\sim 27$ is found."
+ This should be considered. as a lower limit. as the existence of an LL? broad. component (1.0. the 15 Seyfert character of 1113201551) can only be established via constrained. parameter fitting.," This should be considered as a lower limit, as the existence of an $\beta$ broad component (i.e. the 1.8 Seyfert character of H1320+551) can only be established via constrained parameter fitting."
+ Wo this Balmer decrement is interpreted in terms of reddening over a standard. value of 3. a value of EDΕpipe~neg is found. which for a stancard dust to gas ratio corresponds to a LE column density of Ng(BLR)~1077.," If this Balmer decrement is interpreted in terms of reddening over a standard value of 3, a value of $E(B-V)_{\rm BLR}\sim 2^{mag}$ is found, which for a standard dust to gas ratio corresponds to a H column density of $N_{H}({\rm BLR})\sim 10^{22}$."
+ In the framework of the AGN unilie model the cillerence between optical spectroscopic Sevíer types is due to an orientation elfect. which results in. both reddening of the BLIt and absorption of N-ravs., In the framework of the AGN unified model the difference between optical spectroscopic Seyfert types is due to an orientation effect which results in both reddening of the BLR and absorption of X-rays.
+ Under these circumstances the above absorption column density shoulc be seen in the X-ray data., Under these circumstances the above absorption column density should be seen in the X-ray data.
+ We have also tried. to fit the broad. band optica spectra obtained. here with a model mixing a reddenec QSO. from the Francis et al (1991) template and a I5/50 galaxy template. from. the Coleman. Wu Weedman (1980) model.," We have also tried to fit the broad band optical spectra obtained here with a model mixing a reddened QSO, from the Francis et al (1991) template and a E/S0 galaxy template, from the Coleman, Wu Weedman (1980) model."
+ A good simultaneous description of the blue and red spectrum is achieved with the QSO template. which can accomodate a very small amount of reccdening (ECD1)« 01777) but it certainly needs some host ealaxv light.," A good simultaneous description of the blue and red spectrum is achieved with the QSO template, which can accomodate a very small amount of reddening $E(B-V)<0.1^{mag}$ ) but it certainly needs some host galaxy light."
+ The latter contributes about 40-60. per cent of the optical spectrum at. our reference. point at 5550A., The latter contributes about 40-60 per cent of the optical spectrum at our reference point at 5550.
+ Significantly larger reddening over the Francis et al (1991) QSO template simply does not match the data., Significantly larger reddening over the Francis et al (1991) QSO template simply does not match the data.
+ That suggests that whatever causes the large BLIt Balmer decrement does not appear to be nuclear reddening., That suggests that whatever causes the large BLR Balmer decrement does not appear to be nuclear reddening.
+ As already mentioned. X-ray emission from 111320|551 was discovered in the AIC-LASS survey. which assigned it a count rate of (4.1+0.8)105 et/s. Observations were carried out curing 1977.," As already mentioned, X-ray emission from H1320+551 was discovered in the MC-LASS survey, which assigned it a count rate of $(4.1\pm
+0.8)\times 10^{-3}$ ct/s. Observations were carried out during 1977."
+" Ceballos Barcons (1996) assumed a b=L.7 power law spectrum and computed a 2-10 keV. Dux of 2.0189Hoorecmmst, 1"," Ceballos Barcons (1996) assumed a $\Gamma=1.7$ power law spectrum and computed a 2-10 keV flux of $2.0\times 10^{-11}\, {\rm erg}\, {\rm
+cm}^{-2}\, {\rm s}^{-1}$."
+113201551. was detected in the ROSA all-sky survey (1990/91) as source TRNAS J132248.5|545526 (Schwope οἱ al 2000) with a PSPC count rate of 0.23+0.024 ctἐς. and a Hardness Ratio where (0.10.4) and 60.52.0) are the PSPC counts collected in the 0.1-0.4. keV and 0.5-2.0 keV. bands respectively.," H1320+551 was detected in the $ROSAT$ all-sky survey (1990/91) as source 1RXS J132248.5+545526 (Schwope et al 2000) with a PSPC count rate of $0.23\pm 0.024$ ct/s, and a Hardness Ratio where $C(0.1-0.4)$ and $C(0.5-2.0)$ are the PSPC counts collected in the 0.1-0.4 keV and 0.5-2.0 keV bands respectively."
+" The 0.5-2.0 keV absorption-corrected [ux was computed in Ceballos Barcons (1996) by assuming a b=1.9 power law spectrum with galactic absorption. resulting in (1.9+0.2).1077erecmT.d, "," The 0.5-2.0 keV absorption-corrected flux was computed in Ceballos Barcons (1996) by assuming a $\Gamma=1.9$ power law spectrum with galactic absorption, resulting in $(1.9\pm 0.2)\times 10^{-12}\, {\rm erg}\,
+{\rm cm}^{-2}\, {\rm s}^{-1}$."
+ASCUL observed 1113201551 on May. LO of 1999. for. 10 ks., $ASCA$ observed H1320+551 on May 10 of 1999 for 10 ks.
+ We have retrieved the pipeline-reclucecl data from the HEASARC public archive., We have retrieved the pipeline-reduced data from the HEASARC public archive.
+ Phe SISO anc SISI data over the 0.5-S keV hand can be well fitted with a single power law with galactic absorption. resulting in a 47=17.4 [or 26 degrees of Ereedom (see fig.," The SIS0 and SIS1 data over the 0.5-8 keV band can be well fitted with a single power law with galactic absorption, resulting in a $\chi^2=17.4$ for 26 degrees of freedom (see fig."
+ 3. for the 49€ spectrum)., \ref{fig-ascaspec} for the $ASCA$ spectrum).
+ The data do not require additional absorption or à soft excess., The data do not require additional absorption or a soft excess.
+ Fie 4 shows the confidence contours for the power-law photon index (P= 1.53+0.1) and the normalisation required bv the ASCUL cata., Fig \ref{fig-cont-asca} shows the confidence contours for the power-law photon index $\Gamma=1.53\pm0.1$ ) and the normalisation required by the $ASCA$ data.
+ We have also overlayed lines with various 2-10 keV flux levels. for comparison with the NMM-Newton observations.," We have also overlayed lines with various 2-10 keV flux levels, for comparison with the XMM-Newton observations."
+ Indeed the ASCUL 2-10 keV [lux is 10 times, Indeed the $ASCA$ 2-10 keV flux is 10 times
+"absorption. while only of sources with 10!L.<Ly,1017?Lo. and of sources with LOMO""Lo<LyLOM?L. show absorption.","absorption, while only of sources with $10^{11.00}~L_{\odot}
+\leq L_{\rm IR} \leq 10^{11.49}~L_{\odot}$, and of sources with $10^{10.50}~L_{\odot} \leq L_{\rm IR} \leq 10^{10.99}~L_{\odot}$ show absorption."
+" No absorption is seen in sources with Ly,<1047L. (see Table 8a)."," No absorption is seen in sources with $L_{\rm IR}
+\leq 10^{10.49}~L_{\odot}$ (see Table 8a)."
+" Thus. while we cannot sav anything about traditionally defined ULIRGs (Ly,>107L,) because of their rarity in our observed sample. it appears that galaxies with Lig>LOMO!L, have a greater likelihood of showing Labsorption."," Thus, while we cannot say anything about traditionally defined ULIRGs $L_{\rm IR} \geq
+10^{12}~L_{\odot}$ ) because of their rarity in our observed sample, it appears that galaxies with $L_{\rm IR} \geq 10^{11.50}~L_\odot$ have a greater likelihood of showing absorption."
+ We now explore whether this trend could arise due to a selection effect., We now explore whether this trend could arise due to a selection effect.
+ While the flux densityof emission is proportional to Liyj/D?. where D is the distance to (he galaxy. (he flux density clip of an absorption corresponding (o a given optical depth is proportional to the background continuum flux density (hat is being absorbed.," While the flux densityof emission is proportional to $L_{\rm
+HI}/D^2$, where $D$ is the distance to the galaxy, the flux density dip of an absorption corresponding to a given optical depth is proportional to the background continuum flux density that is being absorbed."
+" la Table 3b we present the calculated mean and median NVSS flux densities for (he galaxies in our observed sample as a [function of Lii bins. to investigate whether galaxies having Ly,>LOMO?L, preferentially have higher flux densities."," In Table 8b we present the calculated mean and median NVSS flux densities for the galaxies in our observed sample as a function of $L_{\rm IR}$ bins, to investigate whether galaxies having $L_{\rm IR}\geq 10^{11.50}~L_{\odot}$ preferentially have higher flux densities."
+ Table 8b shows that the radio [lux densities in our sample do not correlate with LB. huminositv., Table 8b shows that the radio flux densities in our sample do not correlate with IR luminosity.
+ However. we do find for those galaxies showing absorption in the two highest luminosity bins. that the mean flux densities are about double the mean values for all galaxies in those bins.," However, we do find for those galaxies showing absorption in the two highest luminosity bins, that the mean flux densities are about double the mean values for all galaxies in those bins."
+" As a second possible effect. the Ly,>1017""Lo absorbers might have a higher covering factor than those of lower luminosity e.g.. their contünuun enission may mostly be in a compact nucleus vs. more extended continuum emission for the others. with the absorption arising principally in the nuclear regions."," As a second possible effect, the $L_{\rm IR} \geq 10^{11.50}~L_{\odot}$ absorbers might have a higher covering factor than those of lower luminosity e.g., their continuum emission may mostly be in a compact nucleus vs. more extended continuum emission for the others, with the absorption arising principally in the nuclear regions."
+" llowever. whether this is (he case or not. we do find a difference in the incidence of absorption between IR galaxies wilh Lig>10!""L. and those with lower IR luminosities. i.e.. (he galaxies with higher IR. luminosities have higher column densities on the lines of sieht to the continuum sources. and/or their continuum enission is confined lo more conipacl regions."," However, whether this is the case or not, we do find a difference in the incidence of absorption between IR galaxies with $L_{\rm IR} \geq 10^{11.50}~L_{\odot}$ and those with lower IR luminosities, i.e., the galaxies with higher IR luminosities have higher column densities on the lines of sight to the continuum sources, and/or their continuum emission is confined to more compact regions."
+" In Table 9. we present the statistics of OII detections (absorption or emission) as a funelion of Ly, bins."," In Table 9, we present the statistics of OH detections (absorption or emission) as a function of $L_{\rm IR}$ bins."
+ The total number of sources per Lii bin in this table excludes galaxies with OI] main line spectra severely affected by REI., The total number of sources per $L_{\rm IR}$ bin in this table excludes galaxies with OH main line spectra severely affected by RFI.
+" Here we note that all 3 of the OII emitters have Li,>LOMOL.. while the two detections for LOM""!Lo<Lip101799p. are both absorbers."," Here we note that all 3 of the OH emitters have $L_{\rm IR} \geq 10^{11.50}~L_{\odot}$, while the two detections for $10^{10.50}~L_{\odot} \leq L_{\rm IR} \leq
+10^{10.99}~L_{\odot}$ are both absorbers."
+ ILowever. despite the small iunmbers in Chis suberoup. il is stronely suggestive (hal Oll-detected sources (emitters and absorbers) are found primarily in galaxies with Ligος.," However, despite the small numbers in this subgroup, it is strongly suggestive that OH-detected sources (emitters and absorbers) are found primarily in galaxies with $L_{\rm IR} \geq 10^{11.50}~L_{\odot}$."
+ This galaxy is at a recdshift of 2=0.0564. ancl is one of two sources with no detection in our observed sample. with the other source being IRAS 23410+0228 (2= 0.0912: see below).," This galaxy is at a redshift of $z=0.0564$, and is one of two sources with no detection in our observed sample, with the other source being IRAS 23410+0228 $z=0.0912$ ; see below)."
+ Its IR huninosity is 10199779£L. and the derived 30 upper limit for, Its IR luminosity is $10^{11.90 \pm 0.09}~L_{\odot}$ and the derived $\sigma$ upper limit for
+We find that the stellar density drops olf much more quickly than estimated by ?.. with approximately one third as many stars detected by us in the 1. 2.5 M. mass range (where our data remain complete) at distances of 215'.. compared with the ing profile determined bv ?..,"We find that the stellar density drops off much more quickly than estimated by \citet{knod00}, with approximately one third as many stars detected by us in the 1 – 2.5 $_{\odot}$ mass range (where our data remain complete) at distances of $>$, compared with the King profile determined by \citet{knod00}."
+ Using our radial density. profile and assuming that Cvg OB? is spherically svuunetric we have estimated a new Nine prolile for Cvg OD2., Using our radial density profile and assuming that Cyg OB2 is spherically symmetric we have estimated a new King profile for Cyg OB2.
+ Using our peak stellar density and the mass function slope estimated above. extrapolated down to 0.01 AL. using (he multi-stage power-law LMF from? we calculate a total stellar content of5.7x104 stars and a total mass of 2.8x10! M...," Using our peak stellar density and the mass function slope estimated above, extrapolated down to 0.01 $_{\odot}$ using the multi-stage power-law IMF from \citet{krou01a} we calculate a total stellar content of$5.7 \times 10^4$ stars and a total mass of $2.8 \times 10^4$ $_{\odot}$."
+ However it is clear from the spatial distribution of A-tvpe dwarfs in Cve OB? presented by ? (hat the structure of (νο OB? is Far from spherical., However it is clear from the spatial distribution of A-type dwarfs in Cyg OB2 presented by \citet{drew08} that the structure of Cyg OB2 is far from spherical.
+ The authors note a distinct lack of A-twpe dwarls to the east of Cvg OD2 and a pronounced concentration towards (he southern part of the cluster (hat mar increase (he size of the association., The authors note a distinct lack of A-type dwarfs to the east of Cyg OB2 and a pronounced concentration towards the southern part of the cluster that may increase the size of the association.
+ We therefore estimate a total stellar mass of (341)x10! M. for the entire association., We therefore estimate a total stellar mass of $(3 \pm 1) \times 10^4$ $_{\odot}$ for the entire association.
+ Usine (his estimated densitv profile aud our derived AIF we can estimate the expected nunber of stars of different masses in Cvg ΟΡΟ., Using this estimated density profile and our derived MF we can estimate the expected number of stars of different masses in Cyg OB2.
+ A total mass of 3x10! M. implies a total of 71200 OB stars. approximately half that predicted by ?..," A total mass of $3 \times 10^4$ $_{\odot}$ implies a total of $\sim$ 1200 OB stars, approximately half that predicted by \citet{knod00}."
+ Our current. ME predicts a total number of O-(vpe stars of 15. only slightly larger (han the currently known munber of 65 ," Our current MF predicts a total number of O-type stars of $\sim$ 75, only slightly larger than the currently known number of 65 O-type stars in Cyg OB2 \citep{mass91,hans03,kimi07,come08,negu08}."
+be discovered in Cyg OD2., This may suggest that there are more O-type stars to be discovered in Cyg OB2.
+ If. we assume that the high-mass MIF slope has been influenced by the demise of the most massive stars in the region and that it was originally [=—1.3 (?).. we estimate the number of stars that have evolved (o their end states to be ~12.," If we assume that the high-mass MF slope has been influenced by the demise of the most massive stars in the region and that it was originally $\Gamma = -1.3$ \citep{mass98}, we estimate the number of stars that have evolved to their end states to be $\sim$ 12."
+" There are no known supernova remnants in Cvg ΟΡΟ, although there is some evidence for stellar remnants: ? and ? identify (wo pulsars within Cvg ΟΡΟ, with ages ~2x10! vr (2).. that likely represent recently expired massive stars. while ? identified an O4IE runaway star (hat appears to have been ejected Lom (νο ΟΡΟ ~1.6 Myr ago. possibly following the supernova explosion of a more massive binary companion."," There are no known supernova remnants in Cyg OB2, although there is some evidence for stellar remnants: \citet{bedn03} and \citet{abdo09} identify two pulsars within Cyg OB2, with ages $\sim 2 \times 10^4$ yr \citep{cami09}, that likely represent recently expired massive stars, while \citet{come07} identified an O4If runaway star that appears to have been ejected from Cyg OB2 $\sim$ 1.6 Myr ago, possibly following the supernova explosion of a more massive binary companion."
+ Table 2. lists some of the properties of Cve OB? estimated in (his work along with those of a number of other massive star forming regions., Table \ref{msfrs} lists some of the properties of Cyg OB2 estimated in this work along with those of a number of other massive star forming regions.
+ With the exception of the ONC (he majority of these regions have not been studied to as low a stellar mass. wilh properties determined in different wavs by different authors (the number of O-tvpe stars is particularly poorly known for many of these regions).," With the exception of the ONC the majority of these regions have not been studied to as low a stellar mass, with properties determined in different ways by different authors (the number of O-type stars is particularly poorly known for many of these regions)."
+ Despite this some comparisons may be nde., Despite this some comparisons may be made.
+ (νο OD2 is certainly a very high-mass cluster. comparable to the Arches and Quintuplet clusters near the Galactic center or the giant HE regions Westerlund 2 and NGC 3603 in the Carina spiral arm.," Cyg OB2 is certainly a very high-mass cluster, comparable to the Arches and Quintuplet clusters near the Galactic center or the giant H regions Westerlund 2 and NGC 3603 in the Carina spiral arm."
+ Onlv regions such as Westerlund | cluster. W49A. and R136 in the," Only regions such as Westerlund 1 cluster, W49A, and R136 in the"
+Wide-angle tail radio galaxies are an interesting sub-class of the population of FRI (Fanaroff Riley 1974)) objects. which typically lie. at the centres of clusters and show narrow. well-collimated jets which flare abruptly into broad. diffuse plumes.,"Wide-angle tail radio galaxies are an interesting sub-class of the population of FRI (Fanaroff Riley \cite{fr}) ) objects, which typically lie at the centres of clusters and show narrow, well-collimated jets which flare abruptly into broad, diffuse plumes."
+" Their radio power is normally intermediate between the more typical jet-dominated FRI and the “classical double"" FRII classes of extragalactic radio source. and an understanding of their dynamics ts important to our knowledge of the relationships between these two classes and the possible evolution between them."," Their radio power is normally intermediate between the more typical jet-dominated FRI and the `classical double' FRII classes of extragalactic radio source, and an understanding of their dynamics is important to our knowledge of the relationships between these two classes and the possible evolution between them."
+ In an earlier paper (Hardeastle 1998.. hereafter Paper D I presented new radio maps of the wide-angle tail radio galaxy (;= 0.109).," In an earlier paper (Hardcastle \cite{h98}, hereafter Paper I) I presented new radio maps of the wide-angle tail radio galaxy $z=0.109$ )."
+ Two-frequency spectral Index mapping in that paper showed flat-spectrum jets (1n this paper. the term is reserved for the narrow. well-collimated features seen in the inner 50 kpe of the sources) and a hotspot (a compact. sub-kpe feature at the end of the northern jet) together with steeper-spectrum material at the edges of the plumes (the broader. more diffuse features seen between 50 and ~500 kpe from the nucleus).," Two-frequency spectral index mapping in that paper showed flat-spectrum jets (in this paper, the term is reserved for the narrow, well-collimated features seen in the inner 50 kpc of the sources) and a hotspot (a compact, sub-kpc feature at the end of the northern jet) together with steeper-spectrum material at the edges of the plumes (the broader, more diffuse features seen between 50 and $\sim 500$ kpc from the nucleus)."
+ This steep-spectrum naterial is particularly clear in the southern plume. and is referred to here as a ‘sheath’. although | emphasise that. unlike the sheaths seen in some twin-jet FRI radio galaxies. this region has πο polarization properties to distinguish it from the rest of the plume.," This steep-spectrum material is particularly clear in the southern plume, and is referred to here as a `sheath', although I emphasise that, unlike the sheaths seen in some twin-jet FRI radio galaxies, this region has no polarization properties to distinguish it from the rest of the plume."
+ The sheath is the only feature of the two-point radio spectrum which is not obviously consistent with a fairly simple model for the source's dynamics. in which particles are accelerated at the base of the plumes and the spectral steepening along the plumes is à consequence of outflow of an ageing electron population.," The sheath is the only feature of the two-point radio spectrum which is not obviously consistent with a fairly simple model for the source's dynamics, in which particles are accelerated at the base of the plumes and the spectral steepening along the plumes is a consequence of outflow of an ageing electron population."
+ The additional data presented here show that the situation is more complicated than that simple model would imply., The additional data presented here show that the situation is more complicated than that simple model would imply.
+" B1950 co-ordinates are used throughout this paper. and spectral index o is defined in the sense Sxv"" "," B1950 co-ordinates are used throughout this paper, and spectral index $\alpha$ is defined in the sense $S \propto \nu^{-\alpha}$."
+The X-band (8.4-GHz) and L-band (1.4-GHz) data used in this paper were described in Paper I. Data from a short (0.5 h) C-band (4.9-GHz) observation with the NRAO Very Large Array (VLA) in its C configuration. taken on 1984 Jun 11. were kindly provided by Alan Bridle.," The X-band (8.4-GHz) and L-band (1.4-GHz) data used in this paper were described in Paper I. Data from a short (0.5 h) C-band (4.9-GHz) observation with the NRAO Very Large Array (VLA) in its C configuration, taken on 1984 Jun 11, were kindly provided by Alan Bridle."
+ The U-band (15-GHz) images presented here are the result of a 3-hour observation with the VLA in its D configuration taken on 1999 Mar 11., The U-band (15-GHz) images presented here are the result of a 3-hour observation with the VLA in its D configuration taken on 1999 Mar 11.
+ All data were reduced and analysed using the software package., All data were reduced and analysed using the software package.
+ As described in Paper L. both and were used as primary flux calibrators for the L- and X-band observations. which were taken between 1994 Nov IO and 1995 Nov 28.," As described in Paper I, both and were used as primary flux calibrators for the L- and X-band observations, which were taken between 1994 Nov 10 and 1995 Nov 28."
+ Specifically. was used for the A-configuration X-band observations and for all others.," Specifically, was used for the A-configuration X-band observations and for all others."
+ The data analysis in Paper | used the older version of the SETJY task. which incorrectly rounded coefficients in the analytic expression for the flux densities of calibrator sources.," The data analysis in Paper I used the older version of the SETJY task, which incorrectly rounded coefficients in the analytic expression for the flux densities of calibrator sources."
+ However. the effect is very small when combined with the change between the old (1990) values of the coefficients and the more appropriate new (1995.2) values.," However, the effect is very small when combined with the change between the old (1990) values of the coefficients and the more appropriate new (1995.2) values."
+ L-band fluxes should be reduced by and the X-band B. C and D-configuration fluxes by256: the A-contiguration X-band flux. based on 2286. should be reduced by when the effect of the partial resolution of the source is incorporated. but we use for all the X-band data in what follows.," L-band fluxes should be reduced by and the X-band B, C and D-configuration fluxes by; the A-configuration X-band flux, based on 286, should be reduced by when the effect of the partial resolution of the source is incorporated, but we use for all the X-band data in what follows."
+ The primary flux calibrator for the, The primary flux calibrator for the
+The Convection Rotation and planetary Transits (CoRoT) space telescope. launched in December 2006. has discovered 15 planets so far using the method of transits (see. e.g.. http://exoplanet.eu).,"The Convection Rotation and planetary Transits (CoRoT) space telescope, launched in December 2006, has discovered 15 planets so far using the method of transits (see, e.g., http://exoplanet.eu)."
+ The second planet discovered by CoRoT orbits a young and active star similar to the Sun. now named CoRoT-2. of spectral type G7V with mass and radius of 0.97 M... and 0.902 R... respectively.," The second planet discovered by CoRoT orbits a young and active star similar to the Sun, now named CoRoT-2, of spectral type G7V with mass and radius of 0.97 $M_{\odot}$ and 0.902 $R_{\odot}$, respectively."
+ CoRoT-2 was continuously monitored for about 135 days. during which 77 transits of its planet were observed with high temporal resolution (32 s).," CoRoT-2 was continuously monitored for about 135 days, during which 77 transits of its planet were observed with high temporal resolution (32 s)."
+ The first two transits were discarded because of the lower temporal resolution., The first two transits were discarded because of the lower temporal resolution.
+ The data were reduced as explainec in Alonsoetal.(2008) with a resulting rms of the points of the white light curve of 0.0006 in units of the relative flux of the out-of-transit data., The data were reduced as explained in \cite{alonso08} with a resulting rms of the points of the white light curve of 0.0006 in units of the relative flux of the out-of-transit data.
+ The CoRoT white passband ranges from 300 to 1100 nm and its transmission profile can be found in Auvergneetal.(2009).. together with a detailed description of the instrument. its operation. and the data processing pipeline.," The CoRoT white passband ranges from 300 to 1100 nm and its transmission profile can be found in \citet{auvergneetal09}, together with a detailed description of the instrument, its operation, and the data processing pipeline."
+ The basic assumptions and the parameters of the planet and its orbit used for the present modelling were taken from previous observations (Alonsoetal.2008).. and are listed below: The model used here simulates à star with quadratic. limb darkening as a 2-D image such as the one shown in the left panel of Figure 1..," The basic assumptions and the parameters of the planet and its orbit used for the present modelling were taken from previous observations \citep{alonso08}, and are listed below: The model used here simulates a star with quadratic limb darkening as a 2-D image such as the one shown in the left panel of Figure \ref{star}."
+ A detailed description of the model is given in Silva(2003)., A detailed description of the model is given in \cite{silva03}.
+. Applications of such à model are described in Silva-Valio(2008). and Silva-Valioetal.(2010).. the latter specifically for CoRoT-2.," Applications of such a model are described in \cite{silva-valio08} and \cite{silva-valio10}, the latter specifically for CoRoT-2."
+ The model simulates the passage of the planet. a completely dark disk. in front of the spotted star.," The model simulates the passage of the planet, a completely dark disk, in front of the spotted star."
+ The light curve is computed as the sum of the fluxes coming from all the pixels in the frame. for a sequence of planet positions along its orbit at intervals of 2 minutes.," The light curve is computed as the sum of the fluxes coming from all the pixels in the frame, for a sequence of planet positions along its orbit at intervals of 2 minutes."
+ Unless otherwise stated. the planet and star parameters are taken from Alonsoetal.(2008).," Unless otherwise stated, the planet and star parameters are taken from \cite{alonso08}."
+. Moreover. the model for the spots assumes that: This model was applied to all transit light curves and the three parameters (radius. intensity. and longitude) for each spot were sought simultaneously by minimizing the y between the model light curve and the data.," Moreover, the model for the spots assumes that: This model was applied to all transit light curves and the three parameters (radius, intensity, and longitude) for each spot were sought simultaneously by minimizing the $\chi^{2}$ between the model light curve and the data."
+ All fits were performed using the AMOEBA routine (Pressetal..1992)., All fits were performed using the AMOEBA routine \citep{press92}.
+. Initial guesses of the parameters were obtained through the genetic algorithm PIKAIA (Charbonneau.1995)., Initial guesses of the parameters were obtained through the genetic algorithm PIKAIA \citep{charbonneau95}.
+. The minimum number of spots per transit. M. needed to fit the light curve within the data uncertainty. was estimated as follows.," The minimum number of spots per transit, $\bar{M}$, needed to fit the light curve within the data uncertainty, was estimated as follows."
+ A model with M spots has 3M free parameters. because each spot corresponds to 3 free parameters (longitude. radius. and intensity).," A model with $M$ spots has $3 M$ free parameters, because each spot corresponds to 3 free parameters (longitude, radius, and intensity)."
+ Therefore. the number of degrees of freedom of the model is s=N—3M. where N ts the number of data points per transit (usually N= 217).," Therefore, the number of degrees of freedom of the model is $s=N-3 M$, where $N$ is the number of data points per transit (usually $N=217$ )."
+ Now. the problem is to find the minimum number of spots that yields an adequate best fit for each transit according to the criterion for hypothesis testing given by. e.g.. Lamptonetal.(1976).," Now, the problem is to find the minimum number of spots that yields an adequate best fit for each transit according to the criterion for hypothesis testing given by, e.g., \citet{lampton76}."
+. We compute the minimum value of the y for each model with a different M starting from the case with a single spot (M.= 1)., We compute the minimum value of the $\chi^{2}$ for each model with a different $M$ starting from the case with a single spot $M=1$ ).
+ A model, A model
+Recent observational ellorts to detect hieh-recshilt (2) supernovae (SNe) have demonstrated their value as cosmological probes.,Recent observational efforts to detect high-redshift $z$ ) supernovae (SNe) have demonstrated their value as cosmological probes.
+ The discovery ancl svstematic study of faint. distant. type Ia supernovac lla) has led. to renewed. progress in constraining the cosmic expansion history (Riessetal.1905:Perlmutter1999).," The discovery and systematic study of faint, distant type Ia supernovae Ia) has led to renewed progress in constraining the cosmic expansion history \cite{riess98,perlmutter99}."
+. Two independent corrected. Hubble. diagrams based on SNella discovered. via the Supernovac Cosmology. Project (Perlmutter et al., Two independent corrected Hubble diagrams based on Ia discovered via the Supernovae Cosmology Project (Perlmutter et al.
+ 1999) and the Lieh-Recshift Supernovae Search ‘Team (Riess et al., 1999) and the High-Redshift Supernovae Search Team (Riess et al.
+ 1998). vield a trend that strongly excludes the hitherto popular Einstein de Sitter universe and. for a spatially Hat inflationary universe consistent with recent microwave background measurements (Hancock1998:deBernarcisetal. 2000).. sugeests a significant cosmological constant. A0.7 (Perlmutteretal.1990:Riessetαἱ. 1998).," 1998) yield a trend that strongly excludes the hitherto popular Einstein de Sitter universe and, for a spatially flat inflationary universe consistent with recent microwave background measurements \cite{hancock98,debernardis00}, , suggests a significant cosmological constant, $\Lambda\simeq0.7$ \cite{perlmutter99,riess98}."
+. Such studies have now identified over 100 z0.2 La. and. using controlled subsets of these La samples. the irst constraints are now emerging on the rate of their occurrence. (Pain οἱ 11996. 20002).," Such studies have now identified over 100 $z>0.2$ Ia, and, using controlled subsets of these Ia samples, the first constraints are now emerging on the rate of their occurrence (Pain et 1996, 2000a)."
+ The observational «etermination of these rates is also important in many cosmological applications. [or example as a diagnostic of the cosmic star formation history (SELL) and metal enrichment at. high-z.," The observational determination of these rates is also important in many cosmological applications, for example as a diagnostic of the cosmic star formation history (SFH) and metal enrichment at $z$."
+ SNe are independent of some of the biases associated with traditional tests based. on Ilux-limited. galaxy samples., SNe are independent of some of the biases associated with traditional tests based on flux-limited galaxy samples.
+ Such biases are increasingly of concern given the steep luminosity functions now being witnessed for star-forming galaxies at all redshilts (Steideletal.1999:Sullivan2000) and. though there remain uncertainties in the correct treatment of dust in high-z type," Such biases are increasingly of concern given the steep luminosity functions now being witnessed for star-forming galaxies at all redshifts \cite{steidel99,sullivan00} and, though there remain uncertainties in the correct treatment of dust in $z$ type"
+support with many valuable numerical gadgets.,support with many valuable numerical gadgets.
+ This work has been supported by the Deutsche Forschungsgemeinschaft within the Priority Programme 1177 under the project PO 1454/1-1 and by NSF grant AST 07-08849., This work has been supported by the Deutsche Forschungsgemeinschaft within the Priority Programme 1177 under the project PO 1454/1-1 and by NSF grant AST 07-08849.
+Lwideuthe there are muy possible alterations and extensions that can be made to Phurbas that will sienificautly alter both the nature of the scheme aud the capabilities of the code.,"Evidently, there are many possible alterations and extensions that can be made to Phurbas that will significantly alter both the nature of the scheme and the capabilities of the code."
+ We believe that Phurbas is not just a new method for MIID. but one of the first practical exanrples of a new class of schemes for mathematical inodeliug of similar physical problems.," We believe that Phurbas is not just a new method for MHD, but one of the first practical examples of a new class of schemes for mathematical modeling of similar physical problems."
+" Simon C.O. Clover wrote the initial version of the routine to couple the Phurba ANID module toCADCET-2.. which lies at the core of Ῥ]ας,"," Simon C.O. Glover wrote the initial version of the routine to couple the Phurba MHD module to, which lies at the core of Phurbas."
+ The tests shown here are of the version of the code incorporating interpolating fits ax advised by the anonvinous referce., The tests shown here are of the version of the code incorporating interpolating fits as advised by the anonymous referee.
+ As noted in Paper L this was a contribution equivalent το co-authorship.," As noted in Paper I, this was a contribution equivalent to co-authorship."
+" We are indebted to Volker Spriusel for makine the source uininiziug a"" ⋅ ⋅ ≼⋅⋝⊀≷↱", We are indebted to Volker Springel for making the source code of publicly available \citep{2005MNRAS.364.1105S}.
+ also: thank the authors of Athena: for malsine⋅↜⋅ itiue publicly: available: (Stoue+etal.E2008)..1, We also thank the authors of Athena for making it publicly available \citep{2008ApJS..178..137S}.
+ AL-A\LALL. aud CoPAL acknowledge hospitality from the Max-Plauck-Tustitut finr Astronomic. and ΔΕΑΕΑΙ additionally acknowledges hospitality of the Institut far Theoretische Astroplivsik der Weidelhere.," M.-M.M.L. and C.P.M. acknowledge hospitality from the Max-Planck-Institut fürr Astronomie, and M.-M.M.L additionally acknowledges hospitality of the Institut fürr Theoretische Astrophysik der Heidelberg."
+ This work has been supported by National Science Foundation CDI eraut AST-083573| and allocation TC-AICA99S021. originally roni the Teragrid. aud now from the Extreme Scicuce and Eneimecring Discovery Euviroument (XSEDE). which is supported by National Science Foundation eraut nunber OCT-1053575.," This work has been supported by National Science Foundation CDI grant AST-0835734 and allocation TG-MCA99S024, originally from the Teragrid, and now from the Extreme Science and Engineering Discovery Environment (XSEDE), which is supported by National Science Foundation grant number OCI-1053575."
+ The test in refseciéiiacesisaconcergenectesttothesolutionofthelinearivedeersionoftheiodi FicdM II DequationsthatPh urbassolves, The test in \\ref{sec_linwaves} is a convergence test to the solution of the linearized version of the modified MHD equations that Phurbas solves.
+ W, We derive here the dispersion relation and solutions in the cases used in the convergence test.
+"M ede, added to the moment equation."," We start with the MHD mass, momentum, and induction equations, with a bulk viscosity term $\rho\nabla(\zeta \nabla\cdot\mathbf{V})$ added to the momentum equation,"
+Aneular momentum diffusion in well-ionized disks is likely driven by magnetohyvdrodynamie (MILD) turbulence.,Angular momentum diffusion in well-ionized disks is likely driven by magnetohydrodynamic (MHD) turbulence.
+ Analvtic analyses. numerical experiments. and laboratory evidence strongly suggest that well-coupled plasmas in differentiallv-rotating flows are subject to (he magnetorotational instabilitv (MIRI: Balbus&Hawley1991.1998: 2003)).," Analytic analyses, numerical experiments, and laboratory evidence strongly suggest that well-coupled plasmas in differentially-rotating flows are subject to the magnetorotational instability (MRI; \citealt{bh91,bh98,bal03}) )."
+ But AHID turbulence is initiated’ bv the MBRI only so long as the plasma is sulliciently ionized to couple to the magnetic field (Ixunz&Balbus2004:Deseh2004).," But MHD turbulence is initiated by the MRI only so long as the plasma is sufficiently ionized to couple to the magnetic field \citep{kb04,des04}."
+. In disks around voung stars. variable and X-ray binary disks in quiescence. ancl possibly the outer parts of AGN disks. the plasma may be (oo neutral (ο support magnetic activity (Ganunie&MenouL998; 2002).," In disks around young stars, cataclysmic-variable and X-ray binary disks in quiescence, and possibly the outer parts of AGN disks, the plasma may be too neutral to support magnetic activity \citep{gm98,men00,sgbh00,mq01,ftb02}."
+. This motivates interest in non-MIID angular momentum transport mechanisms., This motivates interest in non-MHD angular momentum transport mechanisms.
+ Within the last few vears. a body of work has been developed suggesting that vortices can be generated as a result of global hydrodynamic instability (Haley1987:Blaes&Hawley or local hydrodynamic instability (IXIahr&Boclenheimer2003).. that vortices in disks may be long-lived (Godon&Livio1999.2000;Unmurhan 2005).. and that these vortices may be related to an outward fIux of angular momentum (Li.2001:Barranco&Marcus 2005).," Within the last few years, a body of work has been developed suggesting that vortices can be generated as a result of global hydrodynamic instability \citep{haw87,bh88,haw90,love99,li00} or local hydrodynamic instability \citep{kb03}, that vortices in disks may be long-lived \citep{gl99,gl00,ur04,bm05}, and that these vortices may be related to an outward flux of angular momentum \citep{li01,bm05}."
+. IE these claims can be verified then the consequences lor low-ionization disks would be profound., If these claims can be verified then the consequences for low-ionization disks would be profound.
+ ]lere we investigatee the evolution of a disk that is egiven a largee initial vortical velocity perturbation., Here we investigate the evolution of a disk that is given a large initial vortical velocity perturbation.
+ Our study is done in the context of a (two-dimensional) shearing-sheet model. which permits us to resolve (he dynamics to a degree that is not currently possible in a global disk model.," Our study is done in the context of a (two-dimensional) shearing-sheet model, which permits us to resolve the dynamics to a degree that is not currently possible in a global disk model."
+ Our model is also filly compressible. unlike previous work using a local model (Umurhan&Reeey2004:BarrancoMarcus 2005)..," Our model is also fully compressible, unlike previous work using a local model \citep{ur04,bm05}. ."
+ The lormer assume incompressible flow and the latter use the anelastic approximation (e.g.. Gough 1969)). which filters out the hieh-frequency acoustic waves.," The former assume incompressible flow and the latter use the anelastic approximation (e.g., \citealt{gou69}) ), which filters out the high-frequency acoustic waves."
+ We will show that compressibility auc acoustic waves play an essential part in the angular momentum transport., We will show that compressibility and acoustic waves play an essential part in the angular momentum transport.
+ Our paper is organized as follows., Our paper is organized as follows.
+ In. 822 we describe (he model., In 2 we describe the model.
+ In. 833. we describe the evolution of a fiducial. high-resolution model.," In 3 we describe the evolution of a fiducial, high-resolution model."
+ In 844 we investigate (he dependence of the results on model parameters., In 4 we investigate the dependence of the results on model parameters.
+ And in 855 we describe implications. with an emphasis on Κον open questions: are (he vortices destroved by three-dimensional instabilities?:," And in 5 we describe implications, with an emphasis on key open questions: are the vortices destroyed by three-dimensional instabilities?;"
+ aud cdo mechanisms exist that can inject. vorticity into (he disk?, and do mechanisms exist that can inject vorticity into the disk?
+ The shearing-sheetmodel is obtained via a rigorous expansion of the two-dimensional, The shearing-sheetmodel is obtained via a rigorous expansion of the two-dimensional
+ (PauckCollaboration2010:Màπαetal.2010)...," \citep{PlanckBlueBook,prelauB,prelauM,lamarre2010,Maffei_prelaunch}."
+To estimate the mode parameters. we applied usual spectrum fitting techniques. as in previous similar works (e.g.222)..,"To estimate the mode parameters, we applied usual spectrum fitting techniques, as in previous similar works \citep[e.g.][]{Appourchaux08,Barban09,Deheuvels10}."
+ To summarise the procedure. we assume that the observed PSD is distributed around à mean profile spectrum St) and follows a 2-degree-of-freedom y statistics (e.g.2)..," To summarise the procedure, we assume that the observed PSD is distributed around a mean profile spectrum ${\cal S}(\nu)$ and follows a 2-degree-of-freedom $\chi^2$ statistics \citep[e.g.][]{Duvall86}."
+ (v) is modelled as the sum of the background B(v).considered described in Sect. ??..," ${\cal S}(\nu)$ is modelled as the sum of the background $B(\nu)$, described in Sect. \ref{sec:bg},"
+ and p-mode profiles ην) for each degree / and radial order 51.," and p-mode profiles ${\cal P}_{l,n}(\nu)$ for each considered degree $l$ and radial order $n$ ."
+" μέν) are multiplets of Lorentzian profiles and read as where wi is the azimuthal order. Επι. νε. and Dy, are the height. the frequency. and the width of the mode. v, is the rotational splitting. common to all modes. and «,,(/). the height ratios of multiplet components. are geometrical terms depending only on the inclination angle / (e.g..see?).."," ${\cal P}_{l,n}(\nu)$ are multiplets of Lorentzian profiles and read as where $m$ is the azimuthal order, $H_{l,n}$, $\nu_{l,n}$, and $\Gamma_{l,n}$ are the height, the frequency, and the width of the mode, $\nu_\mathrm{s}$ is the rotational splitting, common to all modes, and $a_{l,m}(i)$, the height ratios of multiplet components, are geometrical terms depending only on the inclination angle $i$ \citep[e.g., see][]{Gizon03}."
+ Owing to cancellation effects of averaging the stellar flux over the whole disc. only modes with /<2 are considered in the present analysis.," Owing to cancellation effects of averaging the stellar flux over the whole disc, only modes with $l \le 2$ are considered in the present analysis."
+ It is worth noticing that we analyse the full light curve disregarding any possible variations in the p-mode parameters due to any magnetic activity effects as the ones recently uncovered in HD 49933 (?).., It is worth noticing that we analyse the full light curve disregarding any possible variations in the p-mode parameters due to any magnetic activity effects as the ones recently uncovered in HD 49933 \citep{Garcia10}.
+ The different parameters have been independently fitted by ten different groups. using either MLE (e.g.?) or Bayesian priors and performing maximum a posterior! estimation (MAP.e.g.?2) or MCMC (e.g. ??)..," The different parameters have been independently fitted by ten different groups, using either MLE \citep[e.g.][]{Appourchaux98} or Bayesian priors and performing maximum a posteriori estimation \citep[MAP, e.g.][]{Gaulme09} or MCMC \citep[e.g.][]{Benomar09Bayes,Handberg11}. ."
+ Generally. global fits of all of the p modes were performed.," Generally, global fits of all of the p modes were performed."
+ Nevertheless. one group has used a local approach by fitting sequences of successive /=2. 0. and | modes. according to the original CoRoT recipes (?)..," Nevertheless, one group has used a local approach by fitting sequences of successive $l=2$, 0, and 1 modes, according to the original CoRoT recipes \citep{Appourchaux06}."
+ Between 8 and 15 orders have been included in the fit. depending on the group.," Between 8 and 15 orders have been included in the fit, depending on the group."
+ To reduce the dimensior of the parameter space. other constraints were used by taking advantage of the smooth variation of heights and widths.," To reduce the dimension of the parameter space, other constraints were used by taking advantage of the smooth variation of heights and widths."
+" As a consequence. only one height and one width. Ho, and Γα,. are fitted for each order."," As a consequence, only one height and one width, $H_{0,n}$ and $\Gamma_{0,n}$, are fitted for each order."
+ We linked the height (resp..," We linked the height (resp.,"
+ width) of/2 mode to the height (width) of the nearby /=0 mode through a relation We denote r» the visibility of the /=2 mode relative to [=0., width) of $l=2$ mode to the height (width) of the nearby $l=0$ mode through a relation We denote $r_2$ the visibility of the $l=2$ mode relative to $l=0$.
+" Concerning /=| modes. we can link them either to the previous or to the following /=0 modes. leading to two fitting situations: where r, the visibility of the /=1 mode relative to /=0."," Concerning $l=1$ modes, we can link them either to the previous or to the following $l=0$ modes, leading to two fitting situations: where $r_1$ the visibility of the $l=1$ mode relative to $l=0$."
+ Even if the variations in heights and widths are limited. they can still be significant over half the large separation.," Even if the variations in heights and widths are limited, they can still be significant over half the large separation."
+ Thus. results can be different for both cases.," Thus, results can be different for both cases."
+ The visibility factors 7; depend mainly on the stellar limb-darkening profile (e.g.?).., The visibility factors $r_l$ depend mainly on the stellar limb-darkening profile \citep[e.g.][]{Gizon03}.
+ It is then possible to fix them to theoretical values. deduced from stellar atmosphere models. or leave them as a free parameters.," It is then possible to fix them to theoretical values, deduced from stellar atmosphere models, or leave them as a free parameters."
+ Six groups have fixed them. whereas four others left them free.," Six groups have fixed them, whereas four others left them free."
+" In the next two sections, we present the results of one of the fits used as reference."," In the next two sections, we present the results of one of the fits used as reference."
+ The reference fitting is based on MLE/MAP. and the fitted frequency range isHz.," The reference fitting is based on MLE/MAP, and the fitted frequency range is."
+" The background component B, is fixed. whereas Bayesian priors for B. and W are derived from the background fit(Sect. ??))."," The background component $B_\mathrm{a}$ is fixed, whereas Bayesian priors for $B_\mathrm{g}$ and $W$ are derived from the background fit(Sect. \ref{sec:bg}) )."
+ There are no Bayesian priors for the mode parameters., There are no Bayesian priors for the mode parameters.
+ The marginal probabilities derived. from MCMC derived by other groups are compatible with the error bars presented hereafter., The marginal probabilities derived from MCMC derived by other groups are compatible with the error bars presented hereafter.
+ For the two fitting configurations (A or B). the fitted frequencies are mm very good agreement.," For the two fitting configurations (A or B), the fitted frequencies are in very good agreement."
+ Table 4 provides modes frequencies obtained with fit A. The fitted spectrum is plotted over the observation in Fig. 8.., Table \ref{tab:freq} provides modes frequencies obtained with fit A. The fitted spectrum is plotted over the observation in Fig. \ref{fig:fit}.
+ The determined frequencies are also plotted over the écchelle diagram (Fig. 9)), The determined frequencies are also plotted over the écchelle diagram (Fig. \ref{fig:echelle}) ).
+ The frequencies listed in the table were found by at least eight groups out of ten within the error bars., The frequencies listed in the table were found by at least eight groups out of ten within the error bars.
+ The three nodes labelled with an exponent in the table correspond to nodes fitted by only four groups., The three modes labelled with an exponent in the table correspond to modes fitted by only four groups.
+ Nevertheless. the four groups have independently obtained consistent frequencies for these modes. so we report them. but they could be less reliable.," Nevertheless, the four groups have independently obtained consistent frequencies for these modes, so we report them, but they could be less reliable."
+ An extra /2 |. aroundj/Hz.. seems to be present in the écchelle diagram and has been fitted by a few groups.," An extra $l=1$ , around, seems to be present in the écchelle diagram and has been fitted by a few groups."
+ Nevertheless. it is very close to oneof the orbital harmonies (see Fig. 9)) ," Nevertheless, it is very close to oneof the orbital harmonics (see Fig. \ref{fig:echelle}) )"
+and has been rejected.,and has been rejected.
+ The radial order 2 provided in Table 4 ts only relative and could be shiftedby +1., The radial order $n$ provided in Table \ref{tab:freq} is only relative and could be shiftedby $\pm 1$ .
+ It is obtained by fitting the relation derived from the asymptotic development for p modes ede, It is obtained by fitting the relation derived from the asymptotic development for p modes \citep[e.g.][]{Tassoul80}. .
+"for both segments. indeed. eives an acceptable combined /dof = 36.9/32. with Nyincreasing from 3+0.1 to 39401<1073 7? and inner disk temperature T;,=0.220.1 keV. The inner disk radius turus out to be 1.50.3«10"" cm corresponding to a ~950AL, black hole.","for both segments, indeed gives an acceptable combined $\chi^2$ /dof = $36.9/32$ , with $N_H$increasing from $3\pm 0.4$ to $3.9\pm 0.4 \times 10^{21}$ $^{-2}$ and inner disk temperature $T_{in} = 0.22 \pm 0.1$ keV. The inner disk radius turns out to be $1.5 \pm 0.3 \times 10^9$ cm corresponding to a $\sim 950 M_\odot$ black hole."
+" If the cohuuu deusitv aud temperature are allowed to vary independently for the two segments. C/dof=29/29. with Ry,=L3+40.1«107 cm corresponding to AL~9003/.."," If the column density and temperature are allowed to vary independently for the two segments, $\chi^2/dof = 29/29$, with $R_{in} = 1.3 \pm 0.4 \times 10^9$ cm corresponding to $M \sim 900 M_\odot$."
+ An additional power-aw component with photon iudex PD—2.0 or D—2.7does nof improve the fit if the column deusitv is fixed at the Galactic value., An additional power-law component with photon index $\Gamma = 2.0$ or $\Gamma = 2.7$does not improve the fit if the column density is fixed at the Galactic value.
+" For Ny~6«1073, a disk black ουν and a power-law component with P=2. provides a good fit to both scements κ/dof=19.90/27."," For $N_H \sim 6 \times 10^{21}$, a disk black body and a power-law component with $\Gamma = 2$, provides a good fit to both segments $\chi^2/dof = 19.9/27$."
+" Tere. he variability is due to a change of normalization of the ΝΟΥανν component. while the temperature Tj,0.13 cV and Ry,~1.2<1019 cni remain uearly constant."," Here, the variability is due to a change of normalization of the power-law component, while the temperature $T_{in} \sim 0.13$ keV and $R_{in} \sim 1.2 \times 10^{10}$ cm remain nearly constant."
+" Again. in these more complicated interpretations Rj, is sienificautly larger plying a larger black hole mass."," Again, in these more complicated interpretations $R_{in}$ is significantly larger implying a larger black hole mass."
+ The N-rav source N-7 in NGC 6916 is located in a region of low diffuse ciuission aud has a count rate which does not cause pile-p inChandra detectors., The X-ray source X-7 in NGC 6946 is located in a region of low diffuse emission and has a count rate which does not cause pile-up in detectors.
+ The light curve of this source reveals a decrease of ~1.5 in the count rate over 5000 sec making it one of the few ULX hat have clearly shown variability ou ksec time-scales (Isvaussctal.2005:Mizunoet2001).," The light curve of this source reveals a decrease of $\sim 1.5$ in the count rate over $5000$ sec making it one of the few ULX that have clearly shown variability on ksec time-scales \citep{Kra05,Miz01}."
+. Spectral fitting of the high aud low count rate parts of he light curve reveals that a simple modelof a disk black )ody. ehussion absorbed by the Calactic column density of Ny=2.0«107! 7 can adequately represent voth scemeuts., Spectral fitting of the high and low count rate parts of the light curve reveals that a simple modelof a disk black body emission absorbed by the Galactic column density of $N_H = 2.1 \times 10^{21}$ $^{-2}$ can adequately represent both segments.
+ The best fit temperature varies frou 26 to 0.29 keV. while the quer disk radius remains constant at ~6<107 cin., The best fit temperature varies from $0.26$ to $0.29$ keV while the inner disk radius remains constant at $ \sim 6 \times 10^8$ cm.
+ This would imply that the nass of the black hole is ~LOOAL.. aud at a luuinosity of ~3«1079 eres/see the system has an Eddinetou ratio of ~0.06., This would imply that the mass of the black hole is $\sim 400 M_\odot$ and at a luminosity of $\sim 3 \times 10^{39}$ ergs/sec the system has an Eddington ratio of $\sim 0.06$.
+ Other more complicated spectral fits ike assunuues intrinsic column density for the source and/or addition of a power-law component results in a arecr estimation of iuner disky. radius aud hence a larecr lack hole mass., Other more complicated spectral fits like assuming intrinsic column density for the source and/or addition of a power-law component results in a larger estimation of inner disk radius and hence a larger black hole mass.
+ Although the huniuositv of the source estimated here is only nuüildlv super-Eddiustou for teu solar mass black hole. the low imner disk teniperature ~0.3 keV implies a larger iuner disk radius aud hence a laree black hole mass ~40037...," Although the luminosity of the source estimated here is only mildly super-Eddington for ten solar mass black hole, the low inner disk temperature $\sim 0.3$ keV implies a larger inner disk radius and hence a large black hole mass $\sim 400 M_\odot$."
+ A low sub-Eddingtou accretion rate on a ten solar imass black bole could produce the observed low inucr disk temperature. but then the predicted huninosity would be siguificautlv less than what is observed.," A low sub-Eddington accretion rate on a ten solar mass black hole could produce the observed low inner disk temperature, but then the predicted luminosity would be significantly less than what is observed."
+ As cautioned by Miller.Fabiau&Miller (2001).. the source could be a background ACN and the soft component observed is actually the soft excess which is detected in mauy AGN.," As cautioned by \cite{Mil04}, the source could be a background AGN and the soft component observed is actually the soft excess which is detected in many AGN."
+ These soft excesses ean be modeled as black body emission at &T.—0.1 keV simular to the softcomponent in ULX., These soft excesses can be modeled as black body emission at $kT \sim 0.1$ keV similar to the softcomponent in ULX.
+ While this may be true for some ULX. this is uulikclv iuthis case because here (unlike soft excess of ACIN) the soft colponcut totally dominates the Iuninosity.," While this may be true for some ULX, this is unlikely inthis case because here (unlike soft excess of AGN) the soft component totally dominates the luminosity."
+ If the emission is beamed sinee£ef2001) by a factor jg. the inner disk radius would be overestimated by a7.," If the emission is beamed \citep{Kin01} by a factor$\eta$ , the inner disk radius would be overestimated by $\eta^{1/2}$ ."
+ However. even for extreme ecometrics Misa& have computed y« 5.," However, even for extreme geometries \cite{Mis03} have computed $\eta < 5$ ."
+ Hence an extreme beaming of — 5. would imply that Πο—2.7«10? cn corresponding to a black hole mass of 18047...," Hence an extreme beaming of $\eta \sim 5$ , would imply that $R_{in} \sim 2.7 \times 10^8$ cm corresponding to a black hole mass of $180 M_\odot$ ."
+ The, The
+one epoch to the next.,one epoch to the next.
+ Thus. it is a reasonable scenario that the variability may be caused by the change of accretion rates.," Thus, it is a reasonable scenario that the variability may be caused by the change of accretion rates."
+ In this work. we investigate the variability amplitude of quasars as a function of black hole mass based on accretion dise models assuming the variability to be caused by the change of accretion rates.," In this work, we investigate the variability amplitude of quasars as a function of black hole mass based on accretion disc models assuming the variability to be caused by the change of accretion rates."
+ We also compare our theoretical caleulations with the observed correlations of variability amplitude with other quantities (i.e.. luminosity or observed wavelength).," We also compare our theoretical calculations with the observed correlations of variability amplitude with other quantities (i.e., luminosity or observed wavelength)."
+ The big blue bump observed in many AGN is believed to be related to some Kind of thermal emission (Malkan&Sargent1982:Shalyapinetal. 2002).," The big blue bump observed in many AGN is believed to be related to some kind of thermal emission \citep{m1982,s2002}."
+. The thermal emission of the standard thin accretion dise can reproduce the general features of AGN continua. provided some other components. such as an extrapolation of infrared continuum or a power-law X-ray continuum. are included (Sun&Malkan1989:Laor1990) (butalsoseeGaskell2007).," The thermal emission of the standard thin accretion disc can reproduce the general features of AGN continua, provided some other components, such as an extrapolation of infrared continuum or a power-law X-ray continuum, are included \citep{s1989,l1990} \citep*[but also
+see][]{g2007}."
+. The standard thin dise is geometrically thin. optically thick. and its emission is close to blackbody (Shakura&Sunyaev1973).," The standard thin disc is geometrically thin, optically thick, and its emission is close to blackbody \citep{s1973}."
+ The gases in the dise rotate at Keplerian velocities. and they are in hydrodynamical equilibrium in the vertical direction.," The gases in the disc rotate at Keplerian velocities, and they are in hydrodynamical equilibrium in the vertical direction."
+ The structure of a standard thin dise can be derived analytically (see.e.g..Katoal. 1998).," The structure of a standard thin disc can be derived analytically \citep*[see,
+e.g.,][]{kato1998}."
+. In this work. we adopt the temperature distribution of an accretion dise as a function of radius /? given by Katoetal. 1905). Here Al is the black hole mass. 2.=2ALGe. Apad=10mesm=AL/AL. anda is the viscosity parameter.," In this work, we adopt the temperature distribution of an accretion disc as a function of radius $R$ given by \citet{kato1998}, , where, and Here $M$ is the black hole mass, $R_{\rm g}=2MG/c^2$, $\dot{M}_{\rm
+Edd}=1.5\times10^{18}m{\rm gs^{-1}}, m=M/M_{\odot}$ , and $\alpha$ is the viscosity parameter."
+ The spectrum of the accretion dise can be calculated with where {10} isgiven by Eq. CD. fin=340.," The spectrum of the accretion disc can be calculated with where $T(R)$ isgiven by Eq. \ref{tempdisc}) ), $R_{\rm in}=3R_{\rm
+g}$,"
+ Hou is the outer radius of disc. / is the Plank's constant. / is the inclination of axis of the dise with respect to the line of sight. and £ is the distance from the observer to the black hole (Franketal.2002).," $R_{\rm out}$ is the outer radius of disc, $h$ is the Plank's constant, $i$ is the inclination of axis of the disc with respect to the line of sight, and $D$ is the distance from the observer to the black hole \citep{f2002}."
+. The optical/UV continuum of a standard thin dise is fix which is inconsistent with fi«xf»+ observed in many quasars.," The optical/UV continuum of a standard thin disc is $f_\nu\propto\nu^{1/3}$ , which is inconsistent with $f_\nu\propto\nu^{-1}$ observed in many quasars."
+ Gaskell(2007) suggested that the observed quasar spectra can be reproduced by accretion dises with a temperature gradient of F(R)xROO instead of (I)xI7 as predicted by standard thin dise model.," \citet{g2007}
+ suggested that the observed quasar spectra can be reproduced by accretion discs with a temperature gradient of $T(R)\propto{R^{-0.57}}$ instead of $T(R)\propto{R^{-0.75}}$ as predicted by standard thin disc model."
+ Thus. we also adopt this modified disc model in our spectral calculations (this model is referred to as moditied accretion dise model hereafter).," Thus, we also adopt this modified disc model in our spectral calculations (this model is referred to as modified accretion disc model hereafter)."
+ The temperature of the dise is then given by where Cu... is a correction factor used to let the resulted disc luminosity be the same as that for a standard thin dise with the same dise parameters., The temperature of the disc is then given by where ${\cal C}_{\rm cor}$ is a correction factor used to let the resulted disc luminosity be the same as that for a standard thin disc with the same disc parameters.
+ The spectra of the modified accretion discs can be calculated with Eqs. (3) , The spectra of the modified accretion discs can be calculated with Eqs. \ref{spectdisc}) )
+and (3))., and \ref{tempdisc2}) ).
+ The spectra of standard thin dises ean be calculated with Eqs. C» , The spectra of standard thin discs can be calculated with Eqs. \ref{tempdisc}) )
+and (O3. provided the values of dise parameters (a7. mi and à) are supplied.," and \ref{spectdisc}) ), provided the values of disc parameters $m$, $\dot{m}$ and $\alpha$ ) are supplied."
+ In Fig. |.," In Fig. \ref{freq},"
+ we plot the spectra of the dises with different black hole masses and accretion rates., we plot the spectra of the discs with different black hole masses and accretion rates.
+ For comparison. we also plot the spectra of the dises ealeulated with Eqs. (3))," For comparison, we also plot the spectra of the discs calculated with Eqs. \ref{tempdisc2}) )"
+ and (2) based on the moditied dise model described in last section., and \ref{spectdisc}) ) based on the modified disc model described in last section.
+ In the figure. we find that both standard thin dise model and modified disc model naturally show the optical/UV variability becomes larger with the increase of black hole mass. if the accretion rate changes at the same percentage.," In the figure, we find that both standard thin disc model and modified disc model naturally show the optical/UV variability becomes larger with the increase of black hole mass, if the accretion rate changes at the same percentage."
+ Assuming the optical variability to be caused by the change of accretion rates in AGN. we calculate the optical variability amplitude as a function of black hole mass with the dise models deseribed in last section.," Assuming the optical variability to be caused by the change of accretion rates in AGN, we calculate the optical variability amplitude as a function of black hole mass with the disc models described in last section."
+" In our calculations. the accretion rate varies in the range of my—rmi,dOris where mi,=0.1 and om=Odmy for standard disc model (970.=0.515 for the modified dise model). are adopted."," In our calculations, the accretion rate varies in the range of $\dot{m}=\dot{m}_0\pm\delta \dot{m}$, where $\dot{m}_0=0.1$ and $\delta\dot{m}=0.4\dot{m}_0$ for standard disc model $\delta\dot{m}=0.5\dot{m}_0$ for the modified disc model), are adopted."
+ In Fig. 2..," In Fig. \ref{varia},"
+ we compare our calculations with the correlation between variability amplitude at R-band and black hole mass discovered by Woldetal.(2007)., we compare our calculations with the correlation between variability amplitude at R-band and black hole mass discovered by \citet{wold2007}.
+. We find that our calculations can reproduce their statistic result quite welltsee Fig. 25., We find that our calculations can reproduce their statistic result quite well(see Fig. \ref{varia}) ).
+ We also calculate the optical variability amplitude as functions of luminosity by changing accretion rates (7 for given black hole mass m (see Fig. 91)., We also calculate the optical variability amplitude as functions of luminosity by changing accretion rates $\dot{m}$ for given black hole mass $m$ (see Fig. \ref{lumin}) ).
+ We find that the optical variability amplitude decreases with increasing luminosity for a given black hole mass. which is qualitatively consistent with the observed anti-correlation betweenthese two quantities (e.g.Cristianietal.1996:Vanden 2007). ," We find that the optical variability amplitude decreases with increasing luminosity for a given black hole mass, which is qualitatively consistent with the observed anti-correlation betweenthese two quantities \citep*[e.g.,][]{c1996,vb2004,w2007}. ."
+The dependence of the optical variability amplitude on the observed wavelength measured in the rest frame can be calculated based on either the standard thin disc model or the modified dise model for given dise parameters (see Fig. 43., The dependence of the optical variability amplitude on the observed wavelength measured in the rest frame can be calculated based on either the standard thin disc model or the modified disc model for given disc parameters (see Fig. \ref{wavelength}) ).
+ It is found, It is found
+deflection law given by general relativity.,deflection law given by general relativity.
+ ILowever. lensing measurements have limitations ol their own.," However, lensing measurements have limitations of their own."
+ Weak lensing mass reconstruction of individual clusters can only obtain a two-climensional surface mass densitv: information about the structure along the line of sight is irretrievably lost., Weak lensing mass reconstruction of individual clusters can only obtain a two-dimensional surface mass density; information about the structure along the line of sight is irretrievably lost.
+ Typically these measurements come [rom pointed observations of known massive clusters. selected by X-ray luminosity or other property.," Typically these measurements come from pointed observations of known massive clusters, selected by X-ray luminosity or other property."
+ Because the observations need {ο go deep enough to obtain the necessary density of background source ealaxies to make spatially resolved mass maps. they are time-consuming and therefore tend to be limited in area. usually not extending much bevond the virial radius.," Because the observations need to go deep enough to obtain the necessary density of background source galaxies to make spatially resolved mass maps, they are time-consuming and therefore tend to be limited in area, usually not extending much beyond the virial radius."
+ Because of this limited range of physical scales covered. there are uncertainties arising [rom (he mass sheet degeneracy (Dradaéetal.2004)...," Because of this limited range of physical scales covered, there are uncertainties arising from the mass sheet degeneracy \citep{bradac:mass-sheet}."
+ In addition. large-scale structure along the line of sight and chance projections with clusters at different. redshifts can influence individual cluster mass determinations. causing ~30% scatter around the (rue mass and possibly a bias in (he mass determination (Cen1997:Metzleretal.1999:Hoekstra2001:Whiteetal.2002:Hoekstra2003:Clowe2004:Doclelson 2004)..," In addition, large-scale structure along the line of sight and chance projections with clusters at different redshifts can influence individual cluster mass determinations, causing $\sim 30$ scatter around the true mass and possibly a bias in the mass determination \citep{cen:projection-effects,metzler:cluster-lss-contam,hoekstra:lss-contam1,metzler:lss-contam,
+white:clusters-completeness,hoekstra:lss-contam,clowe:lensing-proj,dodelson:lss-contam}."
+ In this paper. we demonstrate a new wav to measure (he average density and mass profile of clusters using stacked samples that. in principle. can be selected by any observable proxy.," In this paper, we demonstrate a new way to measure the average density and mass profile of clusters using stacked samples that, in principle, can be selected by any observable proxy."
+ We introduce a completely non-parametric method to deproject the average shear profile and obtain the mean 3D density ancl mass profiles lor a sample., We introduce a completely non-parametric method to deproject the average shear profile and obtain the mean 3D density and mass profiles for a sample.
+ This method only assumes (he statistical isotropy of the cluster-mass correlation Dhunction., This method only assumes the statistical isotropy of the cluster-mass correlation function.
+ We perform tests on an N-hocly simulation to show that the method correctly recovers the average 3D density and mass profiles., We perform tests on an N-body simulation to show that the method correctly recovers the average 3D density and mass profiles.
+ From the derived mass profile of the clusters. one can measure a virial radius and virial mass without assuming a model lor the profile.," From the derived mass profile of the clusters, one can measure a virial radius and virial mass without assuming a model for the profile."
+ Unlike individual cluster lensing reconstructions. these mass estimates are not affected. by large-scale structure or chance projections along the line of sight.," Unlike individual cluster lensing reconstructions, these mass estimates are not affected by large-scale structure or chance projections along the line of sight."
+ We show that these methods will allow one to constrain (the densitwv profiles of galaxy clusters. calibrate mass-observable relations. ancl thereby measure the cluster mass function and bias.," We show that these methods will allow one to constrain the density profiles of galaxy clusters, calibrate mass-observable relations, and thereby measure the cluster mass function and bias."
+ In 82. we derive (he relations between tangential shear measurements and the 2D surface mass densitv.," In 2, we derive the relations between tangential shear measurements and the 2D surface mass density."
+ In 83. we invert the relation to obtain the 3D density. profile.," In 3, we invert the relation to obtain the 3D density profile."
+ In 84. we present methods of deriving 3D mass profiles. and in 85 we describe some of the practical steps Gnterpolation. extrapolation. ancl error propagation) needed to implement the method.," In 4, we present methods of deriving 3D mass profiles, and in 5 we describe some of the practical steps (interpolation, extrapolation, and error propagation) needed to implement the method."
+" In 86. we apply the method to a large N-body simulation. showing that it accurately recovers the correct 3D densitv and mass profiles statistically, and in 87 we discuss why the method is insensitive {ο asphliericity of halos and projected laree-scale structure alone the line of sight."," In 6, we apply the method to a large N-body simulation, showing that it accurately recovers the correct 3D density and mass profiles statistically, and in 7 we discuss why the method is insensitive to asphericity of halos and projected large-scale structure along the line of sight."
+ $5 presents our conclusions and outlook for the future., 8 presents our conclusions and outlook for the future.
+ In future papers. we will apply. this method to measurement of the cluster-mass correlation function for optically selected clusters in the SDSS and in N-body simulations. derive scaling relations between optical richness ancl virial mass. and constrain (he bias aud the mass power spectrum by comparison with the," In future papers, we will apply this method to measurement of the cluster-mass correlation function for optically selected clusters in the SDSS and in N-body simulations, derive scaling relations between optical richness and virial mass, and constrain the bias and the mass power spectrum by comparison with the"
+1990).,.
+". In a previous study, we derived the photospheric horizontal velocity in 14 different quiet regions to justify the Alfvénn wave model for coronal heating 2010)."," In a previous study, we derived the photospheric horizontal velocity in 14 different quiet regions to justify the Alfvénn wave model for coronal heating \citep{mats10}."
+". In this study, we will show the temporal evolution of the horizontal motion of the photosphere in more detail."," In this study, we will show the temporal evolution of the horizontal motion of the photosphere in more detail."
+" When discerning between various coronal heating models (see reviews by, e.g., Mandrini et al."," When discerning between various coronal heating models (see reviews by, e.g., Mandrini et al."
+" 2000, Aschwanden et al."," 2000, Aschwanden et al."
+" 2001, Klimchuk 2006, and references therein), the power spectra derived in the present study will be of high value."," 2001, Klimchuk 2006, and references therein), the power spectra derived in the present study will be of high value."
+" was launched on September 22, 2006 in order to investigate the unsolved problems of solar physics, such as coronal heating."," was launched on September 22, 2006 in order to investigate the unsolved problems of solar physics, such as coronal heating."
+" Solar Optical Telescope (SOT) (TsunetaSuematsuetal.2008;IchimotoShimizu 2008),, one of the three telescopes mounted onHinode,, is an optical telescope whose spatial resolution is unprecedentedly high (0722-0733, or 150-200 km) for a solar telescope in space."," Solar Optical Telescope (SOT) \citep{tsun08,suem08,ichi08,shim08}, one of the three telescopes mounted on, is an optical telescope whose spatial resolution is unprecedentedly high 3, or 150-200 km) for a solar telescope in space."
+" As opposed to ground based observations that are affected by atmospheric seeing, seeing-free data sets over a long time span can be continuously obtained by SOT."," As opposed to ground based observations that are affected by atmospheric seeing, seeing-free data sets over a long time span can be continuously obtained by SOT."
+" We selected 14 data sets of continuously observed quiet regions with the G--band filter between October 31, 2006 and December 29, 2007."," We selected 14 data sets of continuously observed quiet regions with the -band filter between October 31, 2006 and December 29, 2007."
+" Our basic selection criteria was for the data set to have a duration longer than 70 minutes, with a mean time cadence less than 32 s. Duration of the selected data sets ranges from 75 min to 345 min and the cadence is almost constant with the accuracy of 1 second."," Our basic selection criteria was for the data set to have a duration longer than 70 minutes, with a mean time cadence less than 32 s. Duration of the selected data sets ranges from 75 min to 345 min and the cadence is almost constant with the accuracy of 1 second."
+ Our data sets are selected to be located within + 100 arcsec from the disk center so that the line of sight is within 6 degrees from the local normal., Our data sets are selected to be located within $\pm$ 100 arcsec from the disk center so that the line of sight is within 6 degrees from the local normal.
+" Pixel resolution of the CCD is 070054 (40 km) and FOV is larger than 27""x 27""(20,000 km x 20,000 km)."," Pixel resolution of the CCD is 054 (40 km) and FOV is larger than $\times$ (20,000 km $\times$ 20,000 km)."
+ We applied dark current subtraction and flat-fielding in the standard manner for all of the G--band images., We applied dark current subtraction and flat-fielding in the standard manner for all of the -band images.
+" The solar rotation, tracking error, and satellite jitter causes overall"," The solar rotation, tracking error, and satellite jitter causes overall"
+20063.,.
+. They may be responsible for the apparent peak ποσα b—Οἱ in the Bouchetetal.(2008). latitude profiles., They may be responsible for the apparent peak near $b = 0^\circ$ in the \citet{bouchet07} latitude profiles.
+ The hard X-ray coutimmiun is consistent with the predictions iu. beth intensity. aud spectral index., The hard X-ray continuum is consistent with the predictions in both intensity and spectral index.
+" However, the uncertainties in the model are still considerable: the distribution of CR sources and gas in the inner Calaxy which affect both the primary and secondary electrons/positrous. aud the optical aud infrared part of the ΤΠΕ (the CAIB is of course known exactly)."," However, the uncertainties in the model are still considerable: the distribution of CR sources and gas in the inner Galaxy which affect both the primary and secondary electrons/positrons, and the optical and infrared part of the ISRF (the CMB is of course known exactly)."
+ Iu fact. our optimised model overpredicts the SPI data. which could simply reflect these uncertainties. but the agreement in the spectral shape gives coufideuce that the mechanisiua is correctly identified.," In fact, our optimised model overpredicts the SPI data, which could simply reflect these uncertainties, but the agreement in the spectral shape gives confidence that the mechanism is correctly identified."
+ There is still an excess iu the COMPTEL euerev range between d30 MeV which is to be explained., There is still an excess in the COMPTEL energy range between 1–30 MeV which is to be explained.
+ The contribution of the positron aunihilation iu flight may contribute in this energy range (Beacom&Yüksel2006).. but it has to be tested against the iuteusitv of the 511 keV line and positronin continu.," The contribution of the positron annihilation in flight may contribute in this energy range \citep{Beacom06}, but it has to be tested against the intensity of the 511 keV line and positronium continuum."
+ From our modelling. we fud that the total rate of secondary positron production by CRs iu the whole Galaxy is ~2&LO 3! in the optimised model.," From our modelling, we find that the total rate of secondary positron production by CRs in the whole Galaxy is $\sim2\times10^{42}$ $^{-1}$ in the optimised model."
+ The conventional niodel gives a factor of ~2 less positrons., The conventional model gives a factor of $\sim$ 2 less positrons.
+ These values are ~107€ of the positron aunililation rate ~LS&100 3 as denied from INTEGRAL observations of the 511 keV line enssiou (Ixuódlsederetal. 2005)., These values are $\sim 10$ of the positron annihilation rate $\sim 1.8\times10^{43}$ $^{-1}$ as derived from INTEGRAL observations of the 511 keV line emission \citep{Knodlseder2005}.
+. The current CR fux of positrons is nof sufficient to account for the observed aunilülatiou rate., The current CR flux of positrons is not sufficient to account for the observed annihilation rate.
+ ACR origin for the 511 keV annihilation liue can still be reconciled with the production rate if CR intensities in the past were higher., A CR origin for the 511 keV annihilation line can still be reconciled with the production rate if CR intensities in the past were higher.
+ Ow work illustrates the intrinsic connection between the diffuse Calactic ciission in differeut energv ranges., Our work illustrates the intrinsic connection between the diffuse Galactic emission in different energy ranges.
+ Inverse Compton cunission bv CR olectrous aud positrons on starlieht aud infrared radiatiou are the most inuportaut compoucuts of the hard N-rav aud clnission in the 100 keV to few MeV range., Inverse Compton emission by CR electrons and positrons on starlight and infrared radiation are the most important components of the hard X-ray and emission in the 100 keV to few MeV range.
+ A considerable proportion of this enüssion is produced by secondary electrous audpositrons. the spectrum of which depends ou the CR nuclei spectrum at energies ~ few GeV and higher.," A considerable proportion of this emission is produced by secondary electrons and positrons, the spectrum of which depends on the CR nuclei spectrum at energies $\sim$ few GeV and higher."
+" These CRs also produceGLAST z""-decav iat dominate the emission in the range fron LOO MeV to ~10 GeV. Hence. (LAST observations of the z)-decavy diffuse cussion will also constrain in the future 16 contribution by secondary electrous/positrous to the diffuse cussion in the SPI euerev range."," These CRs also produce $\pi^0$ -decay that dominate the emission in the GLAST range from 100 MeV to $\sim$ 10 GeV. Hence, GLAST observations of the $\pi^0$ -decay diffuse emission will also constrain in the future the contribution by secondary electrons/positrons to the diffuse emission in the SPI energy range."
+ With 16 secondary electrous/positrous fixed. SPI observatious xobe the IC cussion of primary CR electrons with energies =LO GeV scattering the infrared compoucut of 1ο ISRF and the CMD.," With the secondary electrons/positrons fixed, SPI observations probe the IC emission of primary CR electrons with energies $\lesssim 10$ GeV scattering the infrared component of the ISRF and the CMB."
+ This will provide information ou the low energv spectzni of primary CR clectrous and the infrared component of the ISRE., This will provide information on the low energy spectrum of primary CR electrons and the infrared component of the ISRF.
+ Iu tun. since uost of the diffuse cussion between -—10 GeV 10 TeV is produced via IC scattering of primary clectrous on the same starleht aud iufraved photons. this xovides a connection to observations of diffuse enission at TeV energies by TESS (Aharonianetal.2006). aud AIILACRO (Abdoetal.2007).," In turn, since most of the diffuse emission between $\sim$ 10 GeV – 10 TeV is produced via IC scattering of primary electrons on the same starlight and infrared photons, this provides a connection to observations of diffuse emission at TeV energies by HESS \citep{aharonian2006} and MILAGRO \citep{abdo2007}."
+. aacknowledges partial support from the US Departineut of Encre., acknowledges partial support from the US Department of Energy.
+ aacknowledges partial support from NASA Astronomy and Plivsics Research aud Analysis Program (APRA) evant., acknowledges partial support from NASA Astronomy and Physics Research and Analysis Program (APRA) grant.
+if the total energy is estimated from the kinetic energv ancl virial theorem. both the NFW and Moore. profiles are in the collapsed phase.,"if the total energy is estimated from the kinetic energy and virial theorem, both the NFW and Moore profiles are in the collapsed phase."
+ Otherwise. small concentration parameters allow either phase (in the bistable region). depending on the detailed route to equilibrium.," Otherwise, small concentration parameters allow either phase (in the bistable region), depending on the detailed route to equilibrium."
+ Without the virial assumption. haloes with energy £ lie in the collapsed or bistable regions as well. mostlv in the former.," Without the virial assumption, haloes with energy $\xi^{(c)}$ lie in the collapsed or bistable regions as well, mostly in the former."
+ We have investigated (he equilibrium configurations of NN’ sell-gravitating collisionless particles. interacting via a softened gravitational potential. in the AICE and mean-field Iimit.," We have investigated the equilibrium configurations of $N$ self-gravitating collisionless particles, interacting via a softened gravitational potential, in the MCE and mean-field limit."
+ Below a critical energy.Ort So£.2—0.335. a system with ez0 exists in a stable. collapsed phase.," Below a critical energy, $\xi_c \simeq -0.335$, a system with $\epsilon\neq 0$ exists in a stable, collapsed phase."
+ This phase is unstable for pure eeravitv (€= 0)., This phase is unstable for pure gravity $\epsilon=0$ ).
+ Above another critical enerevMS €eQ. both softened. ancl unsoftened systems exist in an stable. extended. phase.," Above another critical energy $\xi_c^{(+)} \simeq 0$, both softened and unsoftened systems exist in an stable, extended phase."
+ In (the intermediate region. £.2005«€x2x)&©. both the collapsed and extended phases are accessible:. the detailed. route to equilibrium determines which one is picked out.," In the intermediate region $\xi_c < \xi < \xi_c^{(+)}$, both the collapsed and extended phases are accessible; the detailed route to equilibrium determines which one is picked out."
+ The density profiles for the extended and collapsed phases are «qualitatively different., The density profiles for the extended and collapsed phases are qualitatively different.
+ The extended profile has a flat core and outer envelope., The extended profile has a flat core and near-isothermal outer envelope.
+ The collapsed profile is a centrally condensed Dirac peak. whose logarithmic slope depends on e.," The collapsed profile is a centrally condensed Dirac peak, whose logarithmic slope depends on $\epsilon$."
+ We compare our results with published N-body simulations by using the soltening parameter. e. size. rogg. and concentration parameter. e. (o place simulated haloes on the £-e phase diagram.," We compare our results with published $N$ -body simulations by using the softening parameter, $\epsilon$, size, $r_{200}$, and concentration parameter, $c$, to place simulated haloes on the $\xi$ $\epsilon$ phase diagram."
+ We find that many published simulations inadvertently sample (the collapsed phase only. even (hough this phase is unstable for pure gravity. ancl arguably irrelevant astrophvsically.," We find that many published simulations inadvertently sample the collapsed phase only, even though this phase is unstable for pure gravity and arguably irrelevant astrophysically."
+ We remind the reader (hat we neglect several effects (hat are important in real CDM haloes. such as hierarchical clustering. nonzero angular momentum. and Cosmological expansion.," We remind the reader that we neglect several effects that are important in real CDM haloes, such as hierarchical clustering, nonzero angular momentum, and cosmological expansion."
+ Our results elucidate some of the artificial behaviour (hat a softened potential can introduce: ον are not a substitute for a full N-body calculation., Our results elucidate some of the artificial behaviour that a softened potential can introduce; they are not a substitute for a full $N$ -body calculation.
+ We are grateful to Bruce MelIxellar for extensive discussions on the theoretical basis of (he mean-field equations. and the anouvimous referee Lor useful comments (hat improved the manuscript.," We are grateful to Bruce McKellar for extensive discussions on the theoretical basis of the mean-field equations, and the anonymous referee for useful comments that improved the manuscript."
+ This work was supported by the Australian Research Council Discovery. Project erant 0208613., This work was supported by the Australian Research Council Discovery Project grant 0208618.
+ CAIT acknowledges the funding provided by an Australian. Postgraduate Award., CMT acknowledges the funding provided by an Australian Postgraduate Award.
+svnthesis study by Duliketal.(2011).. however. has demonstrated the formation of BBIIs with high chirp masses (~15.U. ) from the two ΟΙΑΗ svstems. ancl has vielded a similarly high rate of 0.36Mpc.Myr.| corresponding to 3.6 detections a vear.,"synthesis study by \citet{Bulik08}, however, has demonstrated the formation of BBHs with high chirp masses $\sim 15 M_{\odot}$ ) from the two BH-WR systems, and has yielded a similarly high rate of $0.36 \hspace{0.5mm}\rm{Mpc}^{-3}
+\hspace{0.5mm} \rm{Myr}^{-1}$ corresponding to 3.6 detections a year."
+ Thev suggest that either currently emploved searches are insensitive to higher mass DDII inspirals or that there is an additional aspect to the evolution of such svstems that has not so far been considered., They suggest that either currently employed searches are insensitive to higher mass BBH inspirals or that there is an additional aspect to the evolution of such systems that has not so far been considered.
+ To take account of this uncertainty. we take ro as a higher rate.," To take account of this uncertainty, we take $r_{2}$ as a higher rate."
+ The cosmic coalescence rate can be extrapolated [rom the local rate density rj. bv assuming (he rate (racks the star formation rate (SER)., The cosmic coalescence rate can be extrapolated from the local rate density $r_0$ by assuming the rate tracks the star formation rate (SFR).
+" Explicitly. the differential GW event rate in the redshift shell 2 to z+dz can be written as with dV the cosmology dependent co-moving volume element eiven bv where the Hubble parameter //(2)=ον+QO,,(1z)*]U? and r(z) is the comoving distance related to the luminosity distance by dj,=r(1+2)."," Explicitly, the differential GW event rate in the redshift shell $z$ to $z + dz$ can be written as with $dV$ the cosmology dependent co-moving volume element given by where the Hubble parameter $H(z)=
+H_{0}[\Omega_{\Lambda}+\Omega_{m}(1+z)^{3}]^{1/2}$ and $r(z)$ is the comoving distance related to the luminosity distance by $d_{L}=r(1+z)$."
+" We use the parameters |! with h—0.7. 0,,20.3 and Q4—0.7 (Komatsuetal.2009).."," We use the parameters $H_{0}=100 h\cdot
+\rm{km}\hspace{0.5mm} \rm{s}^{-1}\hspace{0.5mm} \rm{Mpc}^{-1}$ with $h=0.7$, $\Omega_{m}=0.3$ and $\Omega_{\Lambda}=0.7$ \citep{cosmology}."
+ source rate density evolution is accounted lor by the dimensionless evolution [actor e(z). normalized to unity in our local intergalactic neighbourhood.," Source rate density evolution is accounted for by the dimensionless evolution factor $e(z)$, normalized to unity in our local intergalactic neighbourhood."
+" Following Regimban (2009).. we define e(z)=f,(z)/p.(0) normalized to unity at z=0 where relates the SFR to the BBIL coalescence rate."," Following \citet{Regimbau09}, we define $e(z)=
+\dot{\rho}_{\ast,c}(z) / \dot{\rho}_{\ast,c}(0)$ normalized to unity at $z=0$ where relates the SFR to the BBH coalescence rate."
+" Here. p, is the SER density in AZ.vr.1Mpe.7. based on the parametric form of Hopkins&Beacom (2006).. derived from recent measurements ol the galaxy luminosity function."," Here, $\dot{\rho}_{\ast}$ is the SFR density in $M_{\odot}
+\hspace{0.5mm}\rm{yr}^{-1} \hspace{0.5mm}\rm{Mpc}^{-3}$, based on the parametric form of \citet{SFR}, , derived from recent measurements of the galaxy luminosity function."
+ To allow for uncertainties in the SFR. we also consider arecent model described in Wilkinsetal.(2003).. obtained through measurements of the stellar mass density this model gives a much lower rate for z>1.," To allow for uncertainties in the SFR, we also consider arecent model described in \citet{Wilkins}, obtained through measurements of the stellar mass density – this model gives a much lower rate for $z
+> 1$."
+ The [actors z and z; represent (he redshift values of BBIIs merger aud DDII svstem formation respectively., The factors $z$ and $z_f$ represent the redshift values of BBHs merger and BBH system formation respectively.
+" The delay tme lor BBIIs. /;. is given by the difference in lookback timesbetween 2, and 2 "," The delay time for BBHs, $t_d$ , is given by the difference in lookback timesbetween $z_f$ and $z$ "
+Galaxy clusters are large gravitationally bound. structures of a size of -«1—3 Mpe. which have arisen in the hierarchical structure formation of the Universe.,"Galaxy clusters are large gravitationally bound structures of a size of $\sim$ 1–3 Mpc, which have arisen in the hierarchical structure formation of the Universe."
+ Intergalactic gas fills the space between galaxies in the clusters., Intergalactic gas fills the space between galaxies in the clusters.
+ The typical number density and temperature values of this intergalactic gas are 10 '-10? * and 2-10 keV. respectively(for a review. see e.g. Sarazin 1986).," The typical number density and temperature values of this intergalactic gas are $10^{-1}$ $10^{-3}$ $^{-3}$ and 2-10 keV, respectively(for a review, see e.g., Sarazin 1986)."
+ Inverse Compton scattering between cosmic microwave background (CMB) photons and hot free electrons in clusters of galaxies causes a change in the intensity of the CMB radiation towards clusters of galaxies (the Sunyaev-Zeldovich effect. hereinafter the SZ etfect: for a review. see Sunyaev Zeldovich 1980).," Inverse Compton scattering between cosmic microwave background (CMB) photons and hot free electrons in clusters of galaxies causes a change in the intensity of the CMB radiation towards clusters of galaxies (the Sunyaev-Zel'dovich effect, hereinafter the SZ effect; for a review, see Sunyaev Zel'dovich 1980)."
+ The SZ ettect is an important tool for the study of clusters of galaxies (for a review. see Birkinshaw 1999). since both the amplitude and spectral form of the CMB distortion caused by the SZ effect depend on the intergalactic gas parameters (such as the gas number density and temperature).," The SZ effect is an important tool for the study of clusters of galaxies (for a review, see Birkinshaw 1999), since both the amplitude and spectral form of the CMB distortion caused by the SZ effect depend on the intergalactic gas parameters (such as the gas number density and temperature)."
+ A relativistically correct formalism for the SZ ettect based on the probability distribution of the photon frequency shift after scattering was given by Wright (1979) to describe the Componization process of soft photons by high temperature plasma., A relativistically correct formalism for the SZ effect based on the probability distribution of the photon frequency shift after scattering was given by Wright (1979) to describe the Comptonization process of soft photons by high temperature plasma.
+ SZ relativistic corrections are significant even for galaxy clusters with temperatures of 3-5 keV and allow us to derive the gas temperaure from SZ observations (see. e.g. Colafrancesco Marehegiani 2010).," SZ relativistic corrections are significant even for galaxy clusters with temperatures of 3-5 keV and allow us to derive the gas temperature from SZ observations (see, e.g., Colafrancesco Marchegiani 2010)."
+ The constraints on the gas temperature of high temperature clusters from measurements of SZ relativistic corrections have been firstly obtained by Pointecouteau et al. (, The constraints on the gas temperature of high temperature clusters from measurements of SZ relativistic corrections have been firstly obtained by Pointecouteau et al. (
+1998) and Hansen et al. (,1998) and Hansen et al. (
+2002).,2002).
+ The imaging anc spectroscopic capabilities of coming SZ experiments will permi us to analyze the gas temperature structure even in relatively coo galaxy clusters (Colafrancesco Marchegiani 2010)., The imaging and spectroscopic capabilities of coming SZ experiments will permit us to analyze the gas temperature structure even in relatively cool galaxy clusters (Colafrancesco Marchegiani 2010).
+ Using hydrodynamic simulations of galaxy clusters. Kay e al. (," Using hydrodynamic simulations of galaxy clusters, Kay et al. ("
+2008) compared the temperatures derived from the X-ray spectroscopy (the X-ray temperature) and from the SZ ettec (the SZ temperature).,2008) compared the temperatures derived from the X-ray spectroscopy (the X-ray temperature) and from the SZ effect (the SZ temperature).
+ They demonstrated that the SZ temperature corresponds to the Compton-averaged electron. temperature for relatively cool galaxy clusters., They demonstrated that the SZ temperature corresponds to the Compton-averaged electron temperature for relatively cool galaxy clusters.
+ Their SZ temperature map is irectly calculated from the simulated temperature and density oeistributions., Their SZ temperature map is directly calculated from the simulated temperature and density distributions.
+ Colafrancesco Marchegiani (2010) show how to obtain detailed information about the cluster temperature distribution for spherically-symmetric galaxy clusters., Colafrancesco Marchegiani (2010) show how to obtain detailed information about the cluster temperature distribution for spherically–symmetric galaxy clusters.
+ However. the approximation of a spherical symmetry cannot properly describe merging galaxy clusters.," However, the approximation of a spherical symmetry cannot properly describe merging galaxy clusters."
+ Signatures of shock waves caused by galaxy cluster mergers are imprinted on the 2D X-ray surface brightness maps of merging galaxy clusters (for a review. see Markevitch Vikhlinin 2007) and. therefore. merger shocks should be included in a realistic modeling of the SZ effect from merging galaxy clusters.," Signatures of shock waves caused by galaxy cluster mergers are imprinted on the 2D X-ray surface brightness maps of merging galaxy clusters (for a review, see Markevitch Vikhlinin 2007) and, therefore, merger shocks should be included in a realistic modeling of the SZ effect from merging galaxy clusters."
+ Prokhorov et al. (, Prokhorov et al. (
+010b) incorporated the relativistic Wright formalism for modeling the SZ etfect from a merging galaxy cluster in a numerical simulation and introduced,2010b) incorporated the relativistic Wright formalism for modeling the SZ effect from a merging galaxy cluster in a numerical simulation and introduced
+calibration.,calibration.
+ Iu order for the aforementioned programs to be successful. new data on a umuber of both polarized standards. at various degrees aud orientations of polarization. and unpolarized standards should be obtained.," In order for the aforementioned programs to be successful, new data on a number of both polarized standards, at various degrees and orientations of polarization, and unpolarized standards should be obtained."
+ Pointines should be eridded across the detector in order to investigate field depeudeuce., Pointings should be gridded across the detector in order to investigate field dependence.
+" A nunuber of exposures, iu cach polarizer. at different roll augles should also be enmploved iu order to replicate the eround characterization."," A number of exposures, in each polarizer, at different roll angles should also be employed in order to replicate the ground characterization."
+ This will allow the derivation of clement independent transimiüssiou co-efficieuts (frou point sources) and will allow tests for higher order calibration effects;, This will allow the derivation of element independent transmission co-efficients (from point sources) and will allow tests for higher order calibration effects.
+ Such thorough observations iav. also shed some light as to the reason for the change iu behavior from Cycles 7 aud 11., Such thorough observations may also shed some light as to the reason for the change in behavior from Cycles 7 and 11.
+ We have conducted. polarimetric analyses ὃν all Ginj)polavized standards available in the NICMOS data archive for the NIC2 camera., We have conducted polarimetric analyses on all (un)polarized standards available in the NICMOS data archive for the NIC2 camera.
+ The principle aim was to determine the behavior of the iustrunmient at low levels of polarization., The principle aim was to determine the behavior of the instrument at low levels of polarization.
+ We lave thoroughly tested ou routine for deriving aperture polarimetry of objects and found it to be robust., We have thoroughly tested our routine for deriving aperture polarimetry of objects and found it to be robust.
+ It has also demoustrated that observed simall radius variation in the polarimetric curves of erowtl cau be attributed to the effects of sub-pixel mis-aligumieuts between poluizing clemeuts iu the point spread function., It has also demonstrated that observed small radius variation in the polarimetric curves of growth can be attributed to the effects of sub-pixel mis-alignments between polarizing elements and the point spread function.
+ We have found uo evidence for persisteuco., We have found no evidence for persistence.
+ Our findines also indicate that there is a measure: polarization of pz1.2% iu unpolarized targets which may indicate an intrinsic mstrunental polarization auk the need for further tealàug of the on-orbit trausiission co-efficients., Our findings also indicate that there is a measured polarization of $p\approx1.2\%$ in unpolarized targets which may indicate an intrinsic instrumental polarization and the need for further tweaking of the on-orbit transmission co-efficients.
+ Asstuming we are detecting an Παππλοία effect we have corrected data for polarized targets in the (CQ.UC )-plaue and found the dispersion around previous (ground based) polarization estimates to have been reduced. but not totally removed.," Assuming we are detecting an instrumental effect we have corrected data for polarized targets in the $Q,U$ )-plane and found the dispersion around previous (ground based) polarization estimates to have been reduced, but not totally removed."
+ We have also derived averaged values of tf; that null the uupolarized staudard IID331591 across all apertures. aud found fj aud fo to ο (STILT and 05211. respectively. when fxius f4=1.9667.," We have also derived averaged values of $t_k$ that null the unpolarized standard HD331891 across all apertures, and found $t_1$ and $t_2$ to be 0.8717 and 0.8341, respectively, when fixing $t_3=0.9667$ ."
+ In addition. we have attempted to investigate the xossibilitv of an observationallv dependent iustrinenutal volarization. but are iuhibited from anv conclusions bv an insignificant sample.," In addition, we have attempted to investigate the possibility of an observationally dependent instrumental polarization, but are inhibited from any conclusions by an insignificant sample."
+ While such levels of residual intrinsic polarization may rot hamper studies of highlv polarized targets. these evels will have a detrimental effect ou studies attempting o nicasure p«ον," While such levels of residual intrinsic polarization may not hamper studies of highly polarized targets, these levels will have a detrimental effect on studies attempting to measure $p<5\%$."
+ It is clear that a more comprehensive calibration study of NICMOS is critical in order for a ΠΙΟ ofLST programs to be carried out successfully., It is clear that a more comprehensive calibration study of NICMOS is critical in order for a number of programs to be carried out successfully.
+" The current post-NCS calibration archive (one polarized. and one unpolarized standard) is insufficient for this to occur effectively,", The current post-NCS calibration archive (one polarized and one unpolarized standard) is insufficient for this to occur effectively.
+ An array of polarized aud unpolarized standards should be observed. at many roll angles aud iu all quadrants of the instruncut. iu order for the low level polarization characteristics of NICMOS to be properly and fully investigated. understood and removed.," An array of polarized and unpolarized standards should be observed, at many roll angles and in all quadrants of the instrument, in order for the low level polarization characteristics of NICMOS to be properly and fully investigated, understood and removed."
+ We would like to thank the referee for his careful and thorough reading of this mauuscript., We would like to thank the referee for his careful and thorough reading of this manuscript.
+ Was put has ereatly improved the paper., His input has greatly improved the paper.
+ Thanks are also extended to Cleun Schueider for his useful comments., Thanks are also extended to Glenn Schneider for his useful comments.
+ Support for this study was provided by NASA through a exaut from the Space Telescope Science Institute. which is operated bv the Association of Uiiversities for Research iu Astronomy. Incorporated. wader NASA contract. NAS5-," Support for this study was provided by NASA through a grant from the Space Telescope Science Institute, which is operated by the Association of Universities for Research in Astronomy, Incorporated, under NASA contract NAS5-26555."
+Following the approach of past authors. we construct the spectrum frou coalescing binary MDIIs frou: the galaxy merger rate. t1e black hole popuation deuographnies amougst galaxies. MDII binary dynamics and MDIT ünary eravitationalwave cussion.,"Following the approach of past authors, we construct the spectrum from coalescing binary MBHs from: the galaxy merger rate, the black hole population demographics amongst galaxies, MBH binary dynamics and MBH binary gravitational-wave emission."
+ Our galaxy merecr rate is based ou observations of close pairs aud au estimate of their dynamical friction time scale., Our galaxy merger rate is based on observations of close pairs and an estimate of their dynamical friction time scale.
+ We convert from the galaxy mass function to the black hole nass function using recent determijaflolis ο‘the correlation between black hole mass aud spheroic luass., We convert from the galaxy mass function to the black hole mass function using recent determinations of the correlation between black hole mass and spheroid mass.
+ Our 7fiducial” mode has rapid evolutioidu the norser rae per unit time por galaxw as a function of redslüft. arxl a constaut nass function of MDIIs out to :2Mi," Our “fiducial” model has rapid evolution in the merger rate per unit time per galaxy as a function of redshift, and a constant mass function of MBHs out to $z=3$."
+ The characteristic strain spectra predicte for this model. hf)~10lOcFvr1j2m35 is just below the atest observational limits at f~0.2vr1," The characteristic strain spectrum predicted for this model, $h_c(f)\sim10^{-16} (f/{\rm yr}^{-1})^{-2/3}$, is just below the latest observational limits at $f\sim0.2{\rm~yr}^{-1}$."
+ WheYOU t1e slope anc general character of his prediction agreo with the earlier work of ?.. our predicte auuplituce js somewhat üeher.," Whereas the slope and general character of this prediction agree with the earlier work of \citet{Rajagopal95}, our predicted amplitude is somewhat higher."
+ This increase is the result of the order of magnitude increase iu the loca AIDII αμνο deusitv. wuch is tempered by a somewhat lower herecr rate at hieh redshift.," This increase is the result of the order of magnitude increase in the local MBH number density, which is tempered by a somewhat lower merger rate at high redshift."
+ Wit1i our formulation we can calcilate other quautitics., With our formulation we can calculate other quantities.
+ T1ο number of biuaries coutributing for unit baavid hidf/f~lja ullz frequeucies is approximately 108., The number of binaries contributing for unit bandwidth $\delta f/f\sim1$ ) at nHz frequencies is approximately $10^6$.
+ It is evident frou our simulations that the variauce iu the strain spectrum in the same frequency band is roughly50%.. considerably largerOo than the 103 that would be expected if all evers were weighted equalvy.," It is evident from our simulations that the variance in the strain spectrum in the same frequency band is roughly, considerably larger than the $10^{-3}$ that would be expected if all events were weighted equally."
+" This incicates that the spread around the nea ris due to the possiημίν of a sual munber of high-ampliticο, nearby eveuts."," This indicates that the spread around the mean is due to the possibility of a small number of high-amplitude, nearby events."
+ The mecian redshift for he iuaries coutributiug to the spectrun is Lif the merger rate evolves relatively slowly with redshift (5XL2 j. but can be large for stroug evoution of the merger rate»," The median redshift for the binaries contributing to the spectrum is $\le 1$ if the merger rate evolves relatively slowly with redshift $\gamma\lesssim2$ ), but can be large for strong evolution of the merger rate."
+ The mean “chirp” mass is 3&LO°AL.. and he neal nass ratio is 11l: the individud amasses are then 10* alc LaeAL..., The mean “chirp” mass is $3\times 10^6 M_\odot$ and the mean mass ratio is $11$; the individual masses are then $10^7$ and $10^6~M_\odot$.
+ Our calculations allow us to sec «irectlv how the remaiune lack of understanding of some of these »plyvsical processes luupacts our predictions., Our calculations allow us to see directly how the remaining lack of understanding of some of these physical processes impacts our predictions.
+ These are areas for fuure research., These are areas for future research.
+ The most important factors are:, The most important factors are:
+determine a best-fitting CO(Q-1) redshift of 2.=4.755+0.001. corresponding to an offset of=370430kmss 1 with respect to the optical redshift derived by Vanzellaetal.(2006). of =4.762+0.002 determined from the asymmetric Lya emission. and a continuum break (see Coppinetal.2009 for the optical spectrum).,"determine a best-fitting CO(2–1) redshift of $z=4.755\pm0.001$, corresponding to an offset of $-370\pm30$ $^{-1}$ with respect to the optical redshift derived by \citet{Vanzella06} of $z=4.762\pm0.002$ determined from the asymmetric $\alpha$ emission, and a continuum break (see \citealt{Coppin09} for the optical spectrum)."
+" Note that this apparent redshifting of Lya compared to the systemic redshift of the source (as measured from the nebular emission) is routinely seen in samples of 2.2:22)3 star-forming galaxies with offsets of <Avis,cc450kkmss + compared to the systemic redshift (Steideletal.2010).", Note that this apparent redshifting of $\alpha$ compared to the systemic redshift of the source (as measured from the nebular emission) is routinely seen in samples of $2 \lsim z \lsim3$ star-forming galaxies with offsets of $<\! \Delta \mathrm{v}_{\mathrm{Ly}\alpha}\!>\simeq 450$ $^{-1}$ compared to the systemic redshift \citep{Steidel10}.
+. The CO redshift also corresponds to the redshift of the [OTT] emission at 2=4.751+0.005 from a VLT spectrum (Alaghband-Zadeh et iin preparation)., The CO redshift also corresponds to the redshift of the [OII] emission at $z=4.751\pm0.005$ from a VLT spectrum (Alaghband-Zadeh et in preparation).
+ This CO detection has provided definitive proof that the previously identified counterpart at >=4.76 for this SMG by Coppinetal.(2009). is correct., This CO detection has provided definitive proof that the previously identified counterpart at $z=4.76$ for this SMG by \citet{Coppin09} is correct.
+ Averaging the remaining channels in the ~1.6 (129-channel) line-free region of the spectrum reveals no GGHzsignificant continuum detection down to a o sensitivity of mmly., Averaging the remaining channels in the $\sim1.6$ GHz (129-channel) line-free region of the spectrum reveals no significant continuum detection down to a $\sigma$ sensitivity of mJy.
+ We calculate the line luminosity and the total cold gas mass (H2+He) from the integrated COQ(-1) line flux following SolomonVandenBout(2005). in order to determine the amount of available fuel for the current star formation episode., We calculate the line luminosity and the total cold gas mass $_{2}$ +He) from the integrated CO(2–1) line flux following \citet{Solomon05} in order to determine the amount of available fuel for the current star formation episode.
+" We find a line luminosity of Li4,=(2.040.4)107 kkmss.! ppc.", We find a line luminosity of $'_\mathrm{CO(2-1)}=(2.0\pm0.4)\times10^{10}$ $^{-1}$ $^{2}$ .
+" We then assume both that the gas is thermalised- for the .= 2— transition aassuming a constant line temperature brightness ratio. Log (οσο (1-00) and a CO-to-H» conversion factor of a= OSMM.(Kkkmss|! ppc) (which is appropriate for ultraluminous infrared galaxies and SMOs: e.g. Downes&Solomon 1998:: Taeconietal. 2008). yielding a total cold. gas mass of M.(1.6E0.3)«101"" MM..."," We then assume both that the gas is thermalised for the $J=2$ --1 transition assuming a constant line temperature brightness ratio, $'_\mathrm{CO}(2$ $'_\mathrm{CO}(1$ –0)) and a $_{2}$ conversion factor of $\alpha=0.8$ $_{\odot}$ $^{-1}$ $^{2}$ $^{-1}$ (which is appropriate for ultraluminous infrared galaxies and SMGs; e.g. \citealt{Downes98}; \citealt{Tacconi08}) ), yielding a total cold gas mass of $_\mathrm{gas}$ $(1.6\pm0.3)\times10^{10}$ $_\odot$."
+ This CO(Q-1) gas mass is comparableto that derived for SMGs from higher-./ transitions (typically 4= 3-2 or J= 4-3)., This CO(2–1) gas mass is comparableto that derived for SMGs from $J$ transitions (typically $J=3$ –2 or $J=4$ –3).
+ If JJ033229.4 is a typical SMG. these CO ./= 2-1 observations suggest that we are unlikely to be missing massive reservoirs of molecular gas around SMGs (see also e.g. Carillietal.2010: Ivisonetal. 2010: ef. al. 2006:: Papadopoulosetal. 20103).," If J033229.4 is a typical SMG, these CO $J=2$ –1 observations suggest that we are unlikely to be missing massive reservoirs of molecular gas around SMGs (see also e.g. \citealt{Carilli10}; \citealt{Ivison10}; cf. \citealt{Hainline06}; \citealt{Papa10}) )."
+ We calculate a dynamical mass of the system based on the observed CO(OC-1) line width following Solomon&VandenBout (2008): Masa ssin/2233.54Nejt R (M i. where i. is the inclination of the gas disk. Nery is the CO FWHM tin 1). and R is the radius of the CO emitting region (in pc).," We calculate a dynamical mass of the system based on the observed CO(2–1) line width following \citet{Solomon05}: $_\mathrm{dyn}$ $^{2}i$ $\Delta\,v^{2}_\mathrm{FWHM}\times$ R $_\odot$ ), where $i$ is the inclination of the gas disk, $\Delta\,v_\mathrm{FWHM}$ is the CO FWHM (in $^{-1}$ ), and R is the radius of the CO emitting region (in pc)."
+" Since we do not have a direct constraint on the spatial extent of the CO emission from our unresolved detection. we assume that the CO traces the stellar light (ας seems an appropriate assumption for ;~2 SMOs: Swinbanketal. 2010). which has a seeing-corrected FWHM of 0.5""XE0.2"" in the HAWK-I A-band imaging (or equivalently a FWHM = +tkkpe at 2= 4.76)."," Since we do not have a direct constraint on the spatial extent of the CO emission from our unresolved detection, we assume that the CO traces the stellar light (as seems an appropriate assumption for $z\sim2$ SMGs; \citealt{Swinbank10}) ), which has a seeing-corrected FWHM of $0.5''\pm0.2''$ in the HAWK-I $K$ -band imaging (or equivalently a FWHM $\simeq4$ kpc at $z=4.76$ )."
+ The best-fitting Gaussian FWHM of (160440) kkmss+ implies a strict lower limit on the dynamical mass of Mate2kpe) ssin?/2( 1.20.6).101 MM.. assuming a kkpe radius.," The best-fitting Gaussian FWHM of $(160\pm40)$ $^{-1}$ implies a strict lower limit on the dynamical mass of $_\mathrm{dyn}(<2\,\mathrm{kpc})$ $^{2}i$ $1.2\pm0.6)\times10^{10}$ $_\odot$, assuming a kpc radius."
+" Since we do not have any constraints on ὁ, in this letter we assume ¢=30° (appropriate for randomly inclined disks in a sample of galaxies) and explicitly include the 7 and R dependeney in our derived value: Miavul< .. "," Since we do not have any constraints on $i$, in this letter we assume $i=30^\circ$ (appropriate for randomly inclined disks in a sample of galaxies) and explicitly include the $i$ and R dependency in our derived value: $_\mathrm{dyn}(<2\,\mathrm{kpc})=(4.8\pm2.4)\times10^{10}\,(0.25/\mathrm{sin}^{2}i)(R/2\,\mathrm{kpc})$ $_\odot$."
+"We caution thatM, carries significant. uncertainties which will require higher-resolution observations to accurately determine these quantities.", We caution that$_\mathrm{dyn}$ carries significant uncertainties which will require higher-resolution observations to accurately determine these quantities.
+ For instance. we note that the CO line FWHM is about a factor of 3 narrower than typically observed for 2.~2 SMGs (Greveetal. 2005).. although there are examples of SMGs and high-redshift Lyman-break galaxies (LBGs) with CO emission as narrow as our observed FWHM (e.g. Frayeretal. 1999::," For instance, we note that the CO line FWHM is about a factor of 3 narrower than typically observed for $z\sim2$ SMGs \citep{Greve05}, , although there are examples of SMGs and high-redshift Lyman-break galaxies (LBGs) with CO emission as narrow as our observed FWHM (e.g. \citealt{Frayer99}; ;"
+We have identitied a small number of GRBs which display a period of time during which the X-ray emission shows a smooth plateau followed by a steep decline.,We have identified a small number of GRBs which display a period of time during which the X-ray emission shows a smooth plateau followed by a steep decline.
+ The internal plateau is challenging to interpret using accretion models as it requires a constant powerjet component with a roughly constant radiation efficiency., The internal plateau is challenging to interpret using accretion models as it requires a constant power jet component with a roughly constant radiation efficiency.
+ This possibility has been examined by Kumar (20082). who suggest that the prompt emission of a GRB may be caused by the accretion of the outer regions of a stellar core and that the X-ray plateau could be caused by the fall-back and accretion of the stellar envelope.," This possibility has been examined by Kumar (2008a), who suggest that the prompt emission of a GRB may be caused by the accretion of the outer regions of a stellar core and that the X-ray plateau could be caused by the fall-back and accretion of the stellar envelope."
+ This model has problems accounting for the steep declines seen after the plateau., This model has problems accounting for the steep declines seen after the plateau.
+ Even assuming a sharp edge to the region being accreted. the steepest decline expected is a~2.5 (Kumar et al.," Even assuming a sharp edge to the region being accreted, the steepest decline expected is $\alpha \sim 2.5$ (Kumar et al."
+ 2008b)., 2008b).
+ Here we argue that à more natural explanation may come from the magnetar model which predicts a period of constant spin-down power., Here we argue that a more natural explanation may come from the magnetar model which predicts a period of constant spin-down power.
+ This model starts with the assumption that the neutron star accretor can power the GRB prompt emission which while not certain. is feasible (Usov 1992: Thompson 1994: Bucciantini 22007).," This model starts with the assumption that the neutron star accretor can power the GRB prompt emission which while not certain, is feasible (Usov 1992; Thompson 1994; Bucciantini 2007)."
+ Comparison of the luminosity and duration of the internal plateaus observed in our GRB sample with the dipolar spindown law (Zhang Mésszárros 2001) implies upper limits to the magnetic field strengths close to the maximum allowed for such objects and initial spin periods also close to the maximum allowed to maintain neutron star structural integrity., Comparison of the luminosity and duration of the internal plateaus observed in our GRB sample with the dipolar spindown law (Zhang Mésszárros 2001) implies upper limits to the magnetic field strengths close to the maximum allowed for such objects and initial spin periods also close to the maximum allowed to maintain neutron star structural integrity.
+ The upper limits for the dipolar magnetic field of the magnetar are particularly strong if the emission is strongly beamed., The upper limits for the dipolar magnetic field of the magnetar are particularly strong if the emission is strongly beamed.
+ The largest magnetic fields implied for isotropic emission are consistent with field strengths of ©107G which can be generated in magnetars born with spin of a few milliseconds (Thompson Duncan 1993: Duncan 1998)., The largest magnetic fields implied for isotropic emission are consistent with field strengths of $\times 10^{16}$ G which can be generated in magnetars born with spin of a few milliseconds (Thompson Duncan 1993; Duncan 1998).
+ A giant flare from SGR 1806-20 on 27 December 2004 demonstrated that unless such flares are much rarer than the rate implied by detecting one. magnetars must possess a magnetic field strength of ~ 1010 or higher.," A giant flare from SGR 1806-20 on $27^{th}$ December 2004 demonstrated that unless such flares are much rarer than the rate implied by detecting one, magnetars must possess a magnetic field strength of $\sim 10^{16}$ G or higher."
+ Indeed values up to ~101 6 could not be ruled out (Stella 22005)., Indeed values up to $\sim 10^{17}$ G could not be ruled out (Stella 2005).
+ For the GRB sample in this paper this could allow beaming factors corresponding to jet opening angles of 4-10 degrees. consistent with values derived from Frail (2001).," For the GRB sample in this paper this could allow beaming factors corresponding to jet opening angles of 4-10 degrees, consistent with values derived from Frail (2001)."
+ The number of GRBs that display internal plateau behaviour is very small., The number of GRBs that display internal plateau behaviour is very small.
+ This perhaps is not surprising as we would expect them to only be detectable for quite a narrow combination of magnetic field strength and initial spin period., This perhaps is not surprising as we would expect them to only be detectable for quite a narrow combination of magnetic field strength and initial spin period.
+ These rare features do provide limits on the magnetic fields surrounding the central engine around the GRB. and ean help advance understanding of the mechanisms behind prompt emission.," These rare features do provide limits on the magnetic fields surrounding the central engine around the GRB, and can help advance understanding of the mechanisms behind prompt emission."
+ We gratefully acknowledge funding for aat the University of Leicester by the Science and Technology Facilities Council. in the USA by NASA and in Italy by contract ASVINAF I/088/06/0.," We gratefully acknowledge funding for at the University of Leicester by the Science and Technology Facilities Council, in the USA by NASA and in Italy by contract ASI/INAF I/088/06/0."
+ NL and RLCS also acknowledge funding by STFC via a studentship and PDRA respectively., NL and RLCS also acknowledge funding by STFC via a studentship and PDRA respectively.
+ BZ and NL acknowledge funding by NASA grants NNGOSGB67G and NNXOSAN24G. We are also very grateful to our colleagues on the pproject for their help and support. particularly Kim Page for providing the BAT," BZ and NL acknowledge funding by NASA grants NNG05GB67G and NNX08AN24G. We are also very grateful to our colleagues on the project for their help and support, particularly Kim Page for providing the BAT"
+a large Ni abundance. suggesting that the line is most likely not real).,"a large Ni abundance, suggesting that the line is most likely not real)."
+ Due to the low significance of these emission features. we consider their detection as tentative only. and we do not discuss them any further (though they are included in all subsequent fits).," Due to the low significance of these emission features, we consider their detection as tentative only, and we do not discuss them any further (though they are included in all subsequent fits)."
+ In all other objects an additional narrow emission line in the range 6.4 keV—6.97 keV (neutral to highly ionized Fe) does not improve the statistics significantly., In all other objects an additional narrow emission line in the range 6.4 keV–6.97 keV (neutral to highly ionized Fe) does not improve the statistics significantly.
+ However. a narrow («50 eV in σ) and neutral (76.4 keV) Fe emission line is ubiquitous in the spectra of AGN with typical equivalent width of ~100 eV (which is also anti-correlated with X-ray luminosity and Eddington ratio. see e.g. Bianchi et al 2007).," However, a narrow $<50$ eV in $\sigma$ ) and neutral $\sim$ 6.4 keV) Fe emission line is ubiquitous in the spectra of AGN with typical equivalent width of $\sim$ 100 eV (which is also anti–correlated with X–ray luminosity and Eddington ratio, see e.g. Bianchi et al 2007)."
+ We then include an unresolved Fe Ka emission line (with energy fixed at 6.4 keV) in our spectral models to compute upper limits on its equivalent width., We then include an unresolved Fe $\alpha$ emission line (with energy fixed at 6.4 keV) in our spectral models to compute upper limits on its equivalent width.
+ We obtain upper limits in the range of 190—550 eV in all objects. which are consistent with expectations from Bianchi et al. (," We obtain upper limits in the range of 190–550 eV in all objects, which are consistent with expectations from Bianchi et al. ("
+2007). given the low-luminosity of our objects (typically a few times 1077 erg s 1).,"2007), given the low–luminosity of our objects (typically a few times $10^{42}$ erg $^{-1}$ )."
+ Hence. we do not detect any neutral Fe Ka emission line at 6.4 keV. but the quality of our data is not high enough to rule out that a typical Fe line is present in the hard X-ray spectra of the sources.," Hence, we do not detect any neutral Fe $\alpha$ emission line at 6.4 keV, but the quality of our data is not high enough to rule out that a typical Fe line is present in the hard X–ray spectra of the sources."
+ We then extrapolate the hard power law model to soft energies O assess whether the hard spectral shape provides an adequate description of the broadband 0.3—10 keV data., We then extrapolate the hard power law model to soft energies to assess whether the hard spectral shape provides an adequate description of the broadband 0.3–10 keV data.
+ In Fig., In Fig.
+ | we show he result of this exercise which highlights the presence of a soft excess: the data below 1—2 keV lie in all cases above the extrapolation of the hard power law model., 1 we show the result of this exercise which highlights the presence of a soft excess: the data below 1–2 keV lie in all cases above the extrapolation of the hard power law model.
+ We then re-fitted he simple power law model described above in the whole 0.3—10 keV band and compare it with a (phenomenological) broken yower law model., We then re–fitted the simple power law model described above in the whole 0.3--10 keV band and compare it with a (phenomenological) broken power law model.
+ In Table 2 we report the spectral slopes in the soft (Po) and hard (4) band for the broken power law model only (the break energy is in the 1.5—3 keV range for all sources and is not reported)., In Table 2 we report the spectral slopes in the soft $\Gamma_{\rm s}$ ) and hard $\Gamma_{\rm h}$ ) band for the broken power law model only (the break energy is in the 1.5–3 keV range for all sources and is not reported).
+ The statistics for a single power law model is reported in parenthesis in the last column and its slope was ‘ound consistent with Ες within the errors due to the much higher signal-to-noise in the soft band., The statistics for a single power law model is reported in parenthesis in the last column and its slope was found consistent with $\Gamma_{\rm s}$ within the errors due to the much higher signal–to–noise in the soft band.
+ We also give the derived X-ray Huxes and luminosities in the Q.5-2 keV and in the 2-10 keV bands and the Eddington ratio LooLea where Lio=952100 (Greene Ho 2004).," We also give the derived X–ray fluxes and luminosities in the 0.5–2 keV and in the 2–10 keV bands and the Eddington ratio $L_{\rm Bol}/L_{\rm Edd}$ where $L_{\rm Bol} =
+9 L_{5100}$ (Greene Ho 2004)."
+ The spectral titting results show that the broken power law model is a better description of the broadband X-ray spectra of GH 1 (although at the 98.8 per cent level only). GH 8. and GH 12. while a single power law model is adequate or GH 14.," The spectral fitting results show that the broken power law model is a better description of the broadband X–ray spectra of GH 1 (although at the 98.8 per cent level only), GH 8, and GH 12, while a single power law model is adequate for GH 14."
+ In other words. a soft excess is detected in GH I. GH 8 and GH 12. and not significantly present in GH 14.," In other words, a soft excess is detected in GH 1, GH 8 and GH 12, and not significantly present in GH 14."
+ We note that the broken power law model is not entirely adequate o reproduce the soft excess and that results from such spectral tits should not be used to infer the significance of the soft excess detection., We note that the broken power law model is not entirely adequate to reproduce the soft excess and that results from such spectral fits should not be used to infer the significance of the soft excess detection.
+ A blackbody representation is a better. though also phenomenological. parametrization of the soft excess (see caption of Fig.," A blackbody representation is a better, though also phenomenological, parametrization of the soft excess (see caption of Fig."
+ | and Table 3)., 1 and Table 3).
+ We also point out that in no cases the addition of a neutral absorption component at the redshift of the source improved the statistical quality of the fit., We also point out that in no cases the addition of a neutral absorption component at the redshift of the source improved the statistical quality of the fit.
+ Hence. as already mentioned by Greene Ho (2007) in their analysis of the data. neutral intrinsic absorption covering the totality of the X-ray source plays a negligible role in these sources.," Hence, as already mentioned by Greene Ho (2007a) in their analysis of the data, neutral intrinsic absorption covering the totality of the X–ray source plays a negligible role in these sources."
+ The choice of the broken power law model allows to compare directly our results on the soft X-ray slope LL with previous observations of the GH sample (Greene Ho 2007a) which were limited to energies below ~3 keV due to the lack of signal-to-noise in the hard band (source GH 12 is however observed only with XAZM—Newton)., The choice of the broken power law model allows to compare directly our results on the soft X–ray slope $\Gamma_{\rm s}$ with previous observations of the GH sample (Greene Ho 2007a) which were limited to energies below $\sim$ 3 keV due to the lack of signal–to–noise in the hard band (source GH 12 is however observed only with ).
+ The soft spectral slopes are in good agreement with the previousChandra observations (Greene Ho 2007u) despite long-term flux variability: during the previous observations. GH | and GH 14. were 3 and 2 times brighter (respectively) than in the present data. while GH 8 was a factor |.4 fainter.," The soft spectral slopes are in good agreement with the previous observations (Greene Ho 2007a), despite long–term flux variability: during the previous observations, GH 1 and GH 14 were 3 and 2 times brighter (respectively) than in the present data, while GH 8 was a factor 1.4 fainter."
+ However. the photon indexes in the Q.5-2 keV band are consistent with each other within the errors.," However, the photon indexes in the 0.5–2 keV band are consistent with each other within the errors."
+ Moreover. the values of Pin Table 2 are also in line with the typical soft X-ray slope of luminous quasars (e.g. «Oo2.62.7 in the sample of PG quasars observed by (Porquet et al 2004: Piconcelli et al 2005).," Moreover, the values of $\Gamma_{\rm s}$ in Table 2 are also in line with the typical soft X–ray slope of luminous quasars (e.g. $<\Gamma_{\rm s}> \simeq 2.6-2.7$ in the sample of PG quasars observed by (Porquet et al 2004; Piconcelli et al 2005)."
+" Although affected by large error bars. the hard spectral slopes are remarkably similar within each other. and also consistent with the typical 2—10 keV slope of PG quasars (<Dy,c 1.9)."," Although affected by large error bars, the hard spectral slopes are remarkably similar within each other, and also consistent with the typical 2--10 keV slope of PG quasars $<\Gamma_{\rm h}>
+\simeq 1.9$ )."
+ The only exception is represented by the least luminous source (GH 12D. which has a flatter spectral slope in both the soft and the hard bands with respect to the other objects and. as mentioned. is also the only source for which a soft excess is not statistically required.," The only exception is represented by the least luminous source (GH 14), which has a flatter spectral slope in both the soft and the hard bands with respect to the other objects and, as mentioned, is also the only source for which a soft excess is not statistically required."
+ We conclude that the X-ray spectral properties of our small sample of intermediate—mass black hole AGN do not significantly differ from the average properties of luminous AGN powered by accretion onto more massive and more luminous black holes., We conclude that the X–ray spectral properties of our small sample of intermediate–mass black hole AGN do not significantly differ from the average properties of luminous AGN powered by accretion onto more massive and more luminous black holes.
+ The nature of the soft excess origin in Seyfert | galaxies is still unclear and matter of debate., The nature of the soft excess origin in Seyfert 1 galaxies is still unclear and matter of debate.
+ As pointed out by several authors. the idea that the soft excess in AGN represents the high energy tail of the quasi-blackbody emission from the accretion dise is not consistent with the observed properties in luminous samples (Czerny et al 2003: Gierlinsski Done 2004: Crummy et al 2006).," As pointed out by several authors, the idea that the soft excess in AGN represents the high energy tail of the quasi–blackbody emission from the accretion disc is not consistent with the observed properties in luminous samples (Czerny et al 2003; Gierlińsski Done 2004; Crummy et al 2006)."
+ The maximum temperature for a standard Shakura-Sunyaev thin dise and under the assuption that the energy release is dominated by thermal dise emission is achieved at the innermost dise radius ry and can be rewritten (e.g. Peterson 1997) as, The maximum temperature for a standard Shakura–Sunyaev thin disc and under the assuption that the energy release is dominated by thermal disc emission is achieved at the innermost disc radius $r_{\rm in}$ and can be rewritten (e.g. Peterson 1997) as
+motions one can estimate. e.g.. the expected mean position of the methanol masers at the observing epoch of the water masers. and we find that the displacement is only ~8 mas in Κ.Α. and «18 mas in Dec. Values that small are negligible with respect to the region of interest for us (2250 mas. see Fig. 3)).,"motions one can estimate, e.g., the expected mean position of the methanol masers at the observing epoch of the water masers, and we find that the displacement is only $\sim$ 8 mas in R.A. and $\sim$ 18 mas in Dec. Values that small are negligible with respect to the region of interest for us $\ga$ 250 mas, see Fig. \ref{fdj}) )."
+ Moseadelli et al. (2005)), Moscadelli et al. \cite{mosca05}) )
+ interpreted the mmasers as tracing the surface of a conical jet. expanding away from the star.," interpreted the masers as tracing the surface of a conical jet, expanding away from the star."
+ The model could satisfactorily fit the observed line-of-sight velocities and absolute proper motions. after correcting for the parallax. motion of the Sun with respect to the LSR. and galactic rotation.," The model could satisfactorily fit the observed line-of-sight velocities and absolute proper motions, after correcting for the parallax, motion of the Sun with respect to the LSR, and galactic rotation."
+ Knowledge of the star's 3D velocity. obtained in Sect. 5..," Knowledge of the star's 3D velocity, obtained in Sect. \ref{svsys},"
+ permits accurate estimation of the mmaser velocities with respect. to the star., permits accurate estimation of the maser velocities with respect to the star.
+ Using this information. we could correct the proper motions of the mmaser features derived by Moscadelli et al. (2005)).," Using this information, we could correct the proper motions of the maser features derived by Moscadelli et al. \cite{mosca05}) )."
+ The new velocity vectors are shown as arrows in Fig. 3..., The new velocity vectors are shown as arrows in Fig. \ref{fdj}.
+ Comparison with Fig., Comparison with Fig.
+ 2 of Moseadelli et al. (2005)), 2 of Moscadelli et al. \cite{mosca05}) )
+ shows that the correction does not affect significantly the magnitudes and directions of the vectors. apart from the red-shifted feature to the SE.," shows that the correction does not affect significantly the magnitudes and directions of the vectors, apart from the red-shifted feature to the SE."
+ This had negligible proper motion in the old analysis of the data. whereas with the new correction it appears to move towards SE. lending further support to the bipolar jet model of the mmasers.," This had negligible proper motion in the old analysis of the data, whereas with the new correction it appears to move towards SE, lending further support to the bipolar jet model of the masers."
+ Note that the proper motions obtained in the present study for the two features selected for our parallax measurement in Sect., Note that the proper motions obtained in the present study for the two features selected for our parallax measurement in Sect.
+ 4. (-46 in R.A. and 11 s!in Dec. for the feature with LSR velocity of 23.2s!.. and —47 lin Κ.Α. and 17 s'in Dec. for that with LSR velocity of —6.4 s) are consistent with the jet-model proposed by Moscadelli et al. (2005)).," \ref{spara} (–46 in R.A. and 11 in Dec. for the feature with LSR velocity of –3.2, and –47 in R.A. and 17 in Dec. for that with LSR velocity of –6.4 ) are consistent with the jet-model proposed by Moscadelli et al. \cite{mosca05}) )."
+ We have not derived the proper motions of the other features detected in the present experiment. because these will be the subject of a forthcoming paper illustrating the results of à VLBI multi-year monitoring of the water maser emission I12012644104.," We have not derived the proper motions of the other features detected in the present experiment, because these will be the subject of a forthcoming paper illustrating the results of a VLBI multi-year monitoring of the water maser emission in."
+ We used the model of Moscadelli et al. (2005)), We used the model of Moscadelli et al. \cite{mosca05}) )
+ to fit the new. corrected data and obtained new best-fit parameters that are qualitatively consistent with the old ones.," to fit the new, corrected data and obtained new best-fit parameters that are qualitatively consistent with the old ones."
+ We find a position angle (P.A.) of instead of1237.. an inclination with respect to the line of sight of instead of96°.. an opening angle of instead of17°. and a velocity gradient along the jet of 0.116 mas?! instead of 0.255 mas7!.," We find a position angle (P.A.) of instead of, an inclination with respect to the line of sight of instead of, an opening angle of instead of, and a velocity gradient along the jet of 0.116 $^{-1}$ instead of 0.255 $^{-1}$."
+ Also. the position of the star is shifted by +143 mas in right ascension and —49 mas in declination.," Also, the position of the star is shifted by +143 mas in right ascension and –49 mas in declination."
+ In conclusion. despite some differences between the old and new best-fit parameters. the corrected data confirm that the jet is very beamed. Hes close to the plane of the sky. and is approximately perpendicular to the disk (whose P.A. is ~534T: see Cesaroni et al. 2005)).," In conclusion, despite some differences between the old and new best-fit parameters, the corrected data confirm that the jet is very beamed, lies close to the plane of the sky, and is approximately perpendicular to the disk (whose P.A. is $\sim53\pm7\degr$; see Cesaroni et al. \cite{cesa05}) )."
+ So far. we have used the mmasers only to determine the systemic velocity of the startdisk system.," So far, we have used the masers only to determine the systemic velocity of the star+disk system."
+ Now. we analyse their role in tracing the distribution and kinematics of the gas in the disk-Jet system.," Now, we analyse their role in tracing the distribution and kinematics of the gas in the disk-jet system."
+ Looking at Fig. 3..," Looking at Fig. \ref{fdj},"
+" one sees that the mmaser features are roughly outlining two linear structures. one oriented NE-SW and the other SE-NW: hereafter we will refer to these as. respectively. ""group 1"" and ""group 27."," one sees that the maser features are roughly outlining two linear structures, one oriented NE–SW and the other SE–NW; hereafter we will refer to these as, respectively, “group 1” and “group 2”."
+ The features belonging to these groups are (using the labels in Table 1) 2. 3. 5. 6. 7. 10. 1I. 12. 13. 15. 18. 20. 24. 26. 27. 28. 32. 33 for group I and |. 4. 8. 9. 14. 16. 17. 19. 21. 22. 23. 25. 29. 30. 31 for group 2.," The features belonging to these groups are (using the labels in Table \ref{tmeth}) ): 2, 3, 5, 6, 7, 10, 11, 12, 13, 15, 18, 20, 24, 26, 27, 28, 32, 33 for group 1 and 1, 4, 8, 9, 14, 16, 17, 19, 21, 22, 23, 25, 29, 30, 31 for group 2."
+ Comparison with Edris et al. (2005)), Comparison with Edris et al. \cite{edris}) )
+ lends support to this distinction. because both the spots positions and velocities appear consistent with those of our groups ] and 2. with only one exception (spot number 15 in their Table 6).," lends support to this distinction, because both the spots positions and velocities appear consistent with those of our groups 1 and 2, with only one exception (spot number 15 in their Table 6)."
+ We have calculated the proper motions (see Fig. 3)), We have calculated the proper motions (see Fig. \ref{fdj}) )
+4 of the mmaser features of both groups with respect to feature 3. which belongs to group ] and appears to be the most reliable for the persistency of the internal structure over the three observing epochs.," of the maser features of both groups with respect to feature 3, which belongs to group 1 and appears to be the most reliable for the persistency of the internal structure over the three observing epochs."
+ It is worth noting two facts., It is worth noting two facts.
+" First. the masers in group 2 move towards NW, Le. in the same direction as the blue lobe of the mmaser jet."," First, the masers in group 2 move towards NW, i.e. in the same direction as the blue lobe of the maser jet."
+ Despite the different speeds (up to  100 for aand only +5 ffor CH3OH)). this suggests a relationship between the two.," Despite the different speeds (up to $\sim$ 100 for and only $\sim$ 5 for ), this suggests a relationship between the two."
+ Second. the mean velocity of group | features along the line of ⋋⇪∶↔⊺∣↴⊓⋋−↻↜≼⊔∖⊳∏⊔⊣⋅⋅∖∖⇁∣↴≣∁∣∐∐⇈∖⊜∣⋪⋋⇂↴⋝⇁−∍↜−⊔∖⊳∏⊔≓ ∖∖⇁∖∖⇁⇪⋔∣⋪⊜⋋∣⊃⊖∁⊓⋂⋔⊜⋋⋝⇁⋋↾⊜⋯≣∁∟⊱↾≿∖⇁⊖∣," Second, the mean velocity of group 1 features along the line of sight is –6.9, which differs by –3.4 with respect to the systemic LSR velocity of –3.5."
+∪∁⇈⋝⇁⋂↑↴−∍⇤⊰↥∖⊳∏⊔⊣↣ The mean distance of the same features from the nominal position of the star (starred symbol in Fig. 3)), The mean distance of the same features from the nominal position of the star (starred symbol in Fig. \ref{fdj}) )
+ is 07227 or ~460 AU., is 27 or $\sim$ 460 AU.
+ If the masers are undergoing Keplerian rotation about the 7 M.. star. the expected velocity is ~—3.7kms. consistent with the value of 23.4 qquoted above.," If the masers are undergoing Keplerian rotation about the 7 $M_\odot$ star, the expected velocity is $\sim-3.7$, consistent with the value of –3.4 quoted above."
+ Based on this. we propose that group 1 masers trace the Keplerian disk (as proposec by Minier et al.," Based on this, we propose that group 1 masers trace the Keplerian disk (as proposed by Minier et al."
+ 2001. and Edris et al. 2005))., \cite{mini01} and Edris et al. \cite{edris}) ).
+ As for group 2 masers we speculate that they also lie in the disk. but very close to the transition region between the disk surface and the outer part of the jet.," As for group 2 masers we speculate that they also lie in the disk, but very close to the transition region between the disk surface and the outer part of the jet."
+ Although based on admittedly marginal evidence. this hypothesis may explain why the relative proper motions of these features (see Fig. 3))," Although based on admittedly marginal evidence, this hypothesis may explain why the relative proper motions of these features (see Fig. \ref{fdj}) )"
+ are roughly perpendicular to the plane of the disk., are roughly perpendicular to the plane of the disk.
+ Besides. we will show in the following that also the line-of-sight velocities of group 2 masers (whose mean value is —-5.9 )) differ from those of group | masers lending support to our speculation.," Besides, we will show in the following that also the line-of-sight velocities of group 2 masers (whose mean value is –5.9 ) differ from those of group 1 masers lending support to our speculation."
+ Finally. we note that the number of maser spots without measurable proper motions is larger in group 2 (40%)) then in group | (22%)).," Finally, we note that the number of maser spots without measurable proper motions is larger in group 2 ) then in group 1 )."
+ This suggests that masers at the jet interface might be more persistent than those in the disk plane. consistent with the idea that radiative," This suggests that masers at the jet interface might be more persistent than those in the disk plane, consistent with the idea that radiative"
+it at some characteristic size of the svstem.,it at some characteristic size of the system.
+ So doing. the ellects of the collective dressing of the particles are still retained. albeit less precisely. but the elfects of the structure of the svstem are only sketchily accounted for.," So doing, the effects of the collective dressing of the particles are still retained, albeit less precisely, but the effects of the structure of the system are only sketchily accounted for."
+ 1n this limit. the svstem is regarded as homogencously filling a laree cubix box of side £L. on the surface of which periodic boundary conditions apply.," In this limit, the system is regarded as homogeneously filling a large cubix box of side $L$, on the surface of which periodic boundary conditions apply."
+ This is the geometry considered by Weinberg(1993)., This is the geometry considered by \citet{Weinberg93}.
+". Due to the assumed homogencity. the collective force FE"" vanishes and the unperturbed motion is rectilinear and uniform. whatever the state of relaxation of the system."," Due to the assumed homogeneity, the collective force $\mathbf{F}^0$ vanishes and the unperturbed motion is rectilinear and uniform, whatever the state of relaxation of the system."
+ The action variables are then proportional to the components of the momentum p and the angle variables are proportional to the components of the position r., The action variables are then proportional to the components of the momentum $\mathbf{p}$ and the angle variables are proportional to the components of the position $\mathbf{r}$.
+ Since the angles must be variables of period 2x. the angle vector must be w=2gr/L. which implies that the action vector is J=Lp/2s.," Since the angles must be variables of period $2 \pi$ , the angle vector must be ${\mathbf{w}} = 2 \pi {\mathbf{r}} /L$, which implies that the action vector is ${\mathbf{J}} = L {\mathbf{p}}/2\pi$."
+ The angle Fourier vector is k=LI&/2s where KK is the usual wave vector of Fourier transforms with respect to position., The angle Fourier vector is ${\mathbf{k}} = L {\mathbf{K}} /2\pi$ where ${\mathbf{K}}$ is the usual wave vector of Fourier transforms with respect to position.
+ The frequeney. vector Q is 2zv/L. so that k-O22K«v.," The frequency vector ${\mathbf{\Omega}}$ is $2\pi {\mathbf{v}}/L$, so that ${\mathbf{k}} \cdot {\mathbf{\Omega}} = {\mathbf{K}} \cdot {\mathbf{v}}$."
+ The density-potential basis consists of functions proportional to complex exponentials. like exp(?Ir).," The density-potential basis consists of functions proportional to complex exponentials, like $\exp(i\,{\mathbf{K}}\! \cdot\! {\mathbf{r}})$ ."
+ A given clement of the basis. à say. is characterized by its wave vector K.," A given element of the basis, $\alpha$ say, is characterized by its wave vector ${\mathbf{K}}$."
+ This can be accounted for in the notation by writing this wave vector as IK. the corresponding angle wave vector being noted Κω.," This can be accounted for in the notation by writing this wave vector as ${\mathbf{K}}_\alpha$, the corresponding angle wave vector being noted ${\mathbf{k}}_\alpha$."
+" The density function and the potential of the element a of the basis are both proportional toexpCK,.-r)."," The density function and the potential of the element $\alpha$ of the basis are both proportional to$\exp{(i\, {\mathbf{K}}_\alpha\! \cdot\! {\mathbf{r}}) }$."
+" Their normalization factor must be such that the biorthogonality relation (22)) be satisfied. the density 2°Gr) and the potential c""(r) being related by equation (21))."," Their normalization factor must be such that the biorthogonality relation \ref{biortho}) ) be satisfied, the density $D^\alpha({\mathbf{r}})$ and the potential $\psi^\alpha( {\mathbf{r}})$ being related by equation \ref{DeltaetPsi}) )."
+" These constraints result in: ↾∐↕∢⊾∣⋎∃≺∆↼⊳∖⋜⊔⋅⋖⊾∣↓↕⋖⊾↓⊲↸∪⊔↓⋰⊓⋅↓⋅⇂↓⋅⋜⋯⊳∖⇂⋅∪↓⋅⊔↓⊳∖∪⇂⋅∣⋎⊔↿∖↕↘∃∖∖⋰↓↿↓⊔⋅⋖⊾⊳���↓≻⋖⋅≼⇍↿↿∪↿↓↕∢⋅⋜⋯⋏∙≟↓∢⊾⊳∖∇∇⊳⊔⋜⋯↓∢⊾↓∙∖⇁∶ where 9(k,—K) is a triple Kronecker svmbol."," These constraints result in: The $\psi^\alpha_{ \mathbf{k} }$ 's are the Fourier transforms of $\psi^\alpha({\mathbf{r}})$ with respect to the angles $\mathbf{w}$, namely: where $\delta({\mathbf{k}}_\alpha - {\mathbf{k}})$ is a triple Kronecker symbol."
+ In this case the oy os do not depend on the actions and remain fixed while the relaxation proceeds., In this case the $\psi^\alpha_{\mathbf{k}}$ 's do not depend on the actions and remain fixed while the relaxation proceeds.
+ Ehe response matrix 5. caleulated from its definition (34)). is diagonal. its element aa being given. for w in the upper half complex plane. by: For real &. à. |70 should be added to the singular denominator.," The response matrix $\varepsilon$, calculated from its definition \ref{epsilonalphabeta}) ), is diagonal, its element $\alpha\alpha$ being given, for $\omega$ in the upper half complex plane, by: For real $\omega$, a $+i0$ should be added to the singular denominator."
+" Since o enters this relation by its wave vector I. 20""(o) can be regarded as a function z(I&,ic). or. more generally. as a function of a wave vector I& and of a frequeney w."," Since $\alpha$ enters this relation by its wave vector $\mathbf{K}_\alpha$, $\varepsilon^{\alpha\alpha}(\omega)$ can be regarded as a function $\varepsilon({\mathbf{K}}_\alpha, \omega)$, or, more generally, as a function of a wave vector $\mathbf{K}$ and of a frequency $\omega$."
+" Because of the diagonalitv of the response matrix 5 and the simplicity of equation. (45)). the writing of the kinetic equation (38)) simplifies to: where the tensor Q,,D is delined by: Equations (47)) (48)) are identical to equation. (29) of Weinberg(1993) when the quasi homogeneity of the system and associated: absence of collective ancl external forces are accounted for."," Because of the diagonality of the response matrix $\varepsilon$ and the simplicity of equation \ref{psialphakhomograv}) ), the writing of the kinetic equation \ref{LAequation}) ) simplifies to: where the tensor ${\overline{\overline{Q}}}_{ab}$ is defined by: Equations \ref{LAequationhomograv}) \ref{colltensorhomograv}) ) are identical to equation (29) of \citet{Weinberg93} when the quasi homogeneity of the system and associated absence of collective and external forces are accounted for."
+" Equation (47)) can be written explicitly as a Fokker-Planck equation in the form: where the momentum drag and diffusion coellicients are: For electrical instead. of gravitational interactions. the gravitational constant € should be replaced by απο, in AINSA units. ¢ being the dielectric permittivity of vacuum."," Equation \ref{LAequationhomograv}) ) can be written explicitly as a Fokker-Planck equation in the form: where the momentum drag and diffusion coefficients are: For electrical instead of gravitational interactions, the gravitational constant $G$ should be replaced by $1/4\pi \epsilon_o$ in MKSA units, $\epsilon_o$ being the dielectric permittivity of vacuum."
+ The electric force between like charges being repulsive instead: of attractive. the minus sien before the second term of equation (46)) should be changed to a positive sign and the masses replaced by the charges of the particles.," The electric force between like charges being repulsive instead of attractive, the minus sign before the second term of equation \ref{epsilonhomograv}) ) should be changed to a positive sign and the masses replaced by the charges of the particles."
+ Equation (47)) then reduces to the Balescu-Lenard equation forhomogeneous ancl multispecies plasmas (BabuelPevrissac 1974)., Equation\ref{LAequationhomograv}) ) then reduces to the Balescu-Lenard equation forhomogeneous and multispecies plasmas \citep{Babuel}. .
+. Ht implicitly accounts for the screening ellect. which is embodied in the dielectric function.," It implicitly accounts for the screening effect, which is embodied in the dielectric function."
+ The integralS on wave vector space in equationi (48)) then need not be cut at small wave vectors because τῆς.Iv)| diverges," The integral on wave vector space in equation \ref{colltensorhomograv}) ) then need not be cut at small wave vectors because $\mid\! \varepsilon({\mathbf{K}}, {\mathbf{K}}\!\cdot {\mathbf{v}})\!\mid$ diverges"
+abundances used to construct the models (see 113 and 88.1 of Paper I).,abundances used to construct the models (see 13 and 8.1 of Paper I).
+" In reffig:HR,, models are compared to three observed binary systems: V380 Ori (m), HD72106 (A) and RS Cha (e), which are all believed to be HAeBe stars."," In\\ref{fig:HR}, models are compared to three observed binary systems: V380 Ori $\blacksquare$ ), HD72106 $\blacktriangle$ ) and RS Cha $\bullet$ ), which are all believed to be HAeBe stars."
+" For the Ae star V380 Ori (Alecianetal. 2009)), the primary and chemically normal secondary are compared to the 2.80W5E-14 and models respectively."," For the Ae star V380 Ori \citealt{alecian09}) ), the primary and chemically normal secondary are compared to the 2.80W5E-14 and models respectively."
+" Within the error bars, the fit in the HR diagram is very good and the corresponding age of the primary is ~2.7 MMyr (the age given in Alecianetal.2009 is 221 MMyr)."," Within the error bars, the fit in the HR diagram is very good and the corresponding age of the primary is $\sim$ Myr (the age given in \citealt{alecian09}
+ is $\pm 1$ Myr)."
+ The 2.80W5E-14 model has developed overabundances of Fe and Ca which reach ~0.3 ddex and ~0.6ddex respectively., The 2.80W5E-14 model has developed overabundances of Fe and Ca which reach $\sim 0.3$ dex and $\sim$ dex respectively.
+" Though this model could only be converged until MMyr, it has already developed anomalies which explain the high metallicity ((M/H]=0.5) determined in Alecianetal.(2009)."," Though this model could only be converged until Myr, it has already developed anomalies which explain the high metallicity ([M/H]=0.5) determined in \citet{alecian09}."
+". As illustrated in reffig:V380,, most metals heavier than Z>15 are overabundant by a factor~2—3,, while CNO are barely underabundant."," As illustrated in \\ref{fig:V380}, most metals heavier than $Z \geq 15$ are overabundant by a factor, while CNO are barely underabundant."
+ Lithium becomes —0.2 ddex underabundant., Lithium becomes $-$ dex underabundant.
+" At the best fitting age for the primary, the 1.50W5.14 model still has its initial abundances, which also agrees with observations."," At the best fitting age for the primary, the 1.50W5.14 model still has its initial abundances, which also agrees with observations."
+" The potential effect of a magnetic field, which is observed in the primary, will be discussed in refsec:conclusions.."," The potential effect of a magnetic field, which is observed in the primary, will be discussed in \\ref{sec:conclusions}."
+" For 772106, the primary and the secondary are compared to the 2.50W5E-14 and 1.90W5E-14 models respectively."," For 72106, the primary and the secondary are compared to the 2.50W5E-14 and 1.90W5E-14 models respectively."
+" In the HR diagram, the fit is not perfect, especially for the primary, though this was also problematic in 44 of Folsometal.(2008)."," In the HR diagram, the fit is not perfect, especially for the primary, though this was also problematic in 4 of \citet{folsom08}."
+". In fact, these authors suggest that at a determined age of 6-13 MMyr (our best fit model is between MMyr), the primary is most likely on the ZAMS rather than the PMS."," In fact, these authors suggest that at a determined age of $-$ Myr (our best fit model is between Myr), the primary is most likely on the ZAMS rather than the PMS."
+" Nonetheless, they found that the primary was chemically anomalous with important overabundances of iron peak elements and a strong underabundace of He, while the secondary is almost solar."," Nonetheless, they found that the primary was chemically anomalous with important overabundances of iron peak elements and a strong underabundace of He, while the secondary is almost solar."
+" In the best fitting models for the primary MMyr), iron overabundances vary from about 0.1 to 0.3ddex, while Ca overabundances vary from about ddex down to about ddex."," In the best fitting models for the primary Myr), iron overabundances vary from about 0.1 to dex, while Ca overabundances vary from about dex down to about dex."
+" For the same age interval, the He underabundance varies from —0.1 to —0.15 ddex."," For the same age interval, the He underabundance varies from $-$ 0.1 to $-$ dex."
+ These amplitudes are smaller than the observed values., These amplitudes are smaller than the observed values.
+ This could be due to the mass loss rate being too large (see reffig:absurf1.9))., This could be due to the mass loss rate being too large (see \\ref{fig:absurf1.9}) ).
+ The presence of a magnetic field and of phase variations (see Fig.99 of Folsometal.2008)) do not justify trying to achieve a better fit., The presence of a magnetic field and of phase variations (see 9 of \citealt{folsom08}) ) do not justify trying to achieve a better fit.
+ A precise model would require 2 or DD calculations., A precise model would require 2 or D calculations.
+" Furthermore, some uncertainty remains on the exact amplitudes since all their abundances were determined simultaneously by fitting the observed spectra for a single averaged phase."," Furthermore, some uncertainty remains on the exact amplitudes since all their abundances were determined simultaneously by fitting the observed spectra for a single averaged phase."
+" Over this same age interval, the secondary still has its initial abundances, which agrees with observations."," Over this same age interval, the secondary still has its initial abundances, which agrees with observations."
+" For RS Cha, the primary and the secondary may be compared to the 1.90W5E-14 model."," For RS Cha, the primary and the secondary may be compared to the 1.90W5E-14 model."
+" In the HR diagram, the slight discrepancy with the model can be explained by the slight difference in stellar mass."," In the HR diagram, the slight discrepancy with the model can be explained by the slight difference in stellar mass."
+" Within reasonable error bars, the model is either chemically normal (up to MMyr), or has developed an overabundance of Ca of about ddex accompanied by small underabundances of TLi and He of —0.11 ddex and —0.05 ddex respectively."," Within reasonable error bars, the model is either chemically normal (up to Myr), or has developed an overabundance of Ca of about dex accompanied by small underabundances of $^7$ Li and He of $-$ dex and $-$ dex respectively."
+ This model does not seem to explain the iron enrichment factor of 1.5 obtained by the authors., This model does not seem to explain the iron enrichment factor of 1.5 obtained by the authors.
+" However, in contrast to the two previous binary systems, both these stars have the same composition; therefore, there is no difference in composition between the stars to explain, and so the initial abundances could be responsable for the abundance anomalies with respect to the solar composition."," However, in contrast to the two previous binary systems, both these stars have the same composition; therefore, there is no difference in composition between the stars to explain, and so the initial abundances could be responsable for the abundance anomalies with respect to the solar composition."
+ The X-ray emission observations of Mamajeketal.(1999) suggest that accretion could play a role., The X-ray emission observations of \citet{mamajek99} suggest that accretion could play a role.
+ A smaller mass loss rate would also lead to larger anomalies., A smaller mass loss rate would also lead to larger anomalies.
+ Other young single HAeBe stars which are not shown in refüg:HR may also be compared to our models., Other young single HAeBe stars which are not shown in \\ref{fig:HR} may also be compared to our models.
+" The star 1104237 has a mass of about (Teg=8000 KK), a luminosity of about 1.42LLo (Bohmetal. 2004)) and an approximate age of ((vandenAnckeretal. 1998)), as well as a magnetic field of about GG 2000)."," The star 104237 has a mass of about $\teff=8000$ K), a luminosity of about $_{\odot}$ \citealt{bohm04}) ) and an approximate age of \citealt{vandenAncker98}) ), as well as a magnetic field of about G \citep{donati97,donati00}."
+". Itis an HAeBe star for which Acke&Waelkens found that Si, Cr and Fe abundances were solar, which agrees with our results since even the heavier model is roughly normal until MMyr (see reffig:HR))."," It is an HAeBe star for which \citet{acke04} found that Si, Cr and Fe abundances were solar, which agrees with our results since even the heavier model is roughly normal until Myr (see \\ref{fig:HR}) )."
+" Similarly, 1190073, which has an age of 1.2+0.6 MMyr and a mass of 2.85+0.25 (Te= 9250K), was also found to be roughly solar (Acke&Waelkens 2004))."," Similarly, 190073, which has an age of $\pm$ Myr and a mass of $\pm$ $\teff=9250\,$ K), was also found to be roughly solar \citealt{acke04}) )."
+" Within the timescales shown in reffig:HR,, this is also compatible with our results."," Within the timescales shown in \\ref{fig:HR}, , this is also compatible with our results."
+" Finally, though it has just recently embarked on the MS, the young"," Finally, though it has just recently embarked on the MS, the young"
+us purpose.,this purpose.
+ We have tested some of the most popular algorithms. which cau )o divided in two main categories: the eradieut methods (CAL) aud the least squares methods (LS) In the eradicut based methods we use au estimate of ie cost function at the oth step which is based ou the wth data input. while ιο updating criterion for the learning parameter is derived by uunimizine the value of the cost function.," We have tested some of the most popular algorithms, which can be divided in two main categories: the gradient methods (GAL) and the least squares methods (LS) In the gradient based methods we use an estimate of the cost function at the $n$ th step which is based on the $n$ th data input, while the updating criterion for the learning parameter is derived by minimizing the value of the cost function."
+ Ou the other haud. in the least squares sed methods the optimal least squares prediction is computed at every poit oei time keeping iuto account the whole data history.," On the other hand, in the least squares based methods the optimal least squares prediction is computed at every point in time keeping into account the whole data history."
+ We report iu figure 2. the results obtained with the Least Squares Lattice (LSL) filter which is the best working amone the algorithims we analyzed., We report in figure \ref{fig:lsl} the results obtained with the Least Squares Lattice (LSL) filter which is the best working among the algorithms we analyzed.
+ It is evident the efficiency. of this algorithm in following all the features of the noise Spectruni., It is evident the efficiency of this algorithm in following all the features of the noise spectrum.
+ ποσα The fast convergence of the algoritlun lets us follow non stationarity of the roise Which are slower than one minute., -1truecm The fast convergence of the algorithm lets us follow non stationarity of the noise which are slower than one minute.
+ To check the quality of an estimator we need to put a bound on its performance., To check the quality of an estimator we need to put a bound on its performance.
+ We cau use general results frou he theory of statistical estimators., We can use general results from the theory of statistical estimators.
+ The variance of auv unbiased estimator Is voted. from below by the Cramer-Rao bound., The variance of any unbiased estimator is bounded from below by the Cramer-Rao bound.
+ We checked if the variance of he AR parameters estimated with the LSL estimator attains these limits or rot., We checked if the variance of the AR parameters estimated with the LSL estimator attains these limits or not.
+ We have computed the variance of cach paraucter aud of 07 at different nues to follow how the variance approaches the CR bound as time goes by., We have computed the variance of each parameter and of $\sigma^2$ at different times to follow how the variance approaches the CR bound as time goes by.
+ The times are delayed cach other by 8 secouds. the last time correspouding to oue elapsed minute.," The times are delayed each other by $8$ seconds, the last time corresponding to one elapsed minute."
+thermal spectra. N-rav observations are most scusitive to temperatures Tofew «109 IK. which are cousidered to be typical temperatures for MSP. polar caps. heated by relativistic particles.,"thermal spectra, X-ray observations are most sensitive to temperatures $T\sim$ $\times10^6$ K, which are considered to be typical temperatures for MSP polar caps, heated by relativistic particles."
+" às an MSP (P= Sims) in a binary system with ai white dwarf (WD). with orbital period 5.5dd. ? measured the proper motion of1715.. 1354 Enunas/vr and detected an IL, bow- nebula iu the approximate direction of motion: later ? ereatly improved the accuracy of the proper motion /—ineasurenients usiug VLBI (μμ=121.679+ 0.0521unas/vY. py=τισος-0.086 παςντ). and provided the most accurate distance to1715.. d=1950 ppc (ualking it one of the closest pulsars known)."," is an MSP $P=5.8$ ms) in a binary system with a white dwarf (WD), with orbital period d. \citet{1995ApJ...440L..81B} measured the proper motion of, $\pm$ mas/yr and detected an $_\alpha$ bow-shock nebula in the approximate direction of motion; later \cite{2008ApJ...685L..67D} greatly improved the accuracy of the proper motion measurements using VLBI $\mu_\alpha=121.679\pm0.052$ mas/yr, $\mu_\delta=-71.820\pm0.086$ mas/yr), and provided the most accurate distance to, $d=156.3\pm$ pc (making it one of the closest pulsars known)."
+ ΕΙThe masses of the two binary conrponeuts are Mpag=L76+0.2084.: May0.251dO.018AL. (2). aud the orbital inclination 157.6 ο," The masses of the two binary components are $M_{\rm PSR}=1.76\pm0.20\,M_\odot$; $M_{\rm WD}=0.254\pm0.018\,M_\odot$ \citep{2008ApJ...679..675V}, and the orbital inclination $i=137.6\pm0.2$ ."
+ iis the brightest ΕΕof the MSPs— in the UV. and X-rays., is the brightest of the MSPs in the UV and X-rays.
+ Detailed spectral studies have suggested that the bulk of the X-ray cussion comes from hot polar regions. where the tempcrature is nou-iuifor. decreasing with distauce away from the pole (2:: 73).," Detailed spectral studies have suggested that the bulk of the X-ray emission comes from hot polar regions, where the temperature is non-uniform, decreasing with distance away from the pole \citealt{1998A&A...329..583Z}; \citealt{2002ApJ...569..894Z}) )."
+ For the WD companion. ? give black-body fit of T~ 000K to eround-based —broad-band plotometry.," For the WD companion, \citet{1993A&A...276..382D} give black-body fit of $T\sim$ K to ground-based broad-band photometry."
+ With such a cool WD. its emission no longer dominates in the UV if the NS is hot enough.," With such a cool WD, its emission no longer dominates in the UV if the NS is hot enough."
+ This offers a unique opportunity to observe thermal cussion from the bulk of the neutron star surface in the UV/optical., This offers a unique opportunity to observe thermal emission from the bulk of the neutron star surface in the UV/optical.
+ 7 (hereafter NOL) observed for the first tine in the UV with the Space Telescope —naging Spectrometer (STIS) aboard OST., \cite{2004ApJ...602..327K} (hereafter K04) observed for the first time in the UV with the Space Telescope Imaging Spectrometer (STIS) aboard HST.
+ WOl found a spectrum cousisteut with the Ravleigh-Jeaus tail of retinal ciuission although the measured spectral shape wd a large uncertainty., K04 found a spectrum consistent with the Rayleigh-Jeans tail of thermal emission although the measured spectral shape had a large uncertainty.
+ Assmuineg the maxiunun area rat this could be enmütted from is the surface of the jeutron star. they derived a surface temperature of T.~10 KK. Tere we describe au extensive observational campaign Hu various spectral bands.," Assuming the maximum area that this could be emitted from is the surface of the neutron star, they derived a surface temperature of $T\sim10^5$ K. Here we describe an extensive observational campaign in various spectral bands."
+ IN Section 2 we descibe the data and its analysis. one spectral window at a time.," IN Section 2 we descibe the data and its analysis, one spectral window at a time."
+ The results aud spectral fits from cach dataset are presented in Section 3. with a complete imultiowaveleugth analysis in Section l|.," The results and spectral fits from each dataset are presented in Section 3, with a complete multi-wavelength analysis in Section 4."
+ The results are discussed iu Section 5. followed by a bref παν aud outlook in Section 6.," The results are discussed in Section 5, followed by a brief summary and outlook in Section 6."
+ We utilize a suite of observations of ffrom erounud-. aud. space-based observatorics., We utilize a suite of observations of from ground- and space-based observatories.
+ Here we describe in detail cach part of our campaign. aud the quality of the resulting data.," Here we describe in detail each part of our campaign, and the quality of the resulting data."
+ We obtain spectral data in the mid-IR to 5-avs. with a combination of spectroscopic and dmacine photometry observations (sec Table 1).," We obtain spectral data in the mid-IR to $\gamma$ -rays, with a combination of spectroscopic and imaging photometry observations (see Table 1)."
+ We made use of ACS on UST. FORSL on VLT/UTI1 of ESO/Paranal. aud PANIC on Magellan I. Furthermore. we have re-analwzed archivalChandra.Newton.Spitzer and data.," We made use of ACS on HST, FORS1 on VLT/UT1 of ESO/Paranal, and PANIC on Magellan I. Furthermore, we have re-analyzed archival, and data."
+ The oulv measurements already in the Literature which we have included (Qvithout reanalysis) are the ground-based plotometric points of ?. and the STIS FUV fiuxes of NOL.," The only measurements already in the literature which we have included (without reanalysis) are the ground-based photometric points of \citet{1993A&A...276..382D}, , and the STIS FUV fluxes of K04."
+ There are archival IRAC and MIPSSpitzer observations of1£715., There are archival IRAC and MIPS observations of.
+. Theo observations— iu he fourSpitzer TRAC channels. 3.6jou. 15jnu. Όδι. S.0gan. (2) were carried out in November 2001. the AIIPS jan observations (2?) were done in September 2005.," The observations in the four IRAC channels, $3.6\,\mu$ m, $4.5\,\mu$ m, $5.8\mu$ m, $8.0\,\mu$ m, \citep{Fazio2004} were carried out in November 2004, the MIPS $\mu$ m observations \citep{Rieke2004} were done in September 2005."
+ The pointing accuracy ofSpitzer is reported to 0 better than 05. according to the TRAC instrumentraundbook?.," The pointing accuracy of is reported to be better than $0\farcs5$, according to the IRAC instrument."
+. Astrometric accuracy of better than 0/2 can be reached. especially if cross-correlating witli 2ATASS point sources is successful.," Astrometric accuracy of better than $0\farcs2$ can be reached, especially if cross-correlating with 2MASS point sources is successful."
+ This is the case for he observations of1715.. where around 20 2ATASS point sources with IRAC couuterparts are in the field.," This is the case for the observations of, where around 20 2MASS point sources with IRAC counterparts are in the field."
+ We apply theSpitzer Software. AIOPEN (version anel its point source extraction package. APEN. or the data reduction. which we describe in detail.," We apply the Software, MOPEX (version and its point source extraction package, APEX, for the data reduction, which we describe in detail."
+ The IRAC observations are stronely affected by extended enission. which turus out to be an extended artifact iu the detector. probably lateut from previous observations.," The IRAC observations are strongly affected by extended emission, which turns out to be an extended artifact in the detector, probably latent from previous observations."
+ We select only those Corrected Basic Calibrated Data files (CBCD: they already include some pipeline processing) from the dithered observations where the target position is well separated from the artifact., We select only those Corrected Basic Calibrated Data files (CBCD; they already include some pipeline processing) from the dithered observations where the target position is well separated from the artifact.
+ This reduces the depth of the observations: iu the case of the fourth channel (844242). for example. oulv half of the available coverage can be used.," This reduces the depth of the observations; in the case of the fourth channel $\mu$ m), for example, only half of the available coverage can be used."
+ For this channel we removed one more CBCD file because of a cosmic rav hit right at the target position., For this channel we removed one more CBCD file because of a cosmic ray hit right at the target position.
+ The resulting mosaics at the target position are shown in Figure 1: for the first channel the position at the epoch of the IRAC observation of lis iudicated., The resulting mosaics at the target position are shown in Figure \ref{IRAC}; for the first channel the position at the epoch of the IRAC observation of is indicated.
+Spitzer IRAC pliotoimetric accuracy can reach 2% for sources. provided the proper aperture correction. the correction of the photometry array location effect. and the color correction are taken iuto account (?7)..," IRAC photometric accuracy can reach $2\%$ for sources, provided the proper aperture correction, the correction of the photometry array location effect, and the color correction are taken into account \citep{Reach2005}."
+ As recommended by theSpitzer Science. Center. we use aperture photometry instead of profile fitting photometry for the measurements. using a source aperture radius of 2 native ΠΑΟ pixels (1722. pixel size) aud a sky annulus from2 to 6 native IRAC pixels to avoid bright sources in the neighborhood. (," As recommended by the Science Center, we use aperture photometry instead of profile fitting photometry for the measurements, using a source aperture radius of 2 native IRAC pixels 2 pixel size) and a sky annulus from2 to 6 native IRAC pixels to avoid bright sources in the neighborhood. ("
+This results im a slight over-estination of the skv due to flux from the source,This results in a slight over-estimation of the sky due to flux from the source
+outer boundary condition. we integrate our structure equations with equation (1)) to find the thermal state for an (M) and accreted layer mass.,"outer boundary condition, we integrate our structure equations with equation \ref{eq:heateq}) ) to find the thermal state for an $\timav$ and accreted layer mass."
+ See Figure 1. for examples of the resulting T-P relations., See Figure \ref{fig:T-P} for examples of the resulting $T$ $P$ relations.
+ We then evaluate the luminosity across the chosen location (the right edge of the plot in Figure 15) for different acereted layer masses. up to the unstable ignition which is found by comparing the 7 and p at the base of the accreted layer with analytic ignition curves (Fujimoto 1982).," We then evaluate the luminosity across the chosen location (the right edge of the plot in Figure \ref{fig:T-P}) ) for different accreted layer masses, up to the unstable ignition which is found by comparing the $T$ and $\rho$ at the base of the accreted layer with analytic ignition curves (Fujimoto 1982)."
+" We vary 7. to find an equilibrium model. where the ""core luminosity"" (L,,,,.) averages to zero over the classical nova cycle as shown in Figure 2.."," We vary $T_c$ to find an equilibrium model, where the “core luminosity” $L_{\rm core}$ ) averages to zero over the classical nova cycle as shown in Figure \ref{fig:loop}."
+ The quiescent 7; for the same cycle is also shown in Figure 2.., The quiescent $T_{\rm eff}$ for the same cycle is also shown in Figure \ref{fig:Teff}.
+ At the nova outburst we assume that the accreted shell is expelled. and that. due to the rapidity of this event. 1t does not appreciably heat the WD.," At the nova outburst we assume that the accreted shell is expelled, and that, due to the rapidity of this event, it does not appreciably heat the WD."
+ The resulting equilibrium core temperatures for (Mj=107Myr! νο9.7.5 and 8.5 for M20.4.0.6 and 1.0M...," The resulting equilibrium core temperatures for $\timav=10^{-10}M_\odot\ {\rm yr}^{-1}$ are $T_c/10^6{\rm
+K}=9, 7.5$ and $8.5$ for $M=0.4, 0.6$ and $1.0M_\odot$."
+ The ..0.4M star ts hotter than the 0.6M star because it has a larger maximum.. accumulated mass that ..leads to a longer period of core heating., The $0.4M_\odot$ star is hotter than the $0.6M_\odot$ star because it has a larger maximum accumulated mass that leads to a longer period of core heating.
+ For a 0.6M.. WD. the core temperatures are 7../1069K24.5.3.12.2 and 18.0 for (M3/M..yr!2107132«1071.42©1071? and 1077.," For a $0.6M_\odot$ WD, the core temperatures are $T_c/10^6{\rm K}=4,5.3,12.2$ and 18.0 for $\timav/M_\odot \ {\rm yr^{-1}}=10^{-11},3.2\times 10^{-11},
+4.2\times 10^{-10}$ and $10^{-9}$."
+ The 7; during the classical nova cycle varies over a relatively narrow range that allows us to compare to observations., The $T_{\rm eff}$ during the classical nova cycle varies over a relatively narrow range that allows us to compare to observations.
+ For field CVs. the large set of STIS observations by Szkody et ((2001) and previous observations (Urban et 22000) provide spectra of quiescent WDs in DN.," For field CVs, the large set of STIS observations by Szkody et (2001) and previous observations (Urban et 2000) provide spectra of quiescent WDs in DN."
+ These measurements are made during deep quiescence when the accretion luminosity ts negligible and are long enough after the outbursts that other emission mechanisms (e.g. Pringle's (1988) suggestion of radiative illumination. of the WD) have faded., These measurements are made during deep quiescence when the accretion luminosity is negligible and are long enough after the outbursts that other emission mechanisms (e.g. Pringle's (1988) suggestion of radiative illumination of the WD) have faded.
+ The observed ζωης thus measure the heat directly from the WD interior., The observed $T_{\rm eff}$ 's thus measure the heat directly from the WD interior.
+ This comparison to observations indicates that below the period gap. (M)z107?M..yr! and the WD masses are in the range 0.6-1.0M4...," This comparison to observations indicates that below the period gap, $\timav\approx
+10^{-10}M_\odot\ \rm yr^{-1}$ and the WD masses are in the range $0.6$ $1.0M_\odot$."
+ This agrees with the expectation from Kolb Baraffe (1999). who find (M;z5«107!M..yr! at an orbital period of 2 hours presuming angular momentum losses from gravitational waves alone.," This agrees with the expectation from Kolb Baraffe (1999), who find $\timav\approx 5\times 10^{-11} M_\odot \
+{\rm yr}^{-1}$ at an orbital period of 2 hours presuming angular momentum losses from gravitational waves alone."
+ Above the period gap. the 7; is higher. and we estimate (M)z107M.γι.," Above the period gap, the $T_{\rm eff}$ is higher, and we estimate $\timav\approx
+10^{-9}M_\odot\ \rm yr^{-1}$."
+ Fora graphical comparison. see Townsley Bildsten (2001).," For a graphical comparison, see Townsley Bildsten (2001)."
+ This general agreement with data from field WDs in which the internal luminosity ts directly visible gives us confidence that our calculations can be applied to other quiescent DN systems., This general agreement with data from field WDs in which the internal luminosity is directly visible gives us confidence that our calculations can be applied to other quiescent DN systems.
+ We predict that a 0.647.. WD above the gap has 7.=1.8«107 K and. if in equilibrium below the gap. 7.=7.5«10° K. However. if the WD does not have time to cool as it traverses the gap. it will be hotter than our calculation implies.," We predict that a $0.6M_\odot$ WD above the gap has $T_c=1.8\times 10^7$ K and, if in equilibrium below the gap, $T_c=7.5\times 10^6$ K. However, if the WD does not have time to cool as it traverses the gap, it will be hotter than our calculation implies."
+ We estimate this cooling time from the current WD cooling law (e.g. Chabrier et 22000). Lia©107(0/18«10K). along with the heat capacity of the core.LL M3Kkp/jrgny. giving Arz0.5 Gyr.," We estimate this cooling time from the current WD cooling law (e.g. Chabrier et 2000), $L_{\rm cool} \approx 10^{-2}L_\odot (T_c/1.8\times
+10^7 {\rm K})^{2.5}$, along with the heat capacity of the core, $M\,3k_B/\mu_im_p$, giving $\Delta t\approx 0.5$ Gyr."
+ Since this i$ comparable to the estimated time spent in the gap (Howell et 22001). our equilibrium assumption below the gap Is likely safe.," Since this is comparable to the estimated time spent in the gap (Howell et 2001), our equilibrium assumption below the gap is likely safe."
+ However. note that about 0.2 Gyrs after accretion halts. the WD will enter the ZZ Ceti instability strip!," However, note that about 0.2 Gyrs after accretion halts, the WD will enter the ZZ Ceti instability strip!"
+ Due to the high frequency of stellar interactions in Globular Clusters (GCs). an abundant population of CVs is expected to be found there. especially at low (M;.," Due to the high frequency of stellar interactions in Globular Clusters (GCs), an abundant population of CVs is expected to be found there, especially at low $\timav$."
+ CVs in GCs are commonly searched for via the presence of hydrogen emission lines or X-ray emission (as recent Chandra observations have found: Grindlay et 220012. Grindlay et 22001b). and this method is fruitful.," CVs in GCs are commonly searched for via the presence of hydrogen emission lines or X-ray emission (as recent Chandra observations have found; Grindlay et 2001a, Grindlay et 2001b), and this method is fruitful."
+ We show that these systems (as well as CVs crossing the period gap or those “hibernating” post-novae. Shara et 11986) can also be identified by their position in a color-magnitude diagram (CMD).," We show that these systems (as well as CVs crossing the period gap or those “hibernating” post-novae, Shara et 1986) can also be identified by their position in a color-magnitude diagram (CMD)."
+ By using our theory of the thermal state of the WD. it is possible to predict the broadband colors of quiescent CVs.," By using our theory of the thermal state of the WD, it is possible to predict the broadband colors of quiescent CVs."
+ An excellent example is NGC 6397 (King et 11998: Taylor et 22001)., An excellent example is NGC 6397 (King et 1998; Taylor et 2001).
+ Figure 3 shows a CMD of NGC 6397 with our initial results., Figure \ref{fig:gc} shows a CMD of NGC 6397 with our initial results.
+ The data points are objects which meet the proper- criteria for cluster membership and which are below the MS., The data points are objects which meet the proper-motion criteria for cluster membership and which are below the MS.
+ The lines were produced by superposing a WD with the, The lines were produced by superposing a WD with the
+to parametrize.,to parametrize.
+ Since the ejecta opacity can evolve in many different ways. several possibilities were tested.," Since the ejecta opacity can evolve in many different ways, several possibilities were tested."
+ In Fig. 7..," In Fig. \ref{Opac},"
+ we show the final model selected for the ejecta opacity evolution., we show the final model selected for the ejecta opacity evolution.
+ This model does not have any theoretical justification. but the true opacity should not be very different from the model proposed here.," This model does not have any theoretical justification, but the true opacity should not be very different from the model proposed here."
+ This model optimally fits all the radio data., This model optimally fits all the radio data.
+ In principle. one would expect that the evolution of the ejecta opacity must not be the same at all frequencies higher than GGHz.," In principle, one would expect that the evolution of the ejecta opacity must not be the same at all frequencies higher than GHz."
+ It seems more plausible that the opacity at higher frequencies should begin to decrease before those at lower frequencies., It seems more plausible that the opacity at higher frequencies should begin to decrease before those at lower frequencies.
+ However. the data is not good enough to allow for such a careful modelling of the ejecta opacity.," However, the data is not good enough to allow for such a careful modelling of the ejecta opacity."
+ The opacity at frequencies higher than GGHz linearly decreases from (on day 1500) to (on day 2500). but remains constant at this frequency.," The opacity at frequencies higher than GHz linearly decreases from (on day 1500) to (on day 2500), but remains constant at this frequency."
+ This decrease in the ejecta opacity can explain the wavelength effects reported in the expansion curve of the supernova. but can also explain a slight increase in the 11993J flux densities observed after day ~1500 in the GGHz data (and. to a lower degree. also in the GGHz data) reported by Weiler et al. (2007)).," This decrease in the ejecta opacity can explain the wavelength effects reported in the expansion curve of the supernova, but can also explain a slight increase in the 1993J flux densities observed after day $\sim$ 1500 in the GHz data (and, to a lower degree, also in the GHz data) reported by Weiler et al. \cite{Weiler2007}) )."
+ In Fig. 9..," In Fig. \ref{WeilerRes},"
+ we show the flux-density residuals corresponding to the model used by Weiler et al. (2007)), we show the flux-density residuals corresponding to the model used by Weiler et al. \cite{Weiler2007}) )
+ around day 1500 after explosion., around day 1500 after explosion.
+ In Fig. 10..," In Fig. \ref{RAMSESRes},"
+ we show the residuals of RAMSES for the same time range., we show the residuals of RAMSES for the same time range.
+ We note that the RAMSES residuals are typically half of those of Weiler et al. (2007)).," We note that the RAMSES residuals are typically half of those of Weiler et al. \cite{Weiler2007}) ),"
+ and even ~5 times smaller for some data points., and even $\sim$ 5 times smaller for some data points.
+ The effect of opacity evolution can also be seen in the measured spectral indices., The effect of opacity evolution can also be seen in the measured spectral indices.
+ In Fig. 8...," In Fig. \ref{SPECIND},"
+ we show the spectral indices reported by Weiler et al. (2007)), we show the spectral indices reported by Weiler et al. \cite{Weiler2007}) )
+ between 1.4 and GGHz. comparing them with the model proposed by these authors and our model obtained with RAMSES.," between 1.4 and GHz, comparing them with the model proposed by these authors and our model obtained with RAMSES."
+ In the time window between 500 and 5000 days after the explosion. the systematic offsets between data and the model proposed by Weiler et al.," In the time window between 500 and 5000 days after the explosion, the systematic offsets between data and the model proposed by Weiler et al."
+ can clearly be seen., can clearly be seen.
+ Instead. the RAMSES model predicts remarkably well the evolution of the spectral index for all epochs. including the fattening beginning at an age ~ 1000 days (Pérrez-Torres et al. 2002a:," Instead, the RAMSES model predicts remarkably well the evolution of the spectral index for all epochs, including the flattening beginning at an age $\sim$ 1000 days (Pérrez-Torres et al. \cite{PerezTorres2002};"
+: Bartel et al. 2002))., Bartel et al. \cite{Bartel2002}) ).
+ In Marcaide et al. (2009a)), In Marcaide et al. \cite{Marcaide2009}) )
+ and Paper L we reported a fitted ejecta opacity of for all epochs and frequencies.," and Paper I, we reported a fitted ejecta opacity of for all epochs and frequencies."
+ This result might seem to be in conflict with the model proposed here for the evolution of the ejecta opacity., This result might seem to be in conflict with the model proposed here for the evolution of the ejecta opacity.
+ However. we also noted then that the ejecta opacities reported were too noisy for extracting any robust information about the possible evolution and/or spectral dependence of the ejecta opacity.," However, we also noted then that the ejecta opacities reported were too noisy for extracting any robust information about the possible evolution and/or spectral dependence of the ejecta opacity."
+ Therefore. the value of reported there for the opacity should be taken as an approximate. average. value.," Therefore, the value of reported there for the opacity should be taken as an approximate, average, value."
+ Inside the shell it ts also difficult to parametrize a radial drop in the magnetic field., Inside the shell it is also difficult to parametrize a radial drop in the magnetic field.
+ After extensive testing. we chose the following model to characterize the drop of B. as a function of distance. D. from the contact discontinuity where the parameters a. >. and ο are chosen so that. for a fractional shell width of 0.3. Bo is the mean magnetic field intensity of the shell and the intensity at the forward shock ts half the intensity at the contact discontinuity (see Fig. 7)).," After extensive testing, we chose the following model to characterize the drop of $B$ as a function of distance, $D$, from the contact discontinuity where the parameters $a$, $b$, and $c$ are chosen so that, for a fractional shell width of 0.3, $B_0$ is the mean magnetic field intensity of the shell and the intensity at the forward shock is half the intensity at the contact discontinuity (see Fig. \ref{Opac}) )."
+ Mareaide et al. (20093)), Marcaide et al. \cite{Marcaide2009}) )
+ pointed out that if the flux density per unit beam decreases. the shell size estimate will be biased towards a smaller value. provided the magnetic field drops radially in the shell (see their Sect.," pointed out that if the flux density per unit beam decreases, the shell size estimate will be biased towards a smaller value, provided the magnetic field drops radially in the shell (see their Sect."
+ 7.1.2)., 7.1.2).
+ Thus. the exponential-like decrease in flux densities after day 3100. combined with a magnetic field structure similar to that given in Eq. I..," Thus, the exponential-like decrease in flux densities after day 3100, combined with a magnetic field structure similar to that given in Eq. \ref{EqB},"
+ should translate into progressively biased estimates of the shell size. and therefore an increase in the (observed) deceleration of the expansion curve.," should translate into progressively biased estimates of the shell size, and therefore an increase in the (observed) deceleration of the expansion curve."
+ If we include radial drops in the magnetic field steeper than that corresponding to Eq., If we include radial drops in the magnetic field steeper than that corresponding to Eq.
+ | (such as a linear or a concave-like decay). we obtain poorer fits to the expansion curve.," \ref{EqB} (such as a linear or a concave-like decay), we obtain poorer fits to the expansion curve."
+" Therefore. we conclude that the radial ""Srop in the magnetic field inside the shell must be smooth."," Therefore, we conclude that the radial drop in the magnetic field inside the shell must be smooth."
+ In Fig. I 1.. ," In Fig. \ref{ResExpan}, ,"
+we compare the expansion-curve residuals obtained with a uniform magnetic field inside the shell (a) with those for the magnetic-field structure given by Eq., we compare the expansion-curve residuals obtained with a uniform magnetic field inside the shell (a) with those for the magnetic-field structure given by Eq.
+ | and shown in Fig. 7.., \ref{EqB} and shown in Fig. \ref{Opac}.
+ It can be seen that the residuals for the latest VLBI epochs are more accurately fitted using the radially-decaying magnetic field., It can be seen that the residuals for the latest VLBI epochs are more accurately fitted using the radially-decaying magnetic field.
+ We note however. that the conclusions extracted from this section are based on noisy images (from the latest epochs at which the supernova could be barely imaged).," We note however, that the conclusions extracted from this section are based on noisy images (from the latest epochs at which the supernova could be barely imaged)."
+ Therefore. these conclusions should be considered with caution.," Therefore, these conclusions should be considered with caution."
+ We have developed software to simultaneously model the VLBI expansion curve and the radio light) curves. of supernova 11993J. This software takes into consideration the evolution of the magnetic field energy density and the hydrodynamic evolution of the expanding shock.às well as the relativistic acceleration of CSM electrons as theyinteract with the supernova shock.," We have developed software to simultaneously model the VLBI expansion curve and the radio light curves of supernova 1993J. This software takes into consideration the evolution of the magnetic field energy density and the hydrodynamic evolution of the expanding shock,as well as the relativistic acceleration of CSM electrons as theyinteract with the supernova shock."
+ All these processes have, All these processes have
+"with the data in Figure τας, and the resultant. Davesian posterior distribution. with the input prior distribution. for a is shown in Figure Thb. We see that the posterior distribution is Gaussian. is centered at a significantly different. value than (he prior distribution (1.15 instead of 0.8). and is more tightly. constrained than the prior.","with the data in Figure \ref{fig:pf}a a, and the resultant Bayesian posterior distribution, with the input prior distribution, for $\alpha$ is shown in Figure \ref{fig:pf}b b. We see that the posterior distribution is Gaussian, is centered at a significantly different value than the prior distribution (1.15 instead of 0.8), and is more tightly constrained than the prior."
+ Since we have reproduced (he earlier resulis and the resultant. posterior distributions are well-behaved. we believe that our Bavesian method is correctly. constructed.," Since we have reproduced the earlier results and the resultant posterior distributions are well-behaved, we believe that our Bayesian method is correctly constructed."
+ We will proceed (o use it to understiixd the substellar mass function., We will proceed to use it to understand the substellar mass function.
+ We now extend the above demonstration of the utility of a Davesian approach to include the substellar mass function of the local field., We now extend the above demonstration of the utility of a Bayesian approach to include the substellar mass function of the local field.
+ As described in the Introduction. we use number counts of MT to L8 dwarls taken from Cruzetal.(2003.D)2005).. and space densities ol T5-TS dwarls taken from Bureasser(2002) to study the field substellar mass function.," As described in the Introduction, we use number counts of M7 to L8 dwarfs taken from \citet{kc03,kc05}, and space densities of T5-T8 dwarfs taken from \citet{burgp} to study the field substellar mass function."
+ We combine those data into a joint Z-band luminosity function in order to compare (hem to our models., We combine those data into a joint $J$ -band luminosity function in order to compare them to our models.
+ Cruzetal.(2005). provides a J-band Iuminosity function directly (although known to be incomplete for M cwarls (AL;« 11))., \citet{kc05} provides a $J$ -band luminosity function directly (although known to be incomplete for M7 dwarfs $_J < 11$ )).
+ However. the Bureasser(2002) E dwarl dala provide space density. as a function of spectral type.," However, the \citet{burgp} T dwarf data provide space density as a function of spectral type."
+ We have used the T. dwarf M;- Type relation from Vrbaetal.(2004) to Gransform the spectral (vpe distribution to a J-band luminosity ΠΟΙΟΙ: where Mj is (he absolute /-band magnitude and Sp is a spectral (ype index (I0-T3 = 0-3)., We have used the T dwarf $M_J$ -Spectral Type relation from \citet{vrba} to transform the spectral type distribution to a $J$ -band luminosity function: where $M_J$ is the absolute $J$ -band magnitude and $SpT$ is a spectral type index (T0-T8 = 0-8).
+ Combining those results with the Cruz οἱ ddata gives the empirical ΑςΔΕΟ J-band luminosity function listed in Table 3., Combining those results with the Cruz et data gives the empirical KCAB $J$ -band luminosity function listed in Table 3.
+" The M;-5Spectral Type relation is double-valued [ον spectral types between L5 and Το, reversing ils direction at around L/T (ransiion. with earlv-tvpe T dwarls ancl Iate-tvpe L dwarls having similar M. (Vrbaetal.2004)."," The $_J$ -Spectral Type relation is double-valued for spectral types between $\sim$ L5 and T5, reversing its direction at around L/T transition, with early-type T dwarfs and late-type L dwarfs having similar $_J$ \citep{vrba}."
+". Survey data for this absolute magnitude reeime (14.0<M, 15.5) are incomplete.", Survey data for this absolute magnitude regime $14.0 \le M_J \le 15.5$ ) are incomplete.
+ This reflects both known incompleteness in the Cruzetal.(2005) survey for L5-L8 dwarls and the absence of published density. estimates for TO-T4 cdwarls., This reflects both known incompleteness in the \citet{kc05} survey for L5-L8 dwarfs and the absence of published density estimates for T0-T4 dwarfs.
+" Thus. the space densities between Af;=14 and M,=15.5 listed in Table 3 are lower limits. and (hese data will not be used to constrain the field. substellar mass [unction."," Thus, the space densities between $M_J = 14$ and $M_J = 15.5$ listed in Table 3 are lower limits, and these data will not be used to constrain the field substellar mass function."
+ As outlined in the introduction. three different mass function models will be used to fit to the data: a (wo-segment power law (as used in 822 and A03). à log normal mass function. and a single power law with a low-mass cutoff.," As outlined in the introduction, three different mass function models will be used to fit to the data: a two-segment power law (as used in 2 and A03), a log normal mass function, and a single power law with a low-mass cutoff."
+ The segmented power law, The segmented power law
+"cube, we can subdivide the entire range of bubble sizes into small size bins.","cube, we can subdivide the entire range of bubble sizes into small size bins."
+ The above analysis would then be applicable to each individual size bin and thus to the entire range., The above analysis would then be applicable to each individual size bin and thus to the entire range.
+ Our second toy model considers a slightly more realistic but still quite simple reionization scenario., Our second toy model considers a slightly more realistic but still quite simple reionization scenario.
+" In this toy model, we consider a reionization model in which spherical ionized bubbles of different sizes are placed randomly."," In this toy model, we consider a reionization model in which spherical ionized bubbles of different sizes are placed randomly."
+" If there is no overlap between ionized bubbles then the ionized fraction would be, [4 (8)where n(R) is the number density of bubbles of size R."," If there is no overlap between ionized bubbles then the ionized fraction would be, dR V(R) where $n(R)$ is the number density of bubbles of size $R$."
+ But in practice randomly placed bubbles will overlap with each other and expand further to conserve the emitted photon numbers., But in practice randomly placed bubbles will overlap with each other and expand further to conserve the emitted photon numbers.
+ We neglect the fact of further expansion of bubbles and so the actual ionized fraction which would be less than the above can be calculated using (Furlanettoetal.2004) exp((-Q).," We neglect the fact of further expansion of bubbles and so the actual ionized fraction which would be less than the above can be calculated using \citep{fur04}
+ (-Q)."
+" The spherically averaged 3D power spectrum can be calculated from the two point correlation function using the relation Am drr£((r) Here we would like to mention that the evolution makes the correlation function € anisotropic i.e, a function of both r and µ (Barkana&Loeb2006;SethiHaiman2008)."," The spherically averaged 3D power spectrum can be calculated from the two point correlation function using the relation }(k)=4 dr r^2 (r) Here we would like to mention that the evolution makes the correlation function $\xi$ anisotropic i.e, a function of both $r$ and $\mu$ \citep{barkana06,sethi08}."
+". Since our aim is to study the light cone effect on the spherically averaged power spectrum and qualitatively understand the main features we see in the simulation results, we assume £ to be isotropic."," Since our aim is to study the light cone effect on the spherically averaged power spectrum and qualitatively understand the main features we see in the simulation results, we assume $\xi$ to be isotropic."
+ We expect this approximation not to affect our conclusions., We expect this approximation not to affect our conclusions.
+" The correlation function £(r) can be decomposed into (Zaldarriagaetal.2004) where £;4;, £55 and &5 are the correlation functions of ionization field, density field and the cross-correlation between two fields respectively."," The correlation function $\xi(r)$ can be decomposed into \citep{zal04}
+ (r)= where $\xi_{xx}$, $\xi_{\delta \delta}$ and $\xi_{x \delta}$ are the correlation functions of ionization field, density field and the cross-correlation between two fields respectively."
+" In the density field correlation function £55, two quantities change with redshift: 1) density fluctuations grow with time, 2) the mean density decreases because of the expansion of the Universe."," In the density field correlation function $\xi_{\delta \delta}$, two quantities change with redshift: 1) density fluctuations grow with time, 2) the mean density decreases because of the expansion of the Universe."
+ Thus the evolving £55 in principle would contribute to the light cone effect., Thus the evolving $\xi_{\delta \delta}$ in principle would contribute to the light cone effect.
+ In the linear regime the two quantities together scale as ~(1+z)95., In the linear regime the two quantities together scale as $\sim (1+z)^{-0.5}$.
+ For a distance of 163cMpc jointly change ~3% along the LOS.," For a distance of $163 \, \rm{cMpc}$ they jointly change $\sim 3
+\%$ along the LOS."
+ We use very high redshift theysimulation cubes (before the reionization starts) and find <0.1% enhancement in the 3D power spectrum on almost all scales (Figure 10))., We use very high redshift simulation cubes (before the reionization starts) and find $\lesssim 0.1 \%$ enhancement in the 3D power spectrum on almost all scales (Figure \ref{fig:linearpower}) ).
+ This result agrees with McQuinnet.al(2006) who predicted a constant enhancement in the power spectrum (see their Appendix A)., This result agrees with \citet{mcquinn06} who predicted a constant enhancement in the power spectrum (see their Appendix A).
+ The contribution of the evolving £55 to the total light cone effect on the power spectrum is therefore negligible and hence we ignore this term in the rest of our analysis., The contribution of the evolving $\xi_{\delta \delta}$ to the total light cone effect on the power spectrum is therefore negligible and hence we ignore this term in the rest of our analysis.
+ The evolution of £ is thus mainly dominated by the ἔπα on scales larger or comparable to the size of ionized bubbles., The evolution of $\xi$ is thus mainly dominated by the $\xi_{xx}$ on scales larger or comparable to the size of ionized bubbles.
+ For the rest of our analysis we only consider the term ένα., For the rest of our analysis we only consider the term $\xi_{xx}$.
+ The function which is defined as ένα(η)=(πι)—3:2 should be zero both for Tj=O and z;=1., The function which is defined as $\xi_{xx}(r)=\langle x_1x_2\rangle-\bar{x}_i^2$ should be zero both for $\bar{x}_i=0$ and $\bar{x}_i=1$.
+" It also should satisfy the boundary conditions (fordetailsseeZaldarriagaetal.2004) The correlation function can be calculated for a given bubble Psp(k)distribution as (Furlanettoetal.2004) where V;(R,r) is the volume of the overlap region between two ionized regions centered a distance r apart."," It also should satisfy the boundary conditions \citep[for details see][]{zal04}
+ The correlation function can be calculated for a given bubble distribution as \citep{fur04}
+ where $V_o(R,r)$ is the volume of the overlap region between two ionized regions centered a distance $r$ apart."
+ The function can be written as In the next subsection we present some bubble size distributions measured from simulation., The function can be written as In the next subsection we present some bubble size distributions measured from simulation.
+ We will then model the bubble distribution and use that for the subsequent analysis., We will then model the bubble distribution and use that for the subsequent analysis.
+" We calculate the bubble size distribution from the simulation using the spherical average method (seeFriedrichetal.2011,formoredetailsondifferentbubblesize estimates)..", We calculate the bubble size distribution from the simulation using the spherical average method \citep[see][ for more details on different bubble size estimates]{friedrich11}.
+ Fig., Fig.
+" 11 shows the bubble size distribution RdP/dR® for coeval cubes at three redshifts corresponding to the back, middle and front side of a light cone cube centered around redshift 9.09."," \ref{fig:bub-dis-sim} shows the bubble size distribution $R \, dP/dR$ for coeval cubes at three redshifts corresponding to the back, middle and front side of a light cone cube centered around redshift 9.09."
+ It also shows the bubble size distribution in the light cone cube centered around 9.09., It also shows the bubble size distribution in the light cone cube centered around 9.09.
+ The coeval distributions at the three redshifts differ considerably., The coeval distributions at the three redshifts differ considerably.
+" For example the radii at which the bubble distribution peaks are 10, 20, and 90 cMpc."," For example the radii at which the bubble distribution peaks are 10, 20, and 90 cMpc."
+ Interestingly the bubble size, Interestingly the bubble size
+smoothing matrix takes the form of a band which rotates from the vertical to the horizontal as the smoothing scale 3 is increased.,smoothing matrix takes the form of a band which rotates from the vertical to the horizontal as the smoothing scale $\beta$ is increased.
+ Smoothing thus corresponds to à projection of the input modes into output modes of smaller order., Smoothing thus corresponds to a projection of the input modes into output modes of smaller order.
+ The high-order modes indeed have oscillations on small scales and are thus gradually lost when we increase the smoothing scale 3., The high-order modes indeed have oscillations on small scales and are thus gradually lost when we increase the smoothing scale $\beta$.
+" Let us now consider the convolution of two 2-dimensional functions. such as We again first decompose cach function into shapelet coellicients f,2 anal. d=iή). and fy—ish with shapelet scales à. 2 and 7 respectively. and where n=(ni.no) as before."," Let us now consider the convolution of two 2-dimensional functions, such as We again first decompose each function into shapelet coefficients $f_{\mathbf n} \equiv \langle {\mathbf n};\alpha | f \rangle$ , $g_{\mathbf n} \equiv \langle {\mathbf n};\beta | g \rangle$, and $h_{\mathbf n} \equiv \langle {\mathbf n};\gamma | h \rangle$ with shapelet scales $\alpha$, $\beta$ and $\gamma$ respectively, and where ${\mathbf n}=(n_{1},n_{2})$ as before."
+ Because convolution is bilinear we can again relate the convolved to the unconvolved. coefficients by where Chana(s.0.:2) is the 2-dimensional convolution tensor.," Because convolution is bilinear we can again relate the convolved to the unconvolved coefficients by where $C_{{\mathbf n},{\mathbf m},{\mathbf l}}(\gamma,\alpha,\beta)$ is the 2-dimensional convolution tensor."
+ From the separability of the 2-cdimensional basis functions (see Eq. 18)).," From the separability of the 2-dimensional basis functions (see Eq. \ref{eq:phi_n1n2}) ),"
+ it is easy to show that this tensor is cqual to where the tensors appearing on the right-hand. side are. the. I-dimensional convolution tensor defined in Equation (46))., it is easy to show that this tensor is equal to where the tensors appearing on the right-hand side are the 1-dimensional convolution tensor defined in Equation \ref{eq:c_nml}) ).
+ We can also consider the special case of smoothing with a 2-dimensional gaussian., We can also consider the special case of smoothing with a 2-dimensional gaussian.
+ In this case. οκ)= . which is normalised so that {Parg(x)—1.," In this case, $g({\mathbf
+x})=(2\pi\beta^2)^{-1} e^{-\frac{x^2}{2\beta^2}}$ , which is normalised so that $\int d^{2}x~g({\mathbf x})=1$."
+ The smoothed cocllicients are then given by where Cuu(5.0.3) is the 2-cimensional smoothing matrix.," The smoothed coefficients are then given by where $G_{{\mathbf n},{\mathbf m}}(\gamma,\alpha,\beta)$ is the 2-dimensional smoothing matrix."
+ Ht is again easy to show that itis equal to in terms ofthe 1-dimensional smoothing matrix celinec in, It is again easy to show that itis equal to in terms ofthe 1-dimensional smoothing matrix defined in
+Plot. the pixel data were re-examined and the offending variable star was identilied as being about 11 arcseconds away [rom and about 4.8,"Plot, the pixel data were re-examined and the offending variable star was identified as being about 11 arcseconds away from and about 4.8"
+A precise and accurate meastuement of the Hubble constant at the ~2% levels is important for modern precision Cosmology (forexample.seeTegmarkal. 2010).,"A precise and accurate measurement of the Hubble constant at the $\sim2\%$ levels is important for modern precision cosmology \citep[for example, see][]{teg04,hu05,mac06,jac07,oll07,gre09,kom10,lar10,mor10}."
+. This will be feasible after the launch of the(JIWST). which will be capable of routinely observing Cepheids in distaut galaxies (Riessetal.2009:Freecinau&Madore2010) at near- aud iuid-infrared wavelengths.," This will be feasible after the launch of the, which will be capable of routinely observing Cepheids in distant galaxies \citep{rie09,fre10} at near- and mid-infrared wavelengths."
+ Iu. preparation for future distance scale work. it is necessary to derive the Cephlieid Period-Luimiuosity relatious (lierealter P-L relations. also known as the Leavitt Law) at mid-infrared (7340 to ~ Spin) wavelengths.," In preparation for future distance scale work, it is necessary to derive the Cepheid Period-Luminosity relations (hereafter P-L relations, also known as the Leavitt Law) at mid-infrared $\sim3\mu{\mathrm m}$ to $\sim8\mu{\mathrm m}$ ) wavelengths."
+ Some advantages aud potential problems of the mid-infrared P-L relatious were sumauarized. for example. in Ngeowreference therein)..," Some advantages and potential problems of the mid-infrared P-L relations were summarized, for example, in \citet[][and reference therein]{nge09a}."
+ For example. the contribution of extinction to the error budget of," For example, the contribution of extinction to the error budget of"
+"From our simulations we conclude that we should. be able to recover the ""real, abundance values of Vega. i£ this star is indeed. observed with an inclination angle of about τν as suggested by Takedactal.(2008).","From our simulations we conclude that we should be able to recover the `real' abundance values of Vega, if this star is indeed observed with an inclination angle of about $^{\circ}$, as suggested by \citet{takeda2008}."
+. However. if Vega has an inclination angle of about 4. as used by Yoonctal. (2008).. the influence of rotation should. indeed. be visible in the computed abundances.," However, if Vega has an inclination angle of about $^{\circ}$, as used by \citet{yoon2008}, the influence of rotation should indeed be visible in the computed abundances."
+ Note that the residual shape of IL? of Vega is more similar to our simulations with the lowest inclination angle., Note that the residual shape of $\beta$ of Vega is more similar to our simulations with the lowest inclination angle.
+ The procedure for stellar parameter determination and abundance analysis. described. in Sections 2 and 3. and tested on syntjetic spectra ancl Vega. will now be applied to observations ο “A and E stars that we have at our disposal.," The procedure for stellar parameter determination and abundance analysis, described in Sections 2 and 3, and tested on synthetic spectra and Vega, will now be applied to observations of A and F stars that we have at our disposal."
+ This sample of stars contains both single stars and multiple systems. and hese will allow us to study the relationship between pulsaions. multiplicity and. chemical composition.," This sample of stars contains both single stars and multiple systems, and these will allow us to study the relationship between pulsations, multiplicity and chemical composition."
+ We have also οxXained spectra of suspected + Dor stars. with the aim of bot1 confirming their classification and applying this. abunclzuὉ procedure.," We have also obtained spectra of suspected $\gamma$ Dor stars, with the aim of both confirming their classification and applying this abundance procedure."
+ Results will be. presented. in subsequent puMications., Results will be presented in subsequent publications.
+thrown.,thrown.
+ In this case the metallicity gradient of the Galaxy does not change., In this case the metallicity gradient of the Galaxy does not change.
+ Collins et al. (, Collins et al. (
+2002) suggested that the ballistic model problems could be solved by considering the presence of drag between disk and halo. such as through magnetic tension or viscous interactions between clouds.,"2002) suggested that the ballistic model problems could be solved by considering the presence of drag between disk and halo, such as through magnetic tension or viscous interactions between clouds."
+ Barnabé et al. (, Barnabè et al. (
+2006) presented fluid stationary models able to reproduce the observed negative vertical gradient of the rotation velocity of the extra-planar gas.,2006) presented fluid stationary models able to reproduce the observed negative vertical gradient of the rotation velocity of the extra-planar gas.
+ In these models the gas is described as a stationary rotational fluid., In these models the gas is described as a stationary rotational fluid.
+ In the case of a single cloud. the solutions of this model predict the drag invoked by Collins et al. (," In the case of a single cloud, the solutions of this model predict the drag invoked by Collins et al. ("
+2002). as ram pressure between cloud and fluid halo.,"2002), as ram pressure between cloud and fluid halo."
+ This suggests that a correct deseription. of the extra-planar gas dynamics may be found in hybrid ballistic-fluid stationary models., This suggests that a correct description of the extra-planar gas dynamics may be found in hybrid ballistic-fluid stationary models.
+ Recently. Fraternali Binney (2008) presented a new formulation for viscous cloud-halo interactions in. the framework of the cold gas accretion onto spiral galaxies in the local universe.," Recently, Fraternali Binney (2008) presented a new formulation for viscous cloud-halo interactions in the framework of the cold gas accretion onto spiral galaxies in the local universe."
+ In this paper we will follow the clouds ejected by SN explosions both in the Galaxy in the purely ballistie model and in the hybrid one., In this paper we will follow the clouds ejected by SN explosions both in the Galaxy in the purely ballistic model and in the hybrid one.
+ The paper is organized as follows: in Sect., The paper is organized as follows: in Sect.
+ ?? we will describe our model. in Sect ??. we will show our results concerning the chemical abundances of the supershells and the cloud orbits in the purely ballistie and im the hybrid model.," \ref{model} we will describe our model, in Sect \ref{results} we will show our results concerning the chemical abundances of the supershells and the cloud orbits in the purely ballistic and in the hybrid model."
+ In Sect., In Sect.
+ ??) we compare our results with observations of the HVC Complex C and the IVC VIAre., \ref{hvc} we compare our results with observations of the HVC Complex C and the IVC VIArc.
+ In Sect., In Sect.
+ 22. we report our conclusions., \ref{conclusions} we report our conclusions.
+ Type II SNe usually occur in OB associations containing several dozens of massive stars., Type II SNe usually occur in OB associations containing several dozens of massive stars.
+ Sequential SN explosions create a superbubble. whereas the swept up gas Is concentrated in à thin cold shell called supershell.," Sequential SN explosions create a superbubble, whereas the swept up gas is concentrated in a thin cold shell called supershell."
+ The superbubble expansion in a stratified medium does not follow a spherical evolution and the Kompaneets (1960) approximation well describes the way in which the superbubble grows up in the meridional Galaxy plane., The superbubble expansion in a stratified medium does not follow a spherical evolution and the Kompaneets (1960) approximation well describes the way in which the superbubble grows up in the meridional Galaxy plane.
+ After some interval of time. a region of the supershell can fragment due to the occurrence of the Rayleigh-Taylor instabilities (RTIs). therefore clouds of gas can form.," After some interval of time, a region of the supershell can fragment due to the occurrence of the Rayleigh-Taylor instabilities (RTIs), therefore clouds of gas can form."
+ Once left the stellar disk. the orbit of each cloud can be followed either ballistically or with a hybrid model considering viscous interaction between the cloud and the extra-planar gas.," Once left the stellar disk, the orbit of each cloud can be followed either ballistically or with a hybrid model considering viscous interaction between the cloud and the extra-planar gas."
+ The physical parameters of the superbubble can be fairly well described by the Kompaneets (1960) model., The physical parameters of the superbubble can be fairly well described by the Kompaneets (1960) model.
+ The Kompaneets approximation assumes the following: uniform. pressure within the superbubble. superbubble expansion in a direction normal to the local surface and expansion speed implied by a strong shock (1.e.. the internal pressure dominates the external pressure).," The Kompaneets approximation assumes the following: uniform pressure within the superbubble, superbubble expansion in a direction normal to the local surface and expansion speed implied by a strong shock (i.e., the internal pressure dominates the external pressure)."
+ Kompaneets found an analytic expression for the shape of the bubble during its expansion in an exponential atmosphere with density: where po and H are the disk density and height scale. respectively.," Kompaneets found an analytic expression for the shape of the bubble during its expansion in an exponential atmosphere with density: where $\rho_{0}$ and $H$ are the disk density and height scale, respectively."
+ We considerthe shock expansion in cylindrical coordinates ἐς.R).," We considerthe shock expansion in cylindrical coordinates $(z,R)$."
+ In the work of Kompaneets it was assumed that for z<0 the ISM was denser thar the gas residing above the disk., In the work of Kompaneets it was assumed that for $z<0$ the ISM was denser than the gas residing above the disk.
+ Therefore. in this approximation. the bubble has an asymmetric evolution with respect to the plane of the galaxy.," Therefore, in this approximation, the bubble has an asymmetric evolution with respect to the plane of the galaxy."
+ Kompaneets found analytic expressions for the top and bottom sites of the remnant., Kompaneets found analytic expressions for the top and bottom sites of the remnant.
+Referring to Fig.,Referring to Fig.
+ | we have: and The quantity y is atransformed variable (with units of length) defined by: where Q and £j; and y are the volume. the thermal energy and the ratio of specific heats for the superbubble. respectively.," \ref{bubble} we have: and The quantity $y$ is atransformed variable (with units of length) defined by: where $\Omega$ and $E_{th}$ and $\gamma$ are the volume, the thermal energy and the ratio of specific heats for the superbubble, respectively."
+ The thermal energy ts calculated from the differential equation: where Lo is the wind luminosity and the pressure is given by the equation of state: In our work for the ISM z-profile we assume an exponential law also for z«0., The thermal energy is calculated from the differential equation: where $L_0$ is the wind luminosity and the pressure is given by the equation of state: In our work for the ISM z-profile we assume an exponential law also for $z<0$.
+ Hence the superbubble evolves symmetrically with respect to the galactic plane resulting in a peanut-like shape. as suggested by Mac Low et al. (," Hence the superbubble evolves symmetrically with respect to the galactic plane resulting in a peanut-like shape, as suggested by Mac Low et al. ("
+1988) and by Tenorio-Tagle et al. (,1988) and by Tenorio-Tagle et al. (
+1999).,1999).
+ The volume of the superbubble can be estimated by: where a=(=)+[zu] ts the semi-major axis., The volume of the superbubble can be estimated by: where $a=(z_{L}+|z_{H}|)/2$ is the semi-major axis.
+ A numerical integration of eqs. (4)), A numerical integration of eqs. \ref{y}) )
+ and (5)). with the help of eqs. (6))," and \ref{E}) ), with the help of eqs. \ref{P}) )"
+" and (7)). yields y(t) and Ej, which implicitly give the time evolution of the superbubble."," and \ref{Omega}) ), yields $y(t)$ and $E_{th}$ which implicitly give the time evolution of the superbubble."
+ The solution depends on only three parameters: the scale height H. the density near the source po. and the wind luminosity Zo.," The solution depends on only three parameters: the scale height $H$, the density near the source $\rho_{0}$, and the wind luminosity $L_0$."
+ Hence the problem ts solved most naturally in a system of units which are determined by these three parameters., Hence the problem is solved most naturally in a system of units which are determined by these three parameters.
+ The unit of length is then Ho. the unit of mass is poH. and the unit of time is (poH?/Lo)! .," The unit of length is then $H_{0}$ , the unit of mass is $\rho_{0}H^3$, and the unit of time is $(\rho_{0}H^{5}/L_{0})^{1/3}$ ."
+ We obtain the dimensionless solution by integrating the dimensionless form of eqs. (4)), We obtain the dimensionless solution by integrating the dimensionless form of eqs. \ref{y}) )
+ and (5))., and \ref{E}) ).
+ In Figs., In Figs.
+" 2. and 3. we report the dimensionless evolution of v and the velocity dz,/dr.", \ref{fy} and \ref{fdz} we report the dimensionless evolution of $y$ and the velocity $dz_{L}/dt$.
+ The dimensionless analysis predicts that the superbubble begins to accelerate at (see Fig. 3)):, The dimensionless analysis predicts that the superbubble begins to accelerate at (see Fig. \ref{fdz}) ):
+" namely /[Lo)when dz,/dt has a minimum.", namely when $dz_{L}/dt$ has a minimum.
+ We underline that the simulated superbubbles. characterized by different values of po. H. Lo. reach the accelerated phase at different times. as shown in eq. (8)).," We underline that the simulated superbubbles, characterized by different values of $\rho_0$ , $H$ , $L_0$ , reach the accelerated phase at different times, as shown in eq. \ref{tilde}) ),"
+ butthey have the same shape and volume., butthey have the same shape and volume.
+ This is à direct consequence of theself-similar nature of the Kompaneets, This is a direct consequence of theself-similar nature of the Kompaneets
+in the standard manner using two parameters.,in the standard manner using two parameters.
+ The first xwameter. f. models the elleet of large-scale coherent infall velocities.," The first parameter, $f$, models the effect of large-scale coherent infall velocities."
+ Ln linear theory for the erowth of IHuctuations. fois related. to the linear growth factor (e) at scale actor e by. f=din λάμνα.," In linear theory for the growth of fluctuations, $f$ is related to the linear growth factor $D(a)$ at scale factor $a$ by $f = d\ln{D}/d\ln{a}$ ."
+ This functional form is well-approximated by f=Que)? where μις). is the matter density measured. by an observer at. redshift. ο and ςτε1.55 for a standard ACDAL cosmology.," This functional form is well-approximated by $f = \Omega_{\rm m}(z)^\gamma$ where $\Omega_{\rm m}(z)$ is the matter density measured by an observer at redshift $z$, and $\gamma \approx 0.55$ for a standard $\Lambda$ CDM cosmology."
+ Coherent velocities can also be parameterized by Ξ— ffb. introducing a dependence on galaxy tvpe via the bias factor., Coherent velocities can also be parameterized by $\beta = f/b$ introducing a dependence on galaxy type via the bias factor.
+ The second »wameter. a. (in units of f km +). describes. small-scale random. virialized: velocities. for which we assumed an exponential pairwise distribution.," The second parameter, $\sigma_v$ (in units of $h$ km $^{-1}$ ), describes small-scale random virialized velocities, for which we assumed an exponential pairwise distribution."
+ The overall effect. on he galaxy power spectrum ZA is to induce an anisotropic distortion dependent on the angle 8 of the wavevector to the inc-of-sight. parameterized by (/=cos: We fitted these models to the nine independent. 2D. power spectra split into tangential and radial components. in order to determine the redshilt-space distortion. parameters. and test the ACDAL self-consisteney relation μμ. as à function of redshift.," The overall effect on the galaxy power spectrum $P_{\rm gal}$ is to induce an anisotropic distortion dependent on the angle $\theta$ of the wavevector to the line-of-sight, parameterized by $\mu =
+\cos{\theta}$: We fitted these models to the nine independent 2D power spectra split into tangential and radial components, in order to determine the redshift-space distortion parameters and test the $\Lambda$ CDM self-consistency relation $\beta \approx \Omega_{\rm m}(z)^{0.55}/b$ as a function of redshift."
+" In. this initial analysis we fixed the matter density parameter Qu,=0.3 ancl barvon fraction OQ),/Ou=0.166 (similar to the parameters listed in Komatsu et ((2009) that provide the best fit to the 5-vear data set of the Wilkinson Microwave Anisotropy Probe (WALA): we ix O4,=0.3 to match the value assumed for the distance-redshift relation in our power spectrum analysis).", In this initial analysis we fixed the matter density parameter $\Omega_{\rm m} = 0.3$ and baryon fraction $\Omega_{\rm b}/\Omega_{\rm m} = 0.166$ (similar to the parameters listed in Komatsu et (2009) that provide the best fit to the 5-year data set of the Wilkinson Microwave Anisotropy Probe (WMAP); we fix $\Omega_{\rm m} = 0.3$ to match the value assumed for the distance-redshift relation in our power spectrum analysis).
+" We then varied the redshift-space parameters f and a, and the bias actor b. convolving cach model power spectrum. with the survey selection [function for each region ancl minimizing he chi-squared statistic calculated using the full covariance matrices."," We then varied the redshift-space parameters $f$ and $\sigma_v$ and the bias factor $b$, convolving each model power spectrum with the survey selection function for each region and minimizing the chi-squared statistic calculated using the full covariance matrices."
+ For cach redshift: slice we combined the results rom the three survey regions., For each redshift slice we combined the results from the three survey regions.
+ The best-littine moclels srocluce a good fit to the measured. power spectrum. as evidenced by the values of the chi-squarecl statistic in the hree redshift) slices: v—(1113.9.1229.6.1069.2) for 1137 degrees of freedom.," The best-fitting models produce a good fit to the measured power spectrum, as evidenced by the values of the chi-squared statistic in the three redshift slices: $\chi^2 = (1113.9, 1229.6, 1069.2)$ for 1137 degrees of freedom."
+ The measurements of the growth rates at each ellective redshift’ sap=(0.42.0.59.0.78) ave f=(0.73x0.09.0.75+1.09.0.71 £0.14). mareinalizing over the parameters σε aud b. where l-o error ranges are indicated.," The measurements of the growth rates at each effective redshift $z_{\rm eff} = (0.42, 0.59, 0.78)$ are $f = (0.73 \pm 0.09, 0.75 \pm
+0.09, 0.71 \pm 0.14)$ , marginalizing over the parameters $\sigma_v$ and $b$, where $\sigma$ error ranges are indicated."
+" These should be compared to the ACDAL prediction for O4,(0)=0.3: f=(0.72.0.78. 0.83)."," These should be compared to the $\Lambda$ CDM prediction for $\Omega_{\rm m}(0) = 0.3$: $f = (0.72,0.78,0.83)$ ."
+ We find that the model successfully. passes his test of sell-consistency., We find that the model successfully passes this test of self-consistency.
+ Our measurements of the growth rate are compared with the model in Figure 15 hy plotting he quantity fCOxauass(2) versus ος.," Our measurements of the growth rate are compared with the model in Figure \ref{figgrowth} by plotting the quantity $f \times \sigma_{8,{\rm mass}}(z)$ versus redshift."
+ Previouslv- results are also indicated.as summarized by Song Percival (2009).," Previously-published results are also indicated,as summarized by Song Percival (2009)."
+ The WigeleZ Survey data provides a growth measurement across the redshift range θε«2<05 with a precision similar to previous workat z« 0.4., The WiggleZ Survey data provides a growth measurement across the redshift range $0.4 < z < 0.8$ with a precision similar to previous workat $z < 0.4$ .
+ We highlight in particular the 20% accuracy of our measurement, We highlight in particular the $20\%$ accuracy of our measurement
+"From the exact solutions. we know that rg(0)€ryp(a/2)=rggo. Equation (AL0)}) therefore suggests that for each model. there is a minimum augle between By aud vo: Since. lor a given Byoy (or %,). sinall 0 corresponds to large upstream magnetic field strength. a sliock is not possible for very small 0.","From the exact solutions, we know that $r_f (\theta) \leq r_f (\pi/2) \equiv r_{f,90}$, Equation \ref{rBf}) ) therefore suggests that for each model, there is a minimum angle between $\mathbf{B}_0$ and $\mathbf{v}_0$: Since, for a given $B_{y,0}$ (or $\beta_y$ ), small $\theta$ corresponds to large upstream magnetic field strength, a shock is not possible for very small $\theta$."
+" More practically. this cau also be written as Or 0>Onin (assuming sin(yi,«1/ Ou) for where exdoud.=Betout/¥πρι."," More practically, this can also be written as or $\theta > \theta_\mathrm{min}$ (assuming $\sin\theta_\mathrm{min} \ll 1/\sin\theta_\mathrm{min}$ ) for where $v_{\mathrm{A,cloud}} = B_\mathrm{cloud}/\sqrt{4\pi\rho_0}$."
+" The reason for the condition 0>Oy, is to ensure that the iullow is strong enough to produce shocks iu the magnetized gas.", The reason for the condition $\theta > \theta_\mathrm{min}$ is to ensure that the inflow is strong enough to produce shocks in the magnetized gas.
+ To obtain rr(0). we need to simultaneously solve and Equation (ALO)). which can only be done numerically.," To obtain $r_f (\theta)$, we need to simultaneously solve and Equation \ref{rBf}) ), which can only be done numerically."
+ Alternatively. we can also use Equation (ALL)) to write (assuming M7?o»rp)> 1) Substituting Equation (A15)) into Equation (ALO)) gives us a quadratic equation for ry(0): Since Mo>1. keeping only ME| terms gives This is an analytical approximation for rj(0) (see Fig. 10)).," Alternatively, we can also use Equation \ref{rBf_rf}) ) to write (assuming ${\cal M}^2 \gg r_f(\theta) \gg 1$ ) Substituting Equation \ref{rBfApprox}) ) into Equation \ref{rBf}) ) gives us a quadratic equation for $r_f(\theta)$ : Since ${\cal M}\gg 1$, keeping only ${\cal M}^{-1}$ terms gives This is an analytical approximation for $r_f(\theta)$ (see Fig. \ref{ObShock}) )."
+ The compression [actor rpoo for the case with magnetic field parallel tothe shock front. Gand— 2€) is roo7VER! (see," The compression factor $r_{f,90}$ for the case with magnetic field parallel tothe shock front $\tan\theta\rightarrow\infty$ ) is $r_{f,90} \approx \sqrt{\beta_y}{\cal M}$ (see"
+et al.,et al.
+ 1993). at a ealactoceutric radius of a few huudred pe or less.," 1993), at a galactocentric radius of a few hundred pc or less."
+ De Vaucouleurs (1961) sugeestedOO that [1915 contains a central bar. which provides a very efficient mechanisin to quickly trausport imatter towards the uucleus while releasing angular momentum outwards.," De Vaucouleurs (1964) suggested that 4945 contains a central bar, which provides a very efficient mechanism to quickly transport matter towards the nucleus while releasing angular momentum outwards."
+ Can such a bar explain the velocity anomalics outlined in 33.5.1?, Can such a bar explain the velocity anomalies outlined in 3.5.4?
+ A classical signature of a bar is an S-shaped distortion in the isovelocity contours of the gas (e.g. Ikaluajs 1975: Duval Monnet 1985)., A classical signature of a bar is an S-shaped distortion in the isovelocity contours of the gas (e.g. Kalnajs 1978; Duval Monnet 1985).
+ This is observed iu NGCLOLS (soe refFIG.IILDRANDTMODELaa: 33.5.1)., This is observed in NGC4945 (see \\ref{FIG.HI.BRANDTMODEL}a a; 3.5.4).
+ Tf a bar is preseut. the velocity residuals in refFICG.IILDRANDTMODELec could be interpreted. in ternis of svstematic overall eas flows along it. approaching at positive offsets. receding at negative offsets.," If a bar is present, the velocity residuals in \\ref{FIG.HI.BRANDTMODEL}c c could be interpreted in terms of systematic overall gas flows along it, approaching at positive offsets, receding at negative offsets."
+ The orbit velocities in the bar near the nucleus may be estimated frou the py diagrams aud the velocity field residuals., The orbit velocities in the bar near the nucleus may be estimated from the pv diagrams and the velocity field residuals.
+ The contimmun and CO distributions in ΤΟΕΙMAPS and EP show cuhauced emission extending out from the uucleus alone PA ~ ie. along the axis that is also showing the velocity anomalies refFIC.IILDR ANDTMODELcc).," The continuum and CO distributions in \\ref{FIG.MAPS} and \ref{FIG.CO21.AT43} show enhanced emission extending out from the nucleus along $PA$ $\sim$ $^{\circ}$, i.e. along the axis that is also showing the velocity anomalies \\ref{FIG.HI.BRANDTMODEL}c c)."
+ Assunüug that the putative bar is associated with the enission Ίος introducedl above aud tha this is located within the plane of the ealasx. we obtain an azimuthal angele of ~Lo (counterclockwise) with respect to the line-of-seht projected onto the plane of the ealaxy.," Assuming that the putative bar is associated with the emission ridge introduced above and that this is located within the plane of the galaxy, we obtain an azimuthal angle of $\sim40\degr$ (counterclockwise) with respect to the line-of-sight projected onto the plane of the galaxy."
+ The total inclination to the line-ofsight would then be , The total inclination to the line-of-sight would then be $^{\circ}$.
+There exists a stroug correlation between he unclear absorbing column density aud the presence of a bar iu Sevtert 2 ealaxies., There exists a strong correlation between the nuclear absorbing column density and the presence of a bar in Seyfert 2 galaxies.
+ Strougly barred Sevtert 2 ealaxies have an average Ny that is two orders of maeuitude higher than in nou-harred Sv 2s., Strongly barred Seyfert 2 galaxies have an average $N_{\rm H}$ that is two orders of magnitude higher than in non-barred Sy 2s.
+ More than of the ‘Compton thick Sevfert 24 (Απ x 107!cnisqi most of this columm density must arise from the iunnennost few ppc) are barre and alinost of these are stronglv barred (Miolino e al., More than of the `Compton thick Seyfert 2s' $N_{\rm H}$ $\ga$ $^{24}$; most of this column density must arise from the innermost few pc) are barred and almost of these are `strongly' barred (Maiolino et al.
+ 1999)., 1999).
+ 11915 is a Compton thick Sevtert 2 galaxy. (c.g. Cainmazzi et al., 4945 is a Compton thick Seyfert 2 galaxy (e.g. Guainazzi et al.
+ 2000: Madejski et al., 2000; Madejski et al.
+ 2000)., 2000).
+ Nonduteractiug spiral salaxies with moderate inclination aud substantial far infrared ciission (for the ceutral- part of 11915. Leug~ 11012solum: IRAS. 1989) are known to show both (Martinet. Friccdli 1997).," Non-interacting spiral galaxies with moderate inclination and substantial far infrared emission (for the central part of 4945, $\sim$ $^{10}$; IRAS 1989) are known to show both (Martinet Friedli 1997)."
+ The presence of απλοτσο morphologies iu iudividual Sevfer ealaxies is positively correlated with thei tendency to exhibit enhanced star formüng activity (ALlaiolne ot al., The presence of asymmetric morphologies in individual Seyfert galaxies is positively correlated with their tendency to exhibit enhanced star forming activity (Maiolino et al.
+ 1997)., 1997).
+ Hence inespective of direct observational evidence. the presence o Pa bar mum 11915 is. at least statistically. very likely.," Hence irrespective of direct observational evidence, the presence of a bar in 4945 is, at least statistically, very likely."
+ The velocity dispersion of iu the ceutral region is also cousistent with the presence of a bar., The velocity dispersion of in the central region is also consistent with the presence of a bar.
+ Estimated for cach observed spectimm during the gaussiau fitting. 16 dispersou iu the outer galaxy is 10 nis. Which is consistent with the dispersion of gas along 1ο spiral arius du the presence of relatively uuifonià circular rotation.," Estimated for each observed spectrum during the gaussian fitting, the disperson in the outer galaxy is 10 –, which is consistent with the dispersion of gas along the spiral arms in the presence of relatively uniform circular rotation."
+" However. within 150"" of the nucleus iere is an elongated region where the dispersion is abovezaus."," However, within $\sim150''$ of the nucleus there is an elongated region where the dispersion is above."
+. This may reflect higher gas turbuleuce or a higher space density of eas clouds near the centre. but could also reflect fast motion on lightly eccentric orbits in à bar.," This may reflect higher gas turbulence or a higher space density of gas clouds near the centre, but could also reflect fast motion on highly eccentric orbits in a bar."
+ Since all this evidence is circtustautial. a definite answer to the question whether 11915 has a bar must come from the near-intrared (preferably K-band) image.," Since all this evidence is circumstantial, a definite answer to the question whether 4945 has a bar must come from the near-infrared (preferably K-band) image."
+ This method proved to be couclisive in the case of another nearby lughly inclined. southern starburst galaxy. 2253 (sec ee. Eugelbracht et al.," This method proved to be conclusive in the case of another nearby highly inclined southern starburst galaxy, 253 (see e.g. Engelbracht et al."
+ 1995)., 1998).
+ The aud ypy diagrams can be analysed usns linear resonance hneorv (e.g. Binney Tremaine 1987) to deduce the ocatious of various eravitational resonances within 15., The and pv diagrams can be analysed using linear resonance theory (e.g. Binney Tremaine 1987) to deduce the locations of various gravitational resonances within 4945.
+ I£ we assiuune the presence of a weak bar (the orbits can be described by a superposition of circular notion around the ceutre and smal oscillations caused x the non-anisviumetric part of the votential). the spiral and barred structure is constramed by the locations of hese resonances.," If we assume the presence of a `weak' bar (the orbits can be described by a superposition of circular motion around the centre and small oscillations caused by the non-axisymmetric part of the potential), the spiral and barred structure is constrained by the locations of these resonances."
+ For a perfectly edee-on barred spiral the ocations cui be onlv approximate. because the rotation curve will be affected bv the resonances as well as the xus de-projected size aud. augle to the line-ofsight.," For a perfectly edge-on barred spiral the locations can be only approximate, because the rotation curve will be affected by the resonances as well as the bar's de-projected size and angle to the line-of-sight."
+ In he case of 11915. the proposed bar. 33.5.5) is sufficiently displaced. from the major axis to provide a reasonable approximation.," In the case of 4945, the proposed bar 3.5.5) is sufficiently displaced from the major axis to provide a reasonable approximation."
+ For eas orbiting the centre of a galaxy with augular velocity Q=ΕΠ at radius B. the radial epievclc frequency αν Is expressed by Application of a Brandt rotation curve 11) iufers Taner and outer ‘Lindblad’ resonances IER aud OLR) occur when the pattern of speedthe bar Op=O.w/in iux OQ|ein. respectively (02 represeuts a barred potential. see e.c. 33.115 of Binney Tremaine 1987).," For gas orbiting the centre of a galaxy with angular velocity $\Omega = 
+V/R$ at radius R, the radial epicyclic frequency $\kappa$ is expressed by Application of a Brandt rotation curve 1) infers Inner and outer `Lindblad' resonances (ILR and OLR) occur when the pattern speed of the bar $\Omega_{\rm P} = \Omega-\kappa/m$ and $\Omega+\kappa/m$, respectively $m$ =2 represents a barred potential, see e.g. 3–115 of Binney Tremaine 1987)."
+refFIC.VRAD.DYN shows the measured variation of aneular velocity as a function of ealactoceutric radius.,\\ref{FIG.VRAD.DYN} shows the measured variation of angular velocity as a function of galactocentric radius.
+ A rotation curve of the form given in 11 with the parameters outlined in Table 5 has been fitted to velocities derived from the py diagram in COMP.PVDIAGRAAL., A rotation curve of the form given in 1 with the parameters outlined in Table \ref{TAB.HI.BRANDT} has been fitted to velocities derived from the pv diagram in \\ref{FIG.COMP.PVDIAGRAM}.
+ The fitted curve provides a reasonable ft to the outer galaxy. {κ kkpe)., The fitted curve provides a reasonable fit to the outer galaxy $R$ $\ga$ kpc).
+For the purpose of this paper we are not interested in the polarization. but rather the total [lax al various wavelengths.,"For the purpose of this paper we are not interested in the polarization, but rather the total flux at various wavelengths."
+ Therelore we adapted the data to make them suitable for this study., Therefore we adapted the data to make them suitable for this study.
+ Since the three images were taken at the same wawelength. we averaged (hem using the IRAF routine INICOAMBINE.," Since the three images were taken at the same wavelength, we averaged them using the IRAF routine IMCOMBINE."
+ Assiuning that the light from the bright knots is unpolarized —a reasonable assumption— (he polarizing filters will transmit half of (he total [Iux., Assuming that the light from the bright knots is unpolarized —a reasonable assumption— the polarizing filters will transmit half of the total flux.
+ Therefore the [Iuxes derived from these data using (he standard photometric calibration [actors were multiplied bv a factor of (vo in oxder to work out the total fixes., Therefore the fluxes derived from these data using the standard photometric calibration factors were multiplied by a factor of two in order to work out the total fluxes.
+" The absolute photometric accuracy of the [/96 relav is 1 20% ο,", The absolute photometric accuracy of the f/96 relay is 10 – 20 \cite{Baum94}.
+ The WEPC2 observations and their reduction are described in detail bv ?.., The WFPC2 observations and their reduction are described in detail by \cite{Surace98}.
+ The PINS1345+12 images were centred in the PC chip (0.046 areseconds pixel |)., The PKS1345+12 images were centred in the PC chip $0.046$ arcseconds $^{-1}$ ).
+ The fillers used were F439W and FSIJW. which are very similar to the standard Johnson B and Cousins I filters (??)..," The filters used were F439W and F814W, which are very similar to the standard Johnson B and Cousins I filters \citep{Holtzman95a, Holtzman95b}."
+ Dark subtraction. bias subtraction aud flat fielding were carried out using the standard data," Dark subtraction, bias subtraction and flat fielding were carried out using the standard data"
+"mean free path model, for which Aj.(xui=0.1)32 Mpc/h, is the same as that for the corresponding fixed mean free path model with As,=32 Mpc/h, Ao=170 Mpc.","mean free path model, for which $\lambda_{\rm abs}(x_{\rm
+ HI}=0.1)\simeq 32$ $/h$, is the same as that for the corresponding fixed mean free path model with $\lambda_{\rm abs}=32$ $/h$, $\lambda_0=170$ Mpc."
+ Fig., Fig.
+" 4 shows the ionization field in a 5-Mpc/h slice at x20.75 for Mai,=10°Mo, with Ay,=256 Mpc/h (top) and Aabs=8 Mpc/R (bottom)."," \ref{panels} shows the ionization field in a $/h$ slice at ${x}=0.75$ for $M_{\rm min}=10^8 M_\odot$, with $\mfp=256$ $/h$ (top) and $\mfp=8$ $/h$ (bottom)."
+" The two cases are quite different, with many more small neutral patches for Agbs=8 Mpc/h relative to Aj,=256 Mpc/h."," The two cases are quite different, with many more small neutral patches for $\mfp=8$ $/h$ relative to $\mfp=256$ $/h$."
+ Also shown in 4, Also shown in Fig.
+ is the reioniztion redshift in a slice., \ref{panels} is the reioniztion redshift in a $/h$ slice.
+" Lower Fig.values of A45, result in a more extended 0.5-Mpc/hreionization overlap period — i.e. the maxima in reionization redshift are higher, and the minima are lower."," Lower values of $\mfp$ result in a more extended reionization overlap period – i.e. the maxima in reionization redshift are higher, and the minima are lower."
+" We have carried out large-scale simulations of reionization in a 2 Gpc/h volume, including a finite mean free path to absorption systems."," We have carried out large-scale simulations of reionization in a 2 $/h$ volume, including a finite mean free path to absorption systems."
+ Absorption systems have a significant effect on the characteristic scales at the end of reionization., Absorption systems have a significant effect on the characteristic scales at the end of reionization.
+" For the Thomson scattering optical depth reported by WMAP, Tes£x 0.09, we find that the characteristic bubble size when the universe is 90 cent ionized is sensitive to Aabs, with Ao~440 Mpc perfor Aabs=256 Mpc/h quiteand Mmin=10°Mo, while Ao~94 Mpc for Abs=8 Mpc/h and Minin= 10°Mo."," For the Thomson scattering optical depth reported by , $\tau_{es}\simeq 0.09$ , we find that the characteristic bubble size when the universe is 90 per cent ionized is quite sensitive to $\mfp$, with $\lambda_0\sim 440$ Mpc for $\mfp=256$ $/h$ and $M_{\rm
+ min}=10^9 M_\odot$, while $\lambda_0\sim 94$ Mpc for $\mfp=8$ $/h$ and $M_{\rm min}=10^8 M_\odot$ ."
+" Calculations using a sharp k-space filter lead to a more modest extension of the percolation phase (see Table 1) compared to the reionization history obtained by the solution of equation (1)), on which the reionization histories shown in Figure 1 are based."," Calculations using a sharp $k$ -space filter lead to a more modest extension of the percolation phase (see Table 1) compared to the reionization history obtained by the solution of equation \ref{dxdt}) ), on which the reionization histories shown in Figure 1 are based."
+" For τος~0.09, the difference in the redshift when the universe is 90 per cent ionized between the Aj=8 and λωυς=256 cases is only Azo~0.3, while Mpc/hthe corresponding case Mpc/husing a sharp real space filter and equation (1)) is Azo.9~0.7."," For $\tau_{es}\simeq 0.09$, the difference in the redshift when the universe is 90 per cent ionized between the $\mfp=8$ $/h$ and $\mfp=256$ $/h$ cases is only $\Delta z_{0.9}\sim 0.3$, while the corresponding case using a sharp real space filter and equation \ref{dxdt}) ) is $\Delta z_{0.9}\sim 0.7$."
+" The delay in the overlap time in both cases is Azo,~1.5.", The delay in the overlap time in both cases is $\Delta z_{\rm ov}\sim 1.5$.
+ Using either sharp k-space smoothing or equation (1) to model the effect of absorption systems on the reionization history has drawbacks., Using either sharp $k$ -space smoothing or equation (1) to model the effect of absorption systems on the reionization history has drawbacks.
+" While equation (1) conserves photons in the long mean free path limit, a mean flux and opacity are assumed, and further work will be necessary to validate and/or improve upon the approach layed out in section 2.1.."," While equation (1) conserves photons in the long mean free path limit, a mean flux and opacity are assumed, and further work will be necessary to validate and/or improve upon the approach layed out in section \ref{sec:history}."
+" Using a sharp k-space filter results in a reionization history which also conserves photons in the long mean free path limit (kr—>0), but at the expense of using an oscillatory filter which can “leak” photons from high density regions into lower density ones, and for which there is no unique choice for the relation between Ἆαυς and kr."," Using a sharp $k$ -space filter results in a reionization history which also conserves photons in the long mean free path limit $k_F\longrightarrow 0$ ), but at the expense of using an oscillatory filter which can “leak” photons from high density regions into lower density ones, and for which there is no unique choice for the relation between $\mfp$ and $k_F$."
+" While our results for the timing and photon consumption of the end of reionizaiton are obtained from equations (1) and (9), it is important to note that more work will be necessary to test their accuracy."," While our results for the timing and photon consumption of the end of reionizaiton are obtained from equations (1) and (9), it is important to note that more work will be necessary to test their accuracy."
+" However, results for the of reionization (Figures 1 and 4) do not depend morphologyon the model for the global history, and are therefore more robust."," However, results for the morphology of reionization (Figures 1 and 4) do not depend on the model for the global history, and are therefore more robust."
+" Our results are consistent with those of Furlanetto&Oh(2005), in which of ionizing photons in dense systems extends consumptionthe end of reionization considerably."," Our results are consistent with those of \citet{furlanetto/oh:2005}, in which consumption of ionizing photons in dense systems extends the end of reionization considerably."
+" Choudhuryetal.(2009) used the semi-numerical approach in a volume 100 Mpc/h across, and found a similar trend with decreasing Aap, indicating a transition to an “outside-in” morphology near the end of reionization."," \citet{choudhury/etal:2009} used the semi-numerical approach in a volume $100$ $/h$ across, and found a similar trend with decreasing $\mfp$, indicating a transition to an “outside-in” morphology near the end of reionization."
+" We find that neutral patches may also remain in voids, where formation of ionizing sources is delayed and radiation from the nearest sources is shielded."," We find that neutral patches may also remain in voids, where formation of ionizing sources is delayed and radiation from the nearest sources is shielded."
+" Previous studies that modeled the physical origin of absorption systems as minihalos (e.g.,Ciardietal.2006;McQuinnetal.2007) also found similar effects to those we find here, although these studies were more focused on the intermediate stages of reionization and on smaller simulated volumes, where the photon mean free path was not as small relative to the bubble sizes as we find in our low mean free path cases in the final, percolation phase."," Previous studies that modeled the physical origin of absorption systems as minihalos \citep[e.g.,][]{ciardi/etal:2006,mcquinn/etal:2007} also found similar effects to those we find here, although these studies were more focused on the intermediate stages of reionization and on smaller simulated volumes, where the photon mean free path was not as small relative to the bubble sizes as we find in our low mean free path cases in the final, percolation phase."
+" As mentioned in the introduction, our approach is complementary to that taken by Crocianietal.(2011),, in which the semi-numerical approach was used to determine the distribution of absorbers during reionization."," As mentioned in the introduction, our approach is complementary to that taken by \citet{crociani/etal:2011}, in which the semi-numerical approach was used to determine the distribution of absorbers during reionization."
+" They found that their spatial distribution is quite inhomogenous, owing both to intrinsic density fluctuations in the IGM, as well as fluctuations in the ionizing radition field — regions far from sources have a relatively low flux, resulting in a higher abundance of absorbers."," They found that their spatial distribution is quite inhomogenous, owing both to intrinsic density fluctuations in the IGM, as well as fluctuations in the ionizing radition field – regions far from sources have a relatively low flux, resulting in a higher abundance of absorbers."
+" This affect was also pointed out by McQuinnetal.(2011), who used high-resolution simulations post-processed with radiative transfer to hydrodynamicdetermine the mean free path as a function of flux, finding a strongly nonlinear dependence of the mean free path on the emmissivity of ionizing radiation."," This affect was also pointed out by \citet{mcquinn/etal:2011}, who used high-resolution hydrodynamic simulations post-processed with radiative transfer to determine the mean free path as a function of flux, finding a strongly nonlinear dependence of the mean free path on the emmissivity of ionizing radiation."
+" These effects imply an important improvement to our model is not only a time-varying mean free path, but also a spatially varying one."," These effects imply an important improvement to our model is not only a time-varying mean free path, but also a spatially varying one."
+" To accomplish this, more work wil need to be done on the ""sub-grid"" physics during reionization, using the results of high-resolution cosmological simulations with radiative transfer of a background radiation field, coupled to the hydrodynamics of the gas."," To accomplish this, more work will need to be done on the “sub-grid” physics during reionization, using the results of high-resolution cosmological simulations with radiative transfer of a background radiation field, coupled to the hydrodynamics of the gas."
+" However, given that the final to be ionized are those most distant from the most luminous patchessources, the delay in the very end of reionizationthatwe find here is likely to persist when the inhomogeneity of the IGM opacity due to absorption systems is taken into account."," However, given that the final patches to be ionized are those most distant from the most luminous sources, the delay in the very end of reionizationthatwe find here is likely to persist when the inhomogeneity of the IGM opacity due to absorption systems is properly taken into account."
+ The properlylingering neutral clouds wefind would further complicate interpretation of quasar absorption spectra at z~~6 , The lingering neutral clouds wefind would further complicate interpretation of quasar absorption spectra at $z\sim 6$ 
+We calculate the curent spatial clistribution of projectile delivery to the Earth and. Moon uxing uumerical orbital dyuamics simulatious of caudidate impactors drawn Crom a debiased (NEO) model.,We calculate the current spatial distribution of projectile delivery to the Earth and Moon using numerical orbital dynamics simulations of candidate impactors drawn from a debiased Near-Earth-Object (NEO) model.
+ Surprisingly. we find that the average lunar impact velocity is 20 kin/s. which has r:unifications ln couvertiug observed crater deusities to impactor size distributions.," Surprisingly, we find that the average lunar impact velocity is 20 km/s, which has ramifications in converting observed crater densities to impactor size distributions."
+ We determiue that curreut crate production ou the leadiug hemisphere of the Moon is 1.292:0.01 that of the trailiig whe1 considering the ratio of craters within of the apex to those within of he antapex and that there is virtually no uearside-Darside asymunetry., We determine that current crater production on the leading hemisphere of the Moon is $1.29 \pm 0.01$ that of the trailing when considering the ratio of craters within of the apex to those within of the antapex and that there is virtually no nearside-farside asymmetry.
+ As expected. the degree of eadiug-trailing asyiumetryus increases when the Moou's orbital distauce is decreased.," As expected, the degree of leading-trailing asymmetry increases when the Moon's orbital distance is decreased."
+ We examine he latitude distribution of inpactor sites aud find that for both the Earth aud Moon there is a sinal delicieucy of time-averaged iipact rates at the poles., We examine the latitude distribution of impactor sites and find that for both the Earth and Moon there is a small deficiency of time-averaged impact rates at the poles.
+ The ratio between deliveries within of the vole to that of a baud ceitered on the equator is nearly unity for Earth (<1% )(0.992+0.001 rut detectaJy. non-uniform for the Moon (ον 105€) (0.912+ 0.001).," The ratio between deliveries within of the pole to that of a band centered on the equator is nearly unity for Earth $<
+1\%$ $0.992 \pm 0.001$ ) but detectably non-uniform for the Moon $\sim 10\%$ ) $0.912 \pm 0.004$ )."
+ The terrestrial arrival results are examined to determine te degree of AM/PM asynuuetry to compare with meteorite fall times (Cof which tlere seems to bea PAL excess)., The terrestrial arrival results are examined to determine the degree of AM/PM asymmetry to compare with meteorite fall times (of which there seems to be a PM excess).
+ Our results show that the impact [ux of objects derived from the NEOs in the AM |Ours is  2 times that of the PAL hemisphere. further supporting the assertion that imeteorite-droping objects are recent. ejectious from the main asteroid belt rather than vouug fragments of NEOs.," Our results show that the impact flux of objects derived from the NEOs in the AM hours is $\sim$ 2 times that of the PM hemisphere, further supporting the assertion that meteorite-dropping objects are recent ejections from the main asteroid belt rather than young fragments of NEOs."
+ Cratering: Earh: Meteorites: Moon. surface: Near-Earth objects Subimittecd toIcarus: July 28 2006 Send correspondance to B.C: gladinauephias.ubc.ca," Cratering; Earth; Meteorites; Moon, surface; Near-Earth objects Submitted to: July 28 2006 Send correspondance to B.G.: gladmanphas.ubc.ca"
+orbital velocity and combining this with Iepler’s third law. we ect: Using the masses of both stars. we obtain P = 32 c 6 davs.,"orbital velocity and combining this with Kepler's third law, we get: Using the masses of both stars, we obtain $P$ = 32 $\pm$ 6 days."
+ This is cousiderablv shorter than the nüniuun orbital period we interred from currently available observations. which implies au eccentric orbit.," This is considerably shorter than the minimum orbital period we inferred from currently available observations, which implies an eccentric orbit."
+ Frou observations obtained withAÁepler. we found a pulsating red giant in an eclipsing binary.," From observations obtained with, we found a pulsating red giant in an eclipsing binary."
+" Once a ull orbit has been observed. the orbital paramcters will provide an independent measure of the stellar xuwanieters, such as mass aud radius. with respect to the asteroseisunüc values."," Once a full orbit has been observed, the orbital parameters will provide an independent measure of the stellar parameters, such as mass and radius, with respect to the asteroseismic values."
+ These additional constraints male his a very interesting case aud the star is therefore beime ollowed up by the as well as with eroundl-sed spectroscopy., These additional constraints make this a very interesting case and the star is therefore being followed up by the as well as with ground-based spectroscopy.
+ Stellar parameters of the primary red giant lave con computed from a spectroscopic analysis aud from he solu-like oscillations., Stellar parameters of the primary red giant have been computed from a spectroscopic analysis and from the solar-like oscillations.
+ We did not fiud evidence of he secondary compoucut in the spectra., We did not find evidence of the secondary component in the spectra.
+ The currently available spectra are not vet sufficient to put constraints ou the orbit., The currently available spectra are not yet sufficient to put constraints on the orbit.
+" From the single eclipse observed by and additional iudepeudeut observations from το, SuperWASP and ASAS. we inferred coustraints ou the orbit and secondary star."," From the single eclipse observed by and additional independent observations from TrES, SuperWASP and ASAS, we inferred constraints on the orbit and secondary star."
+ The aunular eclipse. ie. the secondary passing iu front of the red eiut. would cause an eclipse with a depth of ~ assmuing a circular orbit aud without taking lmnib-darkeniug iuto account.," The annular eclipse, i.e., the secondary passing in front of the red giant, would cause an eclipse with a depth of $\sim$, assuming a circular orbit and without taking limb-darkening into account."
+ Such a dip would be observable. but has not oen detected byKepler... TrES or ASAS.," Such a dip would be observable, but has not been detected by, TrES or ASAS."
+ Only. the SuperWASP data show a feature that might be due to an annular eclipse., Only the SuperWASP data show a feature that might be due to an annular eclipse.
+ The SuperWASP and ASAS data lave gaps due to weather conditions., The SuperWASP and ASAS data have gaps due to weather conditions.
+ Although there are also three gaps (L7. 2.9 and 1.2dd) iu the TrES time series data with a duration louger than the duration of the eclipse. we expect the chance that an eclipse falls exactly during one of these gaps to be low.," Although there are also three gaps (4.7, 2.9 and d) in the TrES time series data with a duration longer than the duration of the eclipse, we expect the chance that an eclipse falls exactly during one of these gaps to be low."
+ Therefore. we inter that the orbit is longer than the time spin of the TrES data. ie. at least 75dd. Also. au initial inspection of the raw data of Q2 (additional 3 mouths of data following the QI phase) does not sec to reveal another eclipse.," Therefore, we infer that the orbit is longer than the time span of the TrES data, i.e., at least d. Also, an initial inspection of the raw data of Q2 (additional 3 months of data following the Q1 phase) does not seem to reveal another eclipse."
+ In the case that we missed an eclipse due to a lack of observations. the wait orbital period would be 135 davs.," In the case that we missed an eclipse due to a lack of observations, the minimum orbital period would be 135 days."
+ From the orbital velocity. of the secondary. we conrputed the period for a circular orbit. which would be shorter than half the total time span of the TYES data and we should have seen au occultation iu those data.," From the orbital velocity of the secondary, we computed the period for a circular orbit, which would be shorter than half the total time span of the TrES data and we should have seen an occultation in those data."
+ Therefore we concluded that the orbit is eccentric., Therefore we concluded that the orbit is eccentric.
+ The possible annular eclipse iu the SuperWASP data has a leneth of ~77 ddavs aud depth of ~ 00.03 mae., The possible annular eclipse in the SuperWASP data has a length of $\sim$ days and depth of $\sim$ 0.03 mag.
+ If confirmed this would indeed nuplv a large orbital eccentricity., If confirmed this would indeed imply a large orbital eccentricity.
+ From the eclipse times aud depth we were able to compute the Iuninosityv and the radius of the secondary., From the eclipse times and depth we were able to compute the luminosity and the radius of the secondary.
+ All values seein to be compatible withan FE sequence star., All values seem to be compatible withan F main-sequence star.
+ The mass-huninosity relation for mai-sequence stars then leads to a mass estimate of the secondary as well as its surface gravity., The mass-luminosity relation for main-sequence stars then leads to a mass estimate of the secondary as well as its surface gravity.
+ These results are «παλσος in Table 30 and the position of both stars are jucdicated in an U-R diagram in Fig. 2.., These results are summarized in Table \ref{allres} and the position of both stars are indicated in an H-R diagram in Fig. \ref{fluxosc}.
+ RICSILIOG37 is an extremely intercsting binary for further follow-up., KIC8410637 is an extremely interesting binary for further follow-up.
+ Longer time series from will inuprove dramatically the detection threshold for the derivation of the properties of the oscillations of the red edant., Longer time series from will improve dramatically the detection threshold for the derivation of the properties of the oscillations of the red giant.
+ Funding for this Discovery iissiou is provided by NASA’s Science Mission Directorate., Funding for this Discovery mission is provided by NASA's Science Mission Directorate.
+ We would like to acknowledge the cutire tear for their efforts over μα vears, We would like to acknowledge the entire team for their efforts over many years.
+ Without these efforts it would not have been possible to obtain the results prescuted here., Without these efforts it would not have been possible to obtain the results presented here.
+ SIL WIC. YPE. IRS and DWI) acknowlecdec support by the UK Science aud Technology Facilities Council.," SH, WJC, YPE, IRS and DWK acknowledge support by the UK Science and Technology Facilities Council."
+ The research leading to these results has received finding from the Enropeau Research Council nuder the European Conumuitvs Seventh Framework Proerauune (EPT/20072013)/ERC evant aerecucnt 1n722722| (PROSPERITY). from the Research Council of INULeuven. from the Fund for Scientific Research of Flanders (FWO). and from the Beleian Federal Science Office.," The research leading to these results has received funding from the European Research Council under the European Community's Seventh Framework Programme (FP7/2007–2013)/ERC grant agreement $^\circ$ 227224 (PROSPERITY), from the Research Council of K.U.Leuven, from the Fund for Scientific Research of Flanders (FWO), and from the Belgian Federal Science Office."
+ DS acknowledges support from the Australian Research Council. Mission., DS acknowledges support from the Australian Research Council. .
+LGAL is detected. it will be impossible to infer wy from a single GRB to better than ary~0.3 because of the patchiness of reionisation.,"IGM is detected, it will be impossible to infer $\bar{x}_H$ from a single GRB to better than $\delta \bar{x}_H \sim 0.3$ because of the patchiness of reionisation."
+ Assuming an observation with similar sensitivity to the Subaru POCAS spectrum of CRBO50904. that. the clistribution of DLAs is the same as found at lower redshifts. and that the redshift of the CURB is known. a GRB from a redshift at which UH=0.5 can be used to detect a partly neutral IGM at 98% CLL. mVA of the time (and. for an observation with3 times ivity. &30% of the time).," Assuming an observation with similar sensitivity to the Subaru FOCAS spectrum of GRB050904, that the distribution of DLAs is the same as found at lower redshifts, and that the redshift of the GRB is known, a GRB from a redshift at which $\bar{x}_H \approx 0.5$ can be used to detect a partly neutral IGM at $98\%$ C.L. $\approx 10\%$ of the time (and, for an observation with $3$ times the sensitivity, $\approx
+30\%$ of the time)."
+ μμ<O.5. these percentiles for detection are even smaller.," If $\bar{x}_H < 0.5$, these percentiles for detection are even smaller."
+ Weaker DLAs enhance the probability of detecting a neutral IGAL but too weak of à DLA may prevent a precise redshift determination. which is essential for tight constraints on wry.," Weaker DLAs enhance the probability of detecting a neutral IGM, but too weak of a DLA may prevent a precise redshift determination, which is essential for tight constraints on $\bar{x}_H$ ."
+" Since the z=6.3 burst CRBO5S0904 has a DLA with Nyy=1077"" em7. the absorption on the red side of the line is dominated bv the DLA Clotanietal.2006)."," Since the $z = 6.3$ burst GRB050904 has a DLA with $N_{\rm HI} =
+10^{21.6}$ $^{-2}$, the absorption on the red side of the line is dominated by the DLA \citep{totani06}."
+.. While this burst may favour a model with an ionised universe over a neutral universe (Yotanictal.2006)... a weaker DLA is necessary to be able to constrain wy in the spectrum of a high-redshift. GRB.," While this burst may favour a model with an ionised universe over a neutral universe \citep{totani06}, a weaker DLA is necessary to be able to constrain $\bar{x}_H$ in the spectrum of a high-redshift GRB."
+ This is particularly true if the CRB occurs within a large HILL region., This is particularly true if the GRB occurs within a large HII region.
+ GRBOS5O904 was observed. spectroscopically by. the Subaru telescope 3.4 days alter the prompt gamma rav emission., GRB050904 was observed spectroscopically by the Subaru telescope $3.4$ days after the prompt gamma ray emission.
+ If this afterglow had been observed. hours after he burst. the flux would have been more than an order of maenituce larger.," If this afterglow had been observed hours after the burst, the flux would have been more than an order of magnitude larger."
+ To detect a neutral GM. it is crucial for optical ancl near-infrared spectrographs to observe candidate igh-redshift GRBs as soon as possible after the promp eamm rav enission.," To detect a neutral IGM, it is crucial for optical and near-infrared spectrographs to observe candidate high-redshift GRBs as soon as possible after the prompt gamma ray emission."
+ A high signal-to-noise ratio is critica o distinguish LGAL absorption from that arising from a DLA., A high signal-to-noise ratio is critical to distinguish IGM absorption from that arising from a DLA.
+ Such an observing programme is worthwhile given he promise that GRBs have as probes of the epoch of reionisation., Such an observing programme is worthwhile given the promise that GRBs have as probes of the epoch of reionisation.
+ We would like to thank T. Totani for providing the Subaru afterglow data for GIUDOSDC)004 as well as for his comments on the manuscript and for answering numerous questions about the measurement., We would like to thank T. Totani for providing the Subaru afterglow data for GRB050904 as well as for his comments on the manuscript and for answering numerous questions about the measurement.
+ We also thank Mark. Dijkstra. Claucde-Xndré Faucher-Ciguerre. Alexandre Tehekhovskoi. and Liv Trac for many interesting cliscussions and the anonymous referee for useful comments on the manuscript.," We also thank Mark Dijkstra, Claude-André Faucher-Giguèrre, Alexandre Tchekhovskoi, and Hy Trac for many interesting discussions and the anonymous referee for useful comments on the manuscript."
+ AMAL acknowledges support through an NSE graduate student. fellowship., MM acknowledges support through an NSF graduate student fellowship.
+ Ehe authors are also supported by the David and Lucile Packare the Alfred. P. Sloan Foundation. andgrants AST-0506556 AMand NNCGO5C.J40¢," The authors are also supported by the David and Lucile Packard Foundation, the Alfred P. Sloan Foundation, andgrants AST-0506556 and NNG05GJ40G."
+"The inverse-Compton luminosity due to electrons with given energy 554, Is Therefore. the adiabatic decrease of the EIC. Iuminosity at a given [requeney MeeX55, 18 The ratio of the observed SSC Iuninositv (o the observed EIC: luminosity is (e.g..Stawarzetal.2003).","The inverse-Compton luminosity due to electrons with given energy $\gamma_{\rm eic}$ is Therefore, the adiabatic decrease of the EIC luminosity at a given frequency $\nu_{\rm eic} \propto \gamma_{\rm eic}^2$ is The ratio of the observed SSC luminosity to the observed EIC luminosity is \citep[e.g.,][]{sta03}."
+. The comoving οποιον density of the CAIB radiation at redshift z is where aT!~4x101 erg ., The comoving energy density of the CMB radiation at redshift $z$ is where $a T^4 \sim 4 \times 10^{-13}$ erg $^{-3}$.
+ The comoving energv density of the svnchrotron enission in the clumpy jet. for clamps will radius A4. is The intrinsic svnchrotron humninositv of each chump. L'SMLC7 is related. to the observed svuchrotrvon Iuminosity of the whole knot region. Lay. by where Vis a number of identical clumps.," The comoving energy density of the synchrotron emission in the clumpy jet, for clumps with radius $R_{\rm c}$, is The intrinsic synchrotron luminosity of each clump, $L'_{\rm syn, \, c}$, is related to the observed synchrotron luminosity of the whole knot region, $L_{\rm syn}$, by where $N$ is a number of identical clumps."
+ Hence. one finally obtains," Hence, one finally obtains"
+of the integration as this will average out the scintillation and reduce the intensity variance.,of the integration as this will average out the scintillation and reduce the intensity variance.
+" The scintillation index given in equation 5 is only valid for very short exposures where there is no temporal averaging, i.e. the exposure time has to be less than the pupil crossing time of the intensity fluctuations."," The scintillation index given in equation \ref{eqn:scint_var} is only valid for very short exposures where there is no temporal averaging, i.e. the exposure time has to be less than the pupil crossing time of the intensity fluctuations."
+" The crossing time, t-, can be calculated as te=Dtei/Uw, where v, is the velocity of the turbulent layer."," The crossing time, $t_c$, can be calculated as $t_c = D_{\rm{tel}}/v_w$, where $v_w$ is the velocity of the turbulent layer."
+" If the exposure time, t, is greater than the crossing time then the scintillation index is modified to (Kenyonetal.2006),, where V(h) is the velocity of the turbulent layer at altitude h."," If the exposure time, $t$, is greater than the crossing time then the scintillation index is modified to \citep{Kenyon06}, where $V(h)$ is the velocity of the turbulent layer at altitude $h$."
+ Using this modification to the scintillation index we can calculate an example light curve for a fictional extrasolar planet transit for a given turbulence profile., Using this modification to the scintillation index we can calculate an example light curve for a fictional extrasolar planet transit for a given turbulence profile.
+ Figure 12 shows an example simulated extrasolar planet transit., Figure \ref{fig:planet_example} shows an example simulated extrasolar planet transit.
+ The transit depth is assumed to be 0.05 and has a duration of 2.5 hours., The transit depth is assumed to be 0.05 and has a duration of 2.5 hours.
+ A 2 m telescope and 30 s exposure time are also assumed., A 2 m telescope and 30 s exposure time are also assumed.
+ The optical turbulence profile used in the simulation is the same as that shown in figure 10 as measured by SCIDAR at San Pedro Mártir., The optical turbulence profile used in the simulation is the same as that shown in figure \ref{fig:scidar_prof_2000_05_19} as measured by SCIDAR at San Pedro $\acute{\mathrm{a}}$ rtir.
+ A wind speed of 5 ms~? for the surface layer and 20 ms~! for all other turbulence is assumed., A wind speed of 5 $^{-1}$ for the surface layer and 20 $^{-1}$ for all other turbulence is assumed.
+" The normalised scintillation noise in the visible is reduced from 0.70x10? (0.78 mmag) to 0.21x10? (0.23 mmag), an improvement factor of 3.3."," The normalised scintillation noise in the visible is reduced from $0.70\times 10^{-3}$ (0.78 mmag) to $0.21\times 10^{-3}$ (0.23 mmag), an improvement factor of 3.3."
+ If we assume a target magnitude of 11 then we have reduced the scintillation to a level which is comparable to the shot Although the aperture must be placed at the conjugate altitude of the turbulence the photometry can be done in the focal plane., If we assume a target magnitude of 11 then we have reduced the scintillation to a level which is comparable to the shot Although the aperture must be placed at the conjugate altitude of the turbulence the photometry can be done in the focal plane.
+ This means that we do not expect any of the other noise sources to increase as a result of implementing our conjugate-plane photometry technique., This means that we do not expect any of the other noise sources to increase as a result of implementing our conjugate-plane photometry technique.
+" 'The magnitudes of other noise sources, such as shot noise, readout noise or flat fielding noise, will depend on other factors."," The magnitudes of other noise sources, such as shot noise, readout noise or flat fielding noise, will depend on other factors."
+" There are three possible regimes in which we are interested: scintillation dominated, other noise dominated and a mixture of the two."," There are three possible regimes in which we are interested: scintillation dominated, other noise dominated and a mixture of the two."
+" In the first and last cases the noise will add in quadrature and so a reduction in scintillation noise by a factor of n will reduce the total noise to, στ,=Vei+o2:n,(4s—1), where or is the total noise before scintillation reduction."," In the first and last cases the noise will add in quadrature and so a reduction in scintillation noise by a factor of $n$ will reduce the total noise to, $\sigma_{T_{2}}=\sqrt{\sigma_{T}^{2} + \sigma_{\mathrm{scint}}^{2}\left(\frac{1}{n^2}-1\right)}$, where $\sigma_{T}$ is the total noise before scintillation reduction."
+ Figure 13 shows a 2D plot of the total noise reduction factor as a function of the telescope diameter and the target magnitude assuming the same parameters as before., Figure \ref{fig:mag_D_noise} shows a 2D plot of the total noise reduction factor as a function of the telescope diameter and the target magnitude assuming the same parameters as before.
+ The atmospheric model was the median profile from the SCIDAR data recorded on 2000 May 19., The atmospheric model was the median profile from the SCIDAR data recorded on 2000 May 19.
+ The optimum telescope size is found to be between 1.2 πι and 2 πι. Less than this and the diffraction effects limit the possible scintillation noise reduction and apertures greater than this become shot noise dominated., The optimum telescope size is found to be between 1.2 m and 2 m. Less than this and the diffraction effects limit the possible scintillation noise reduction and apertures greater than this become shot noise dominated.
+ In the latter scenario the scintillation noise is insignificant and so scintillation correction techniques will have no effect., In the latter scenario the scintillation noise is insignificant and so scintillation correction techniques will have no effect.
+ The median reduction in intensity variance for all available SCIDAR data collected over 24 nights in March/April 1997 and May 2000 at San Pedro Martir is a factor of 6., The median reduction in intensity variance for all available SCIDAR data collected over 24 nights in March/April 1997 and May 2000 at San Pedro $\acute{\mathrm{a}}$ rtir is a factor of 6.
+" However, with the limited data available it is difficult to say if this representative; it is possible that other times or sites will yield even better results if the turbulence is more constrained to stratified layers."," However, with the limited data available it is difficult to say if this representative; it is possible that other times or sites will yield even better results if the turbulence is more constrained to stratified layers."
+ Adaptive optics (AO) can be used to reduce the phase aberrations for imaging., Adaptive optics (AO) can be used to reduce the phase aberrations for imaging.
+ Here it is intensity fluctuations which are the problem and so AO systems can not directly reduce the scintillation., Here it is intensity fluctuations which are the problem and so AO systems can not directly reduce the scintillation.
+" However, AO systems can be used in conjunction with this technique to further reduce the scintillation."," However, AO systems can be used in conjunction with this technique to further reduce the scintillation."
+ As shown previously the surface turbulent layer is a major limitation to the conjugate-plane technique., As shown previously the surface turbulent layer is a major limitation to the conjugate-plane technique.
+the RGB observables in the (Aly.J—IX) IR plane.,"the RGB observables in the $\rm (M_K,J-K)$ IR plane."
+ Note that the validity of the relations has been extended up to [AL/LJ++0.4 dex. with respect to the results presented in Paper I. including the recent results on the metallicity of NGC 6791 (Origliaetal.2006).," Note that the validity of the relations has been extended up to ${\sim}$ +0.4 dex, with respect to the results presented in Paper I, including the recent results on the metallicity of NGC 6791 \citep{ori06}."
+. Hence the relations presented here have been obtained in the metallicity range: |[Fe/1I] 40.4 dex?.., Hence the relations presented here have been obtained in the metallicity range: ${\leq}$ ${\leq}$ +0.4.
+" The overall metallicity of a stellar population can be linked (o the slope of the RGB (slopercp) as defined in the near IR-CMD. by Ixuchinskietal.(1995). ancl (1995).. accordingly to the following relation: where slopejb, is the RGB slope measured. following the prescriptions discussed in Paper Lin the (Ας...)— IX)-CMD."," The overall metallicity of a stellar population can be linked to the slope of the RGB $\rm slope_{RGB}$ ) as defined in the near IR-CMD, by \citet{K95} and \citet{KF95}, accordingly to the following relation: where $\rm slope^{JK}_{RGB}$ is the RGB slope measured following the prescriptions discussed in Paper I in the $\rm (M_K,J-K)$ -CMD."
+ The RGDB--üp huninosity is a quantity well predicted by theoretical models and it can be easily measured in populous stellar svstenis like galaxies., The -tip luminosity is a quantity well predicted by theoretical models and it can be easily measured in populous stellar systems like galaxies.
+ The Iuminositv of the RGB tip for stellar populations older than /~1—2 Gyr is nearly independent from the population age., The luminosity of the RGB tip for stellar populations older than $t \sim 1-2$ Gyr is nearly independent from the population age.
+ Moreover. il turns out to be particularly bright in (he near--LR (Mg~—6) hence. it is a verv promising standard caadle up to large distances.," Moreover, it turns out to be particularly bright in the -IR $\rm M_K \sim -6$ ) hence, it is a very promising standard candle up to large distances."
+ The main limitation of this method is (he clear--cut. determination of the RGB—tip luminosity Chat could be seriously affected by low--uunmber statistics., The main limitation of this method is the -cut determination of the -tip luminosity that could be seriously affected by -number statistics.
+ In fact. (he rapid star evolution near the end of the RGB. makes the brightest portion of the RGB intrinsically poorly populated.," In fact, the rapid star evolution near the end of the RGB, makes the brightest portion of the RGB intrinsically poorly populated."
+ However. this problem is marginal in verv populous stellar svstems.," However, this problem is marginal in very populous stellar systems."
+ The dependence of this feature Irom the overall metallitv of the stellar population has been empirically determined by using a sample of Galactic GCs (see Paper 1I) and it turis out to be in good agreement wilh theoretical expectations., The dependence of this feature from the overall metallicity of the stellar population has been empirically determined by using a sample of Galactic GCs (see Paper II) and it turns out to be in good agreement with theoretical expectations.
+ Here we report the relation obtained in Paper II: As well known (seeSalaris&Girareli2002:GrocholskiSarajedini2002) for clusters older than 2 Gyr the level of the IHelium-burning Red clump (RC) is mainly influenced by metallicity and it shows little dependence on the age.," Here we report the relation obtained in Paper II: As well known \citep[see][]{salgir02,grosara02}
+ for clusters older than 2 Gyr the level of the Helium-burning Red clump (RC) is mainly influenced by metallicity and it shows little dependence on the age."
+ Hence for relatively old. metal rich. populations (0>2Gyr.[M/II] —1) it is also possible to use the level of the RC as an additional distance indicator.," Hence for relatively old metal rich populations $\rm t>2~Gyr, [M/H]>-1$ ) it is also possible to use the level of the RC as an additional distance indicator."
+ We estimate the mean level of the RC bv. using a metal-rich «Ρο «-0.3) cluster subsample (namely. 47 Tuc. NGC 6342. NGC 6380.," We estimate the mean level of the RC by using a metal-rich $<$ $<$ -0.3) cluster subsample (namely, 47 Tuc, NGC 6342, NGC 6380,"
+the total mass of the surrounding cloud.,the total mass of the surrounding cloud.
+ This was done bv summing the emission over the central part of the map (shown in Figure 3)) and subtracting the Dux already determined. for the compact sources., This was done by summing the emission over the central part of the map (shown in Figure \ref{regions}) ) and subtracting the flux already determined for the compact sources.
+ Table 4. shows the masses calculated from the extended emission at aand [for various assumed parameters., Table \ref{extmass} shows the masses calculated from the extended emission at and for various assumed parameters.
+ These masses are considerable. representing about half the mass of the system.," These masses are considerable, representing about half the mass of the system."
+ The nipass estimates are about a factor of 2 greater than those al450jn., The mass estimates are about a factor of 2 greater than those at.
+. Calculating mass ratios at aand {from the data in Table 3. and averaging. we obtain: and The ratio for the 4BLE- 4BIL svstem. measured from the map only. is The elongation of 4X could be explained as a circular clisk inclined at an anele of approximately 53° to the plane of the sky.," Calculating mass ratios at and from the data in Table \ref{peakmass} and averaging, we obtain; and The ratio for the 4BI - 4BII system, measured from the map only, is The elongation of 4A could be explained as a circular disk inclined at an angle of approximately $^{o}$ to the plane of the sky."
+ The elongation of 4D in the raw map appears to be EW., The elongation of 4B in the raw map appears to be EW.
+ However. upon deconvolving the beam. from the map we discovered that this elongation is due to 4B being itself à binary system.," However, upon deconvolving the beam from the map we discovered that this elongation is due to 4B being itself a binary system."
+ The primary component of 4B in the deconvolved. mmap appears to have a slight elongation along the same axis as 4A. Vhe amount of dilluse material in which the svsteni is embedded. is dillicult to determine with anv accuracy. since the map does not encompass the entire extended ridge. the flux density levels may be alfected by the chop throw sampling some of the extended emission. and at tthe error beam may. spread Lux density from the compact sources into the surrounding regions.," The primary component of 4B in the deconvolved map appears to have a slight elongation along the same axis as 4A. The amount of diffuse material in which the system is embedded is difficult to determine with any accuracy, since the map does not encompass the entire extended ridge, the flux density levels may be affected by the chop throw sampling some of the extended emission, and at the error beam may spread flux density from the compact sources into the surrounding regions."
+ The separation of the 4B double is approximately 127 on the sky., The separation of the 4B double is approximately 12” on the sky.
+ HE we assume that the triple svstem is coplanar. so that the 4BI-BLL axis has the same inclination to the plane of the sky as the 4X disk. the true DI-DIE separation would then be 167 (5600 AU for a distance of 350pc).," If we assume that the triple system is coplanar, so that the 4BI-BII axis has the same inclination to the plane of the sky as the 4A disk, the true BI-BII separation would then be 16” (5600 AU for a distance of 350pc)."
+ The 4A-4B separation is approximately 307., The 4A-4B separation is approximately 30”.
+ The long axis of the 4A clisk extends some S [rom the centre in the deconvolved images (we measure the disk extent as the cistance at which the [Lux density reaches of its peal: value)., The long axis of the 4A disk extends some 8” from the centre in the deconvolved images (we measure the disk extent as the distance at which the flux density reaches of its peak value).
+ Thus the presence of the 4A disk is not incompatible with the interpretation that 4A. 4Bl and 4BIL have coplanar orbits.," Thus the presence of the 4A disk is not incompatible with the interpretation that 4A, 4BI and 4BII have coplanar orbits."
+ We can also make an assessment of how stable the triple 4A-4BI-4BII system is., We can also make an assessment of how stable the triple 4A-4BI-4BII system is.
+ We first assume that the svsteni is coplanar and the orbits circular., We first assume that the system is coplanar and the orbits circular.
+" A criterion for stability in a hierarchical triple was developed by Harrington (1977). lere. Dis; is the periastron distance of the stanel-alone star. and £2,,,,5 the semi-major axis of the binary orbit."," A criterion for stability in a hierarchical triple was developed by Harrington (1977), Here, $D_{triple}$ is the periastron distance of the stand-alone star, and $D_{binary}$ the semi-major axis of the binary orbit."
+" Ads is the mass of the single component. A, ancl AJ» the masses of the binary components."," $M_3$ is the mass of the single component, $M_1$ and $M_2$ the masses of the binary components."
+ The constants have values A/—3.5 and 24=0. [or à corevolving svsten. or A=2.55. A=0.64 for a counterrotating svsten.," The constants have values $K=3.5$ and $A=0.7$ for a corevolving system, or $K=2.75$, $A=0.64$ for a counterrotating system."
+ Lor our case. the mass ratio is 1/0.63 (Section 6.2)).," For our case, the mass ratio is $1/0.63$ (Section \ref{ratios}) )."
+ The semimajor axis of the BI-BIL system cannot be smaller than the observed separation on the skv of 127., The semimajor axis of the BI-BII system cannot be smaller than the observed separation on the sky of 12”.
+ The true 4A-4DI separation of course depends on the angle of the binary-single star axis., The true 4A-4BI separation of course depends on the angle of the binary-single star axis.
+ We have alreacky argued (as have previous authors) that the 4A elongation is an inclined. disk., We have already argued (as have previous authors) that the 4A elongation is an inclined disk.
+ Lf we assume that the 4A-4B) orbit is coplanar with the 4A clisk (expected if fragmentation formation mocels apply). then the true 4A-4BI cistance would be 307. which would also be the largest possible periastron distance.," If we assume that the 4A-4BI orbit is coplanar with the 4A disk (expected if fragmentation formation models apply), then the true 4A-4BI distance would be 30”, which would also be the largest possible periastron distance."
+ Taking these figures we find that the system fails Harringtons stability test easily. even if we consider the more stable counterrotating case.," Taking these figures we find that the system fails Harrington's stability test easily, even if we consider the more stable counterrotating case."
+ See Figure 9.., See Figure \ref{stabfig}.
+ Lt is possible to envisage accretion of material stabilising initially-unstable triple systems by mocifsving the separation ratios of the components (Smith et al. 1997).," It is possible to envisage accretion of material stabilising initially-unstable triple systems by modifying the separation ratios of the components (Smith et al, 1997)."
+ There certainly seems to be a substantial reservoir of available material in the LRASS envelope (see Table. 4))., There certainly seems to be a substantial reservoir of available material in the IRAS4 envelope (see Table \ref{extmass}) ).
+ In the case of a low-mass binary orbiting a higher-mass single star. as here. this mechanism is unlikely to result. in. stability.," In the case of a low-mass binary orbiting a higher-mass single star, as here, this mechanism is unlikely to result in stability."
+ Low specific angular momentum material will tend to be accreted onto LRASLA. causing the single-binary separation to decrease ancl destabilizing the svstem.," Low specific angular momentum material will tend to be accreted onto IRAS4A, causing the single-binary separation to decrease and destabilizing the system."
+ High. specific angular momentum will tend to accrete onto the binary (1BI-4DID. widening it hy adding. angular momentuni and again destabilizing the system.," High specific angular momentum will tend to accrete onto the binary (4BI-4BII), widening it by adding angular momentum and again destabilizing the system."
+ Furthermore. the 4+A-4BI-4BIL system fails the Harrington. stability test by a comfortable margin.," Furthermore, the 4A-4BI-4BII system fails the Harrington stability test by a comfortable margin."
+ Even if the available envelope. mass were aecreted exclusively onto 4BI and 1211. moving the mass ratio towards L. there is not enough mass to entirely stabilize the system.," Even if the available envelope mass were accreted exclusively onto 4BI and BII, moving the mass ratio towards 1, there is not enough mass to entirely stabilize the system."
+ We conclude that LASS seems to be an unstable multiple which will be disrupted within a few orbits., We conclude that IRAS4 seems to be an unstable multiple which will be disrupted within a few orbits.
+ Could the elongation of the main cloud be explained. as a foreshortenccl disk. in a similar way to the elongation. of 4A?," Could the elongation of the main cloud be explained as a foreshortened disk, in a similar way to the elongation of 4A?"
+ The voung age (approx 107 vears deduced. from the, The young age (approx $^5$ years deduced from the
+is actually the nudi respousible for constraining iu teiiperature the Feh2TU peak value (sec. again. Fig. 1)).,"is actually the main responsible for constraining in temperature the Fe5270 peak value (see, again, Fig. \ref{fig:index_tio}) )."
+ We can probe iudex seuxifivitv also bv iueaus of enipirical data., We can probe index sensitivity also by means of empirical data.
+ Frou the Duzzouictal.(1991). fitting function we can write in fact: that fairly well coyres with the theoretical namely According to Wortlrevetal.(1991) inodoels we have while Gorgasetal.(1993) results Note from the equations the wez dudes seusitivitv to stellar surface graviti*, From the \citet{buzzoni94} fitting function we can write in fact: that fairly well compares with the theoretical namely According to \citet{wortheyetal94} models we have while \citet{gorgas93} results Note from the equations the weak index sensitivity to stellar surface gravity.
+ For example. anv chanee of one dex in loggouly coutributes at most by 0.3 tto. the Feh2T0Onas variation.," For example, any change of one dex in $\log g$only contributes at most by $\pm 0.3$ to the $_{\rm max}$ variation."
+ Tn terms of overall properties of a stellar agerceate. this figure reflects in a change of TO stellar mass of roughly 43 dex. which iuplies a variation of nearlv one order of magnitude iu stellar lifetiune.," In terms of overall properties of a stellar aggregate, this figure reflects in a change of TO stellar mass of roughly 0.3 dex, which implies a variation of nearly one order of magnitude in stellar lifetime."
+" Thereore. as a ΙΟΥ conclusion. we have that Fe5270,4. in a stellar populationaye."," Therefore, as a major conclusion, we have that $_{\rm max}$ in a stellar population."
+ A comparison of the different calibrations is shown in Fie. 2.," A comparison of the different calibrations is shown in Fig. \ref{f3},"
+ where we report the output iudex for two reference values of stellar eravity (uamcly. loggy=1 aud 3 dex) appropriate for red giants;," where we report the output index for two reference values of stellar gravity (namely, $\log g = 1$ and 3 dex) appropriate for red giants."
+ Within a remarkably eood agreement among the differcut datasets. compared to Buzzouietal.(1991). calibration. one notes from the plot that Worthevetal.(1991). fit displavs an even lower dependence on logg. while Gorgasetal.(1993) ft stands out for its steeper trend with metallicity. especially at super-solar regines.," Within a remarkably good agreement among the different datasets, compared to \citet{buzzoni94} calibration, one notes from the plot that \citet{wortheyetal94} fit displays an even lower dependence on $\log g$, while \citet{gorgas93} fit stands out for its steeper trend with metallicity, especially at super-solar regimes."
+ Overall we cau conclude that thedirect Fe5270 dependence on Fe abundance turus out to be AFc5270=onΛΙΓΟΠΠ. with o—1.2>LI ," Overall, we can conclude that the Fe5270 dependence on Fe abundance turns out to be $\Delta {\rm Fe5270} =
+\alpha \Delta {\rm [Fe/H]}$, with $\alpha = 1.2\rightarrow 1.4$."
+Iu addition to its direct scusitivity to |Fo/II]. a major supplementary advantage of using F05270 feature as metallicity tracer in stellay populations resides iu the special property of the index to always peak the temperature range sampled by red elauts in stellar svsteus richer than [Fe/U] >2 Conversely. this is not the case. for stance. of other stronger features. used to derive popular iudices like he Lick Mg» or the near-IR CaT triplet index (Jones 2001). ," In addition to its direct sensitivity to [Fe/H], a major supplementary advantage of using Fe5270 feature as metallicity tracer in stellar populations resides in the special property of the index to always peak the temperature range sampled by red giants in stellar systems richer than [Fe/H] $\gtrsim -2$ Conversely, this is not the case, for instance, of other stronger features, used to derive popular indices like the Lick $_2$ or the near-IR CaT triplet index \citep{jones84,idiart97,cenarro01}. ."
+Both these features peak. in fact. among elants and dwarts at the coolest temperature tail of cach οΠΟ».," Both these features peak, in fact, among giants and dwarfs at the coolest temperature tail of each CMDs,"
+and the induction equation.,and the induction equation.
+ These equations take conserved Lyberbolic forms. for easy iutegration: where U. F. and ο represent the set of conserved variables. fiuxes. aud sources. respectively.," These equations take conserved hyberbolic forms, for easy integration: where U, F, and S represent the set of conserved variables, fluxes, and sources, respectively."
+ ILARM is set to first seed a torus around a black hole with au initial deusity aud poloidal maguetic field. then evolve the space through time. tracking deusity. iuterual energv. L-velocity. Lanaguctic field.," HARM is set to first seed a torus around a black hole with an initial density and poloidal magnetic field, then evolve the space through time, tracking density, internal energy, 4-velocity, 4-magnetic field."
+" This is allowed to continue until the space reaches a ""steady state” - of special note in the code output is the liehhy turbulent naenetic field. au artifact of the magneto-rotational instability. which is an inportaut mechanisin goveruiug he transport of anenlarC» momentum in the disk."," This is allowed to continue until the space reaches a “steady state” - of special note in the code output is the highly turbulent magnetic field, an artifact of the magneto-rotational instability, which is an important mechanism governing the transport of angular momentum in the disk."
+ This instability causes the turbulence which is likely to be at Cst partlv responsible for the energizing of electrons in the disk. though in our MC trials. we only consider lly thermal distributions.," This instability causes the turbulence which is likely to be at least partly responsible for the energizing of electrons in the disk, though in our MC trials, we only consider fully thermal distributions."
+" The MC code used for photon emission aud scatteriug has a loug history and is discussed i a nmuuber of resources, such as Caufieldetal.(1987):Liang&Der-imer(1988):Bottcheretal. (1998)."," The MC code used for photon emission and scattering has a long history and is discussed in a number of resources, such as \cite{can87,lia88,bot98}."
+. For a complete treatment. readers should see these papers.," For a complete treatment, readers should see these papers."
+ Iu general. this code is a coupled MC/FP (FokkerPlauck) code.," In general, this code is a coupled MC/FP (Fokker-Planck) code."
+ For our inteuts. the FP evolution of the electron distribution was unnuecessarv at this stage. so it was turned off to allow a fixed teniperature giveu by the TARAL output.," For our intents, the FP evolution of the electron distribution was unnecessary at this stage, so it was turned off to allow a fixed temperature given by the HARM output."
+ The code is set up ou a 2D axialli-απλο] cevliudrcal exid. creating a (hollow or solid. depending ou whether the inner radius is set to zero) cvlndzrical shape.," The code is set up on a 2D axially-symmetric cylindrical grid, creating a (hollow or solid, depending on whether the inner radius is set to zero) cylindrical shape."
+ Each zone is assigued a deusitv. ion and electron temperatures. magnetic field. amplitude. and thermal and uonthermal distribution compoucuts.," Each zone is assigned a density, ion and electron temperatures, magnetic field amplitude, and thermal and nonthermal distribution components."
+ For our purposes. this is simply set to be a Maxwellianu. but the code allows power law noutheriual distributions as well.," For our purposes, this is simply set to be a Maxwellian, but the code allows power law nonthermal distributions as well."
+ The code allows cussion from the volume and boundaries. aud emitted photons are tracked aud allowed. to scatter or absorb.," The code allows emission from the volume and boundaries, and emitted photons are tracked and allowed to scatter or absorb."
+ Output roni the ITAR\ code ds used to assigu he zone quantities., Output from the HARM code is used to assign the zone quantities.
+ That is. we specify a last density aud the rest of the output scales to give us he ion temperature aud saturated MBRI iiagnuetic field.," That is, we specify a maximum density and the rest of the output scales to give us the ion temperature and saturated MRI magnetic field."
+ Asstuning that inclusion of an initial toroidal maguetic fields would add to the final field. as this dimension is niostlv unaffected. by evolution. allows us to set he magnetic field. as long as it doesut drop below he saturation value.," Assuming that inclusion of an initial toroidal magnetic fields would add to the final field, as this dimension is mostly unaffected by evolution, allows us to set the magnetic field, as long as it doesn't drop below the saturation value."
+" Similarly, as we have a specified ki temperature but no wav to directly evaluate the electrons acceleration. we set a global ratio between the two values."," Similarly, as we have a specified ion temperature but no way to directly evaluate the electrons' acceleration, we set a global ratio between the two values."
+ That is. in one trial the electron temperature in a zone could be set to always be twice the value of the ion teniperature in the same zone.," That is, in one trial the electron temperature in a zone could be set to always be twice the value of the ion temperature in the same zone."
+ For all of the fits presented in figures below. the open circles are data poiuts. triangles are upper limits. aud the bowties denote the Chanudra-obtained flaring aud quiescent x-ray data points aud slopes.," For all of the fits presented in figures below, the open circles are data points, triangles are upper limits, and the bowties denote the Chandra-obtained flaring and quiescent x-ray data points and slopes."
+ The first three figures show fits to the flaring x-ray point., The first three figures show fits to the flaring x-ray point.
+ As shown. the flaring bowtie las almost zero slope associated with it.," As shown, the flaring bowtie has almost zero slope associated with it."
+ Iu order to match this. there were several options;," In order to match this, there were several options."
+ Obviously. the bremisstraliluug conrponeut is very flat out to the high energv cut-off point.," Obviously, the bremsstrahlung component is very flat out to the high energy cut-off point."
+" It is also possible for Conpton-scattered components to reach these energies and have flat spectra - trials are shown for fits usine firs the first-scattered: bump. and secoudhy, the bump arising frou Notons that scatter twice off of hot electrons."," It is also possible for Compton-scattered components to reach these energies and have flat spectra - trials are shown for fits using first the first-scattered bump, and secondly, the bump arising from photons that scatter twice off of hot electrons."
+ These three spectral components provide nearly equally-adequate fits. if only considering the flaring wchavior. though the second Compton trial in Figure 3 seclus to best fit the lower energy compoucuts.," These three spectral components provide nearly equally-adequate fits, if only considering the flaring behavior, though the second Compton trial in Figure 3 seems to best fit the lower energy components."
+ It should © noted that effort was not focused on fittine the very low energy radio points., It should be noted that effort was not focused on fitting the very low energy radio points.
+ It is expected hat these arise primarily frou a very large volume of ow density. temperature. aud maenetic field around he aceretion disc. which is bevoud the scale of this experiment. as our volume cuts off at r = LOX. Ouce the faring trials were completed. consideration had to be eiven to fittine the quiescent points with a consistent inethod. given the flarine fts.," It is expected that these arise primarily from a very large volume of low density, temperature, and magnetic field around the accretion disc, which is beyond the scale of this experiment, as our volume cuts off at r = 40M. Once the flaring trials were completed, consideration had to be given to fitting the quiescent points with a consistent method, given the flaring fits."
+ It was found that very good fits to the quicscent poiut were possible by simply dropping the density or teniperature independeutlv. for the two Compton buup trials.," It was found that very good fits to the quiescent point were possible by simply dropping the density or temperature independently, for the two Compton bump trials."
+ Unfortunately. the breimsstraliluug trial did not fit the quiescent point regardless of how its parameters were uodified. as it will always have a very flat spectrum at he x-ray point.," Unfortunately, the bremsstrahlung trial did not fit the quiescent point regardless of how its parameters were modified, as it will always have a very flat spectrum at the x-ray point."
+ The second Compton buiup quicscent fit was achieved * chaugime only the density in the trial., The second Compton bump quiescent fit was achieved by changing only the density in the trial.
+" This is representativo of a global mass accretion rate change. which could certainly be possible for observed variability ron Ser A*,"," This is representative of a global mass accretion rate change, which could certainly be possible for observed variability from Sgr A*."
+ Conversely. it was found that for the first Compton mn trial. dropping the temperature allowed for a eood ft.," Conversely, it was found that for the first Compton bump trial, dropping the temperature allowed for a good fit."
+" This could indicate that the faring state arises due to some larec-scale magnetic energv dissipation event. such as rapid reconnectionu over a laree area,"," This could indicate that the flaring state arises due to some large-scale magnetic energy dissipation event, such as rapid reconnection over a large area."
+"As seen in the previous sections, the peak in the circular velocity profile of a subhalo is a more stable quantity to recover than the total subhalo mass.","As seen in the previous sections, the peak in the circular velocity profile of a subhalo is a more stable quantity to recover than the total subhalo mass."
+ The origin of this stability is related to the fact that the radius at which the maximum circular velocity is reached is located much closer to the centre of the halo and so is unaffected by truncation., The origin of this stability is related to the fact that the radius at which the maximum circular velocity is reached is located much closer to the centre of the halo and so is unaffected by truncation.
+ Fig., Fig.
+ 6 shows how the position of peak changes with the concentration of a halo., \ref{conc} shows how the position of peak changes with the concentration of a halo.
+" For a NFW halo this can be obtained numerically to give, The values determined by equation (14)) are based on an ideal NFW halo, but for low resolution haloes there will be deviations from this curve."," For a NFW halo this can be obtained numerically to give, The values determined by equation \ref{vmax}) ) are based on an ideal NFW halo, but for low resolution haloes there will be deviations from this curve."
+" For the subhalo used in this work (c= 12) rvmax=0.18ryi; which corresponds to roughly r5ooo (the radius at which the enclosed density is 5,000 times the critical density, pci)."," For the subhalo used in this work $c=12$ ) $r_{\rm vmax}=0.18r_{\rm vir}$ which corresponds to roughly $r_{5000}$ (the radius at which the enclosed density is 5,000 times the critical density, $\rho_{\rm crit}$ )."
+" Stripping occurs in the outer regions of the subhalo and so for it to affect this radius a large amount of material needs to be lost, consistent with Fig. 5.."," Stripping occurs in the outer regions of the subhalo and so for it to affect this radius a large amount of material needs to be lost, consistent with Fig. \ref{deflection}."
+ One of the main issues with using the maximum circular velocity of a halo is how its measurement depends upon resolution., One of the main issues with using the maximum circular velocity of a halo is how its measurement depends upon resolution.
+" To investigate this, we generated a halo with My,=1012Μ9 and c=12 in isolation using a different number of particles within the virial radius each time."," To investigate this, we generated a halo with $M_{\rm vir}=10^{12} \rm M_{\odot}$ and $c=12$ in isolation using a different number of particles within the virial radius each time."
+" For each number of particles within the virial radius, we constructed 1,000 realisations in order to constrain the variation."," For each number of particles within the virial radius, we constructed 1,000 realisations in order to constrain the variation."
+ Fig., Fig.
+ 7 shows how the recovered maximum circular velocity varied with the total particle number., \ref{res} shows how the recovered maximum circular velocity varied with the total particle number.
+ For the sparsely populated realisations the average maximum circular velocity was higher than the analytic value., For the sparsely populated realisations the average maximum circular velocity was higher than the analytic value.
+" As more particles were used, the two values converged."," As more particles were used, the two values converged."
+" For the average value to be within 2.5 per cent of the analytic value, in excess of 500 particles were required in the halo."," For the average value to be within 2.5 per cent of the analytic value, in excess of 500 particles were required in the halo."
+ The variation of the maximum circular velocity between different realisations of the same total virial particle number is strong for the sparsely populated haloes., The variation of the maximum circular velocity between different realisations of the same total virial particle number is strong for the sparsely populated haloes.
+" At all points the curve is within 1 standard deviation of the analytic value, but the variation is clear where for 10 particles the standard deviation is 0.56 compared with 0.002 for 10,000."," At all points the curve is within 1 standard deviation of the analytic value, but the variation is clear where for 10 particles the standard deviation is 0.56 compared with 0.002 for 10,000."
+" To obtain an accurate value for the maximum circular velocity of a recovered subhalo, its resolution has to be taken into account."," To obtain an accurate value for the maximum circular velocity of a recovered subhalo, its resolution has to be taken into account."
+ Halo finders are an important tool for the analysis of cosmological simulations., Halo finders are an important tool for the analysis of cosmological simulations.
+" They are pivotal in the construction of merger trees, which underpin galaxy formation modelling, and their results allow us to characterise, for example, the abundance and spatial distribution of both dark matter haloes and subhaloes."," They are pivotal in the construction of merger trees, which underpin galaxy formation modelling, and their results allow us to characterise, for example, the abundance and spatial distribution of both dark matter haloes and subhaloes."
+ There are as many techniques for identifying haloes and subhaloes in cosmological simulations as there are halo, There are as many techniques for identifying haloes and subhaloes in cosmological simulations as there are halo
+Fig.,Fig.
+ 5. displays the spatial ACT of the ds for the clouds of sample HR. The figure clearly shows that there is no substantial correlation between the orientations of the PAs for the clouds of sample It 0n any physical scale., \ref{fig5} displays the spatial ACF of the $PA$ s for the clouds of sample R. The figure clearly shows that there is no substantial correlation between the orientations of the $PA$ s for the clouds of sample R on any physical scale.
+ On he other hand. Fig.," On the other hand, Fig."
+ 6 displays the spatial correlations of the PAs for the Iam sample., \ref{fig6} displays the spatial correlations of the $PA$ s for the Rm sample.
+ In contrast to the un-mirrored case. he mirrored P.4s for the selected clouds show non-negligible correlations on spatial scales of the order of 100800 pc.," In contrast to the un-mirrored case, the mirrored $PA$ s for the selected clouds show non-negligible correlations on spatial scales of the order of $\sim 100-800$ pc."
+" We have also checked for the consisteney of the results by aking different permutations of /,,5, and fu. and of their imposed. ratio: Le. values of (ini uus (pe))2(2.5). (3.5). (3.7) with and without imposing that frasefain and found no significant variations in the results for Fig."," We have also checked for the consistency of the results by taking different permutations of $l_{min}$ and $l_{max}$ and of their imposed ratio; i.e., values of $l_{min}$ $l_{max}$ (pc))=(2,5), (3,5), (3,7) with and without imposing that $l_{max}/l_{min}\ge 2$, and found no significant variations in the results for Fig."
+ and Fig. 6., \ref{fig5} and Fig. \ref{fig6}.
+ Phese spatial scales of a few hundred: parsecs ave well matched by the sizes of supernova remnants in low clensity environments., These spatial scales of a few hundred parsecs are well matched by the sizes of supernova remnants in low density environments.
+ Indeed. using analytical calculations. AleCray Ixafatos (1987) estimated the cooling radius of a supernova shell/supershell to be given by: where & is the local value of the metallicity. Wy is the number of stars in the stellar cluster with masses greater than 7 AL... Ez is the energy input of one supernova in units of 10 erg. and ny. in 7. is the ambient gas density.," Indeed, using analytical calculations, McCray Kafatos (1987) estimated the cooling radius of a supernova shell/supershell to be given by: where $\zeta$ is the local value of the metallicity, $N_{\star}$ is the number of stars in the stellar cluster with masses greater than $7$ $_{\odot}$, $E_{51}$ is the energy input of one supernova in units of $10^{51}$ erg, and $n_{0}$, in $^{-3}$, is the ambient gas density."
+ The cooling radius marks the end of the adiabatic phase of the expanding shell/supershell and. the time at which cooling becomes important in the hot interior of the remnant., The cooling radius marks the end of the adiabatic phase of the expanding shell/supershell and the time at which cooling becomes important in the hot interior of the remnant.
+" Note that the shell/supershell continues to expand: afterwards according to the zero-pressure snowplow law with its raclius slowly increasing following RU)=(ΕΙ ση. where £, ds the time at which #2. has been reached."," Note that the shell/supershell continues to expand afterwards according to the zero-pressure snowplow law with its radius slowly increasing following $R(t)=R_{c} (t/t_{c})^{1/4}$ , where $t_{c}$ is the time at which $R_{c}$ has been reached."
+" For an ambient gas density ag=0.1 57. IN,1. and a solar moetallicitv. the value of A2. is of the order of ~200 pc. for N,=1. ¢=1. and ny=0.03 em7. A410 pe. whereas for ng=0.03 . N,=1. and a lower metallicity value of &=0.5. 16 cooling radius is A.~765 pc."," For an ambient gas density $n_{0} =0.1$ $^{-3}$, $N_{\star}=1$, and a solar metallicity, the value of $R_{c}$ is of the order of $\sim 200$ pc, for $N_{\star}=1$, $\zeta=1$, and $n_{0}=0.03$ $^{-3}$, $R_{c} \sim 410$ pc, whereas for $n_{0}=0.03$ $^{-3}$, $_{\star}=1$, and a lower metallicity value of $\zeta=0.5$, the cooling radius is $R_{c} \sim 765$ pc."
+ In order to compare the 'orrelations to the noise. Fig.," In order to compare the correlations to the noise, Fig."
+ 6 also displays (dashed line) 16 ACT. for the same sample of clouds but with their Pls rawn from random values (averaged over 10 iterations)., \ref{fig6} also displays (dashed line) the ACF for the same sample of clouds but with their $PA$ s drawn from random values (averaged over 10 iterations).
+ As *xpected in this case. the ACT tends to zero values on all μαxtial scales.," As expected in this case, the ACF tends to zero values on all spatial scales."
+ The similarity between the latter physical scales. and 10 physical correlation scales observed in Fig., The similarity between the latter physical scales and the physical correlation scales observed in Fig.
+ 6 suggests ju supernova driving might be plaving a considerable role in shaping the ISAT structure ancl dynamics even in regions of low star formation activity., \ref{fig6} suggests that supernova driving might be playing a considerable role in shaping the ISM structure and dynamics even in regions of low star formation activity.
+ “Vhis is consistent with the results of Dib ct al. (, This is consistent with the results of Dib et al. (
+2006) who showed. that supernova: [feedback is capable of maintaining turbulent velocity dispersions for the LIL gas of the order of 3.5 km  even in regions where the star formation rate per unit. area Is. as low as 10.71 M. vr71 57. corresponding: to one hundredth of the Galactic SN rate.,"2006) who showed that supernova feedback is capable of maintaining turbulent velocity dispersions for the HI gas of the order of $3-5$ km $^{-1}$ even in regions where the star formation rate per unit area is as low as $10^{-4}$ $_{\odot}$ $^{-1}$ $^{-2}$, corresponding to one hundredth of the Galactic SN rate."
+ The existence of spatial correlations for the sample of mirrored Ps which we associate to SN shells ancl the absence of such correlations for sample It indicates that there are no correlations between the sites of star formation and that gravitational instability occurs in the Galaxy wherever local concitions are favorable., The existence of spatial correlations for the sample of mirrored $PA$ s which we associate to SN shells and the absence of such correlations for sample R indicates that there are no correlations between the sites of star formation and that gravitational instability occurs in the Galaxy wherever local conditions are favorable.
+ Note that. in calculating the AC's. we have mace use of the kinematical distances to each of the clouds published by Llever ct al. (," Note that, in calculating the ACFs, we have made use of the kinematical distances to each of the clouds published by Heyer et al. ("
+2001).,2001).
+ Phe kinematic distance in the outer Galaxy is an overestimate (e.g. Drand Blitz 1993) of the true distance measured by trigonometric parallax. measurements (c.g.. Llachisuka ct al.," The kinematic distance in the outer Galaxy is an overestimate (e.g., Brand Blitz 1993) of the true distance measured by trigonometric parallax measurements (e.g., Hachisuka et al."
+ 2006.2009).," 2006,2009)."
+ For example. more accurate distance estimates using parallax measurements for the W3 region vield a distance of 2.1 kpe (Hachisuka et al.," For example, more accurate distance estimates using parallax measurements for the W3 region yield a distance of $2.1$ kpc (Hachisuka et al."
+ 2006) whieh is smaller than the value of ~3.5 kpe we have used for clouds in that region., 2006) which is smaller than the value of $\sim 3.5$ kpc we have used for clouds in that region.
+ Vhe overestimate of the distance by the kinematic distance in the outer Galaxy will depend on the position of the cloud and on the existence of random velocity components and/or velocity. olfsets at. this location with respect to a given Galactic rotation curve., The overestimate of the distance by the kinematic distance in the outer Galaxy will depend on the position of the cloud and on the existence of random velocity components and/or velocity offsets at this location with respect to a given Galactic rotation curve.
+ We have calculated simple models to quantify the kinematic distance uncertainties when using the [lat rotation curve (with a circular velocity of ej).=220 km 5 similar to the value adopted in Llever et al., We have calculated simple models to quantify the kinematic distance uncertainties when using the flat rotation curve (with a circular velocity of $v_{circ}=220$ km $^{-1}$ similar to the value adopted in Heyer et al.
+ 2001) for objects within the longitude range studied in this paper., 2001) for objects within the longitude range studied in this paper.
+ The method. proceeds as follow: We first calculate a distance assuming pure circular motions for a range of expectet velocities and longitudes in the ranges -5.-100] kms+ ar “100°. 1407]. respectively.," The method proceeds as follow: We first calculate a distance assuming pure circular motions for a range of expected velocities and longitudes in the ranges [-5,-100] km $^{-1}$ and $100^{\circ},140^{\circ}$ ], respectively."
+ In aà second step. we add a velocity ollset (Ae) and/or à random component with a dispersion of σι and recompute the distance.," In a second step, we add a velocity offset $\Delta v$ ) and/or a random component with a dispersion of $\sigma_{v}$ and recompute the distance."
+ The procedure is repeatec 8192 times in order to sample the random component of each longitude and velocity., The procedure is repeated 8192 times in order to sample the random component of each longitude and velocity.
+ With that. we can derive the roo mean square. ap. of the newly calculated distance.," With that, we can derive the root mean square, $\sigma_{D}$, of the newly calculated distance."
+" Fig. shows the fractional clistance error 05/42 for three sets of 0, ancl A, as a Function of the velocity for various values of the longitude /.", \ref{fig7} shows the fractional distance error $\sigma_{D}/D$ for three sets of $\sigma_{v}$ and $\Delta_{v}$ as a function of the velocity for various values of the longitude $l$.
+" Lf there are no streaming motions (A,= 0) but only a ranclom velocity component of dispersion m.=5 km 5 (left plot in Fig.7)). the dilferenees between the kinematic distance and the true distance are of the order of =10 percent formost velocities and get larger for nearby clouds."," If there are no streaming motions $\Delta_{v}=0$ ) but only a random velocity component of dispersion $\sigma_{v}=5$ km $^{-1}$ (left plot in \ref{fig7}) ), the differences between the kinematic distance and the true distance are of the order of $\lesssim 10$ percent formost velocities and get larger for nearby clouds."
+ Such errors could be propagated into the calculation of the inter-cloucl distances using Eq.2 on a two-by-two basis., Such errors could be propagated into the calculation of the inter-cloud distances using \ref{eq2} on a two-by-two basis.
+ However. the overall uncertainty on the ACF is not expected. to be very large ancl might stretch. for a given," However, the overall uncertainty on the ACF is not expected to be very large and might stretch, for a given"
+the soft/high X-ray states.,the soft/high X-ray states.
+" The radio emission during the latter states can become extremely low, but also can flare during the soft state back to the hard-state level, e.g., during the long soft state of MJD 52160-52570."," The radio emission during the latter states can become extremely low, but also can flare during the soft state back to the hard-state level, e.g., during the long soft state of MJD 52160–52570."
+" 11 in ZPS11 shows also the 15-50 keV BAT light curve, normalized in the same way."," 1 in ZPS11 shows also the 15–50 keV BAT light curve, normalized in the same way."
+ We quantify the radio/X-ray correlation in Section 3.3 below., We quantify the radio/X-ray correlation in Section \ref{corr} below.
+" During the 15 GHz monitoring of Cyg X-1, there were four events when the flux rose above 50 mJy."," During the 15 GHz monitoring of Cyg X-1, there were four events when the flux rose above 50 mJy."
+ We show their light curves in detail in Appendix A.., We show their light curves in detail in Appendix \ref{flares}. .
+" Appendix B gives details of our method of calculating the bolometric flux (i.e., that integrated over all energies) based on the ASM, BATSE and BAT data. reff:gamma(("," Appendix \ref{bol} gives details of our method of calculating the bolometric flux (i.e., that integrated over all energies) based on the ASM, BATSE and BAT data. \\ref{f:gamma}( ("
+a) shows the relationship between the resulting bolometric flux and I'(3-12 keV).,a) shows the relationship between the resulting bolometric flux and $\Gamma(3$ –12 keV).
+" In order to allow comparison with commonly plotted hardness-flux diagrams, the horizontal axis shows Γ decreasing to the right."," In order to allow comparison with commonly plotted hardness-flux diagrams, the horizontal axis shows $\Gamma$ decreasing to the right."
+" We see a positive correlation of Fog with Γ in the intermediate state, with the Spearman's rank-order correlation coefficient, r,50.16, and the corresponding probability that this r, is due to a chance, P,~1x10:14."," We see a positive correlation of $F_{\rm bol}$ with $\Gamma$ in the intermediate state, with the Spearman's rank-order correlation coefficient, $r_{\rm s}\simeq 0.16$, and the corresponding probability that this $r_{\rm s}$ is due to a chance, $P_{\rm s}\simeq 1\times 10^{-14}$."
+" The correlation continues to the hard state for Fy.>30 keV cm~ s!, within which r,~0.52, P,~2x107""."," The correlation continues to the hard state for $F_{\rm bol}>30$ keV $^{-2}$ $^{-1}$, within which $r_{\rm s}\simeq 0.52$, $P_{\rm s}\simeq 2\times 10^{-17}$."
+" In the hard state at lower Fyg, it changes by a factor of a few at a given Γ without any apparent trend."," In the hard state at lower $F_{\rm bol}$, it changes by a factor of a few at a given $\Gamma$ without any apparent trend."
+" The soft state data, at (3-12 keV)22.3, show no clear overall correlation."," The soft state data, at $\Gamma(3$ –12 $)\ga 2.3$, show no clear overall correlation."
+" In this state, I(3-12 keV) is not uniquely correlated with the bolometric flux."," In this state, $\Gamma(3$ –12 keV) is not uniquely correlated with the bolometric flux."
+ We note that in a number of occurrences of the soft state Fyg4 was lower than that in the high- hard state., We note that in a number of occurrences of the soft state $F_{\rm bol}$ was lower than that in the high-flux hard state.
+" This happened mostly during the 2010/11 soft state, when the ASM fluxes reached relatively low values, see 4((b)."," This happened mostly during the 2010/11 soft state, when the ASM fluxes reached relatively low values, see \\ref{f:lc}( (b), \ref{f:lc_bol}( (b)."
+ , \\ref{f:lc_bol} presents the light curve of $F_{\rm bol}$.
+We see changes up to a factor of ~10., We see changes up to a factor of $\sim$ 10.
+" We clearly see the superorbital periodicities of Cyg X-1, ~150 d in reff:lc,0l((a)and~300 d in reff:lc,ol((b), see, e.g., Poutanen,Zdziarski&Ibragimov(2008),,ZPS 11. -"," We clearly see the superorbital periodicities of Cyg X-1, $\sim$ 150 d in \\ref{f:lc_bol}( (a) and $\sim$ 300 d in \\ref{f:lc_bol}( (b), see, e.g., \citet*{pzi08}, ZPS11. \\ref{f:X_radio}( ("
+" —ejllustratetheX rayspectralvariability patternso f CygX Ibyshowingtherelationshipbetweenthebolometric fluxandthe f luxinagivenband, Fx.","a–e) illustrate the X-ray spectral variability patterns of Cyg X-1 by showing the relationship between the bolometric flux and the flux in a given band, $F_{\rm X}$."
+" We have fitted these relationships in the hard state (red in reff:lc,ol— —6))byapowerlaw, FyοςFe (see Appendix 3))."," We have fitted these relationships in the hard state (red in \\ref{f:lc_bol}- \ref{f:hot_bol}) ) by a power law, $F_{\rm bol}\propto F_{\rm X}^{p'}$ (see Appendix \ref{fits}) )."
+ The indices found are given in Table 1.., The indices found are given in Table \ref{t:fit}.
+ We see that p’~1 in the 14-150 keV range., We see that $p'\simeq 1$ in the 14–150 keV range.
+ This corresponds to a spectrum of constant shape moving up and down in normalization only., This corresponds to a spectrum of constant shape moving up and down in normalization only.
+" On the other hand, the dependencies are substantially weaker than direct proportionality in the 1.5—3 keV and 150-195 keV bands."," On the other hand, the dependencies are substantially weaker than direct proportionality in the 1.5–3 keV and 150–195 keV bands."
+" The intermediate state (green in reff:lc,ol———6)) show sadif erentvariability pattern.", The intermediate state (green in \\ref{f:lc_bol}- \ref{f:hot_bol}) ) shows a different variability pattern.
+"U ptol2keV, theintermediate stat "," Up to 12 keV, the intermediate state points are quite similar to those in the hard state, though they extend to higher fluxes."
+"keV and it becomes increasingly negative (i.e, the local X-ray flux is anti-correlated with Fy)), as shown in reff:X,adio((e) for75— —100keV. The soft state (blue in reff:lc,ol— —6))showsanapproximatedependenceof Fy€Fx over all the observed energy range."," However, the dependence changes sign at $E\ga 20$ keV and it becomes increasingly negative (i.e, the local X-ray flux is anti-correlated with $F_{\rm bol}$ ), as shown in \\ref{f:X_radio}( (e) for 75–100 keV. The soft state (blue in \\ref{f:lc_bol}- \ref{f:hot_bol}) ) shows an approximate dependence of $F_{\rm bol}\simprop F_{\rm X}$ over all the observed energy range."
+" However, the scatter increases with the increasing energy."," However, the scatter increases with the increasing energy."
+" For coronal models (e.g., Merloni&Fabian 2003)), the flux, F,.,, due to emission of coronal hot electrons is an important quantity."," For coronal models (e.g., \citealt{mf02,merloni03}) ), the flux, $F_{\rm hot}$, due to emission of coronal hot electrons is an important quantity."
+" It is predicted to be constant for coronae above gas-pressure dominated discs, and to decrease as ΛΜ1/2 for radiation-pressure dominated discs, see a discussion in Merlonial. (2003)."," It is predicted to be constant for coronae above gas-pressure dominated discs, and to decrease as $\dot M^{-1/2}$ for radiation-pressure dominated discs, see a discussion in \citet{mhd03}."
+. It is equal to the bolometric X-ray flux minus the disc blackbody contribution., It is equal to the bolometric X-ray flux minus the disc blackbody contribution.
+" Ideally, to calculate Fi; accurately one should fit each spectrum with a model containing a disc blackbody, and then subtract its total flux from the resulting model bolometric flux."," Ideally, to calculate $F_{\rm hot}$ accurately one should fit each spectrum with a model containing a disc blackbody, and then subtract its total flux from the resulting model bolometric flux."
+" Given our method of determining Fs, this is not possible."," Given our method of determining $F_{\rm bol}$, this is not possible."
+" Instead, we use the flux above an energy approximately dividing the spectral region dominated by the blackbody and that dominated by Comptonization."," Instead, we use the flux above an energy approximately dividing the spectral region dominated by the blackbody and that dominated by Comptonization."
+" Specifically, we assume Fi,=F(>3keV) in the soft state, and =F(>1.5keV) in the intermediateand hard states."," Specifically, we assume $F_{\rm hot}= F(>\!\!3\,{\rm keV})$ in the soft state, and $=F(>\!\!1.5\,{\rm keV})$ in the intermediateand hard states."
+ These choices are based on the Cyg X-1 spectra shown in Zdziarski&Gierlifiski (2004).., These choices are based on the Cyg X-1 spectra shown in \citet{zg04}. .
+" reff:hot,olshowsthere sultinge stimateo f thecoronalemissionasa f ractionof Ἐνοι.", \\ref{f:hot_bol} shows the resulting estimate of the coronal emission as a fraction of $F_{\rm bol}$ .
+ We see it is approximately constant in the soft state., We see it is approximately constant in the soft state.
+All of these data have been released in fully processed. form.,All of these data have been released in fully processed form.
+ No additional processing is necessary., No additional processing is necessary.
+ However. the images in each data set have clifferent scales aud sizes than the images in other clata sets.," However, the images in each data set have different scales and sizes than the images in other data sets."
+ There are two drivers to the photometry: in order to compute photometric redshifts. it is necessary to have accurate colors.," There are two drivers to the photometry: in order to compute photometric redshifts, it is necessary to have accurate colors."
+ For this. fixed circular aperture photometry. preferably hrough a small apertire. is sulficien.," For this, fixed circular aperture photometry, preferably through a small aperture, is sufficient."
+ The photometry is also used for computing masses of the &alaxies: for this. accurate total magütudes are necessary.," The photometry is also used for computing masses of the galaxies; for this, accurate total magnitudes are necessary."
+ Total magnitudes are best measured hrough an adaptive eliptical apertire (τοι1980).., Total magnitudes are best measured through an adaptive elliptical aperture \citep{kron}.
+ If there are to be uo systemaic shifts. then these apertures nuts be matchec ---1 all bands.," If there are to be no systematic shifts, then these apertures must be matched in all bands."
+ If the images are registered. this cau be doue easiy Using SExtractor (Bertin&Aruous1996 in “couble-image moc”.," If the images are registered, this can be done easily using SExtractor \citep{hihi} in “double-image mode”."
+ Since all the images lave «illerent scales. {μον tuist be resam21661 to pi them on the same astrometric grid.," Since all the images have different scales, they must be resampled to put them on the same astrometric grid."
+ The resaiipling was «[9]je. witlSWarp (Bertin2001)., The resampling was done with \citep{swarp}.
+. SWarp’s main prpose is to «‘ombine several (not necessarily overlappiug ünages Into a single image., 's main purpose is to combine several (not necessarily overlapping) images into a single image.
+ It. resamples the input images to put them on a common asrometric¢ erid and scales the flux to correct lor avy photometric shifts between the images., It resamples the input images to put them on a common astrometric grid and scales the flux to correct for any photometric shifts between the images.
+ It can also be isel (ο just remap a single image with uo flux scaling., It can also be used to just remap a single image with no flux scaling.
+ The NICMOS A image was used as the base image to which all the other iuages were matched., The NICMOS $H$ image was used as the base image to which all the other images were matched.
+ The NICMOS J image needed no remappiug. as it was already ou the same scae.," The NICMOS $J$ image needed no remapping, as it was already on the same scale."
+ The optical images from the UDF were remapped., The optical images from the UDF were remapped.
+" This changed their scale from 0.03 ""/pixel to 0.09 ""/pixel. making them slightly undersampled."," This changed their scale from 0.03 $''$ /pixel to 0.09 $''$ /pixel, making them slightly undersampled."
+ The ISAAC images were SWarped together so that cillerent pointings in eac1 band were combined into a single mosaic., The ISAAC images were ed together so that different pointings in each band were combined into a single mosaic.
+ The EIS images show a simall but systematic shit, The EIS images show a small but systematic shift
+MOOG.,MOOG.
+ Finally. we thank our anonymous referee for his/her careful reading of our manuscript aud helpful comments.," Finally, we thank our anonymous referee for his/her careful reading of our manuscript and helpful comments."
+ In addition. R. M. C. acknowledges the Natioual Research Council lor support duriug a part of this work. as well as the AAS Small Research Crant program. and the Djehuty Stellar Evolution Project.," In addition, R. M. C. acknowledges the National Research Council for support during a part of this work, as well as the AAS Small Research Grant program, and the Djehuty Stellar Evolution Project."
+ This work was performed uxler the auspices of the U.S. Departinent of Energy. National Nuclear Security Acininistration by the University of California. Lawrence Livermore National Laboratory under contract No.," This work was performed under the auspices of the U.S. Department of Energy, National Nuclear Security Administration by the University of California, Lawrence Livermore National Laboratory under contract No."
+ W-7105-Ene-I8., W-7405-Eng-48.
+ CLA.P. eratefully acknowledges support from the Daniel Ixirkwood Research Fuud at Indiaua University., C.A.P. gratefully acknowledges support from the Daniel Kirkwood Research Fund at Indiana University.
+ This research bas made use of the SIMBAD database. operated at CDS. Strasbourg. France.," This research has made use of the SIMBAD database, operated at CDS, Strasbourg, France."
+n ≽↼ 2()↼⋡ -u0p 2pyKk-Ov — path dey+ ΣΡΙ} cos (54),"+p_0 + i + 2 p_0 i - 3 p_0 ), and +2p_0 + i - p_0 i + p_0 )."
+ As is usual with moment hierarchies. the above equations for op) depend on third moments of the distribution function. q; and q_. the parallel and perpendicular heat fluxes.," As is usual with moment hierarchies, the above equations for $\delta p_{\Par,\Perp}$ depend on third moments of the distribution function, $q_\Par$ and $q_\perp$, the parallel and perpendicular heat fluxes."
+" Snyder οἱ introduced closure approximations lor 4 and g_ that determine dp_ and op, without solving the full kinetic equations of the previous section.", Snyder et introduced closure approximations for $q_\Par$ and $q_\perp$ that determine $\delta p_\Perp$ and $\delta p_\Par$ without solving the full kinetic equations of the previous section.
+ These Landau-fluid approxinations Close” equations (1))-(3)) and allow one to solve for the linear response of the plasma., These Landau-fluid approximations “close” equations \ref{eq:MHD1}) \ref{eq:MHD4}) ) and allow one to solve for the linear response of the plasma.
+ The linearized heat f[Inxes in the perpendicular and parallel directions are given bv =Pot and --= ~SPoGTU, The linearized heat fluxes in the perpendicular and parallel directions are given by =-p_0 and =-8 p_0 c_0^2
+GST.,GST.
+ Oue wav to reintroduce accelerating electric fields in the magnetosphere is to allow finite resistivity of the plasina., One way to reintroduce accelerating electric fields in the magnetosphere is to allow finite resistivity of the plasma.
+ Several formulations of resistive force-free equatious have been proposed. most notably by Lvutikov(2003) aud Ciyuziuov(2007.2008).," Several formulations of resistive force-free equations have been proposed, most notably by \cite{Lyutikov03} and \cite{Gruzinov07,Gruzinov08}."
+. Tn resistive ΕΕ the Oli's law can be unuubieuouslv defined iu the proper frame of the fluid by relating the current in that frame to the electric field through Jamiel Es. where 6 ds plasma conductivity (seesc...Lichnuerowicz1967:Paleuzuclaetal. 2009).," In resistive MHD, the Ohm's law can be unambiguously defined in the proper frame of the fluid by relating the current in that frame to the electric field through $\vec{j}_{\rm fluid}=\sigma \vec{E}_{\rm
+fluid}$ , where $\sigma$ is plasma conductivity \citep[see,
+e.g.,][]{Lichnerowicz67, Palenzuela09}."
+. In the force-free svstem. however. the fluid velocity along the magnetic field is unknown. aud only the transverse velocity can be obtaimed from the electromagnetic fields.," In the force-free system, however, the fluid velocity along the magnetic field is unknown, and only the transverse velocity can be obtained from the electromagnetic fields."
+ This introduces some freedom ia prescribing the Oluu's law in the force-free picture., This introduces some freedom in prescribing the Ohm's law in the force-free picture.
+ Ta the prescription proposed by Lvutikov(2003). the parallel velocity along the Ποιά was taken to be zero.," In the prescription proposed by \cite{Lyutikov03}, the parallel velocity along the field was taken to be zero."
+" In v""Stroug-Field Electrodvuaimics” (Cuuzinov2007.2008.here-after SEE).. the Oluns law was formulated iu the fane that inoves alone the field lines with such speed that the charge deusitv in that frame vanishes."," In ``Strong-Field Electrodynamics” \citep[hereafter
+SFE]{Gruzinov07,Gruzinov08}, the Ohm's law was formulated in the frame that moves along the field lines with such speed that the charge density in that frame vanishes."
+ Despite the fact that such a frame formally exists oulv iu space-lke regions. SEE prescription appears to eive smooth mmucerical solutions throughout the magnetosphere (Cauzinovy2011).. aud. most importantly. does not explode iu current sheets where maguctic fields can eo through zero.," Despite the fact that such a frame formally exists only in space-like regions, SFE prescription appears to give smooth numerical solutions throughout the magnetosphere \citep{Gruzinov11}, and, most importantly, does not explode in current sheets where magnetic fields can go through zero."
+ Both formulations are arbitrary. however. because the real fluid velocity docs not lave to follow cither frame assumption.," Both formulations are arbitrary, however, because the real fluid velocity does not have to follow either frame assumption."
+ As au additional coustraint. it is useful. therefore. to construct aresistive formulation that reproduces physical solutions expected at the extremes of very large aud very siuall conductivity of the plasma. namely the force-free πι (c.» oc) and the limit of vactiun olectromagnuoetisui (c.> 0)," As an additional constraint, it is useful, therefore, to construct a resistive formulation that reproduces physical solutions expected at the extremes of very large and very small conductivity of the plasma, namely the force-free limit $\sigma\to\infty$ ) and the limit of vacuum electromagnetism $\sigma\to 0$ )."
+ Below we describe a resistive prescription hat generalizes these schemes and combines the correct iutiug behavior of Lyutikovs scheme with current-sheet stability of Cauzinovs formulation., Below we describe a resistive prescription that generalizes these schemes and combines the correct limiting behavior of Lyutikov's scheme with current-sheet stability of Gruzinov's formulation.
+ We then apply lis prescription to nunericallv calculate the structure of resistive maguctospleres in pulsars., We then apply this prescription to numerically calculate the structure of resistive magnetospheres in pulsars.
+ Iu the resistive force-free picture of the pulsar uaenetosphere. the magnetized neutron star is thought o be surrounded bv an abundant massless plasma with finite conductivity. so that not all accelerating ficlds are shorted out.," In the resistive force-free picture of the pulsar magnetosphere, the magnetized neutron star is thought to be surrounded by an abundant massless plasma with finite conductivity, so that not all accelerating fields are shorted out."
+ For simplicity aud as a soot of principle. we only cousider the unroalistic case of constant conductivity throughout space.," For simplicity and as a proof of principle, we only consider the unrealistic case of constant conductivity throughout space."
+| More complicated prescriptions will be studied. elsewhere., More complicated prescriptions will be studied elsewhere.
+ Our finding is that uxiug our formulation of the resistive force-free clectrocdvuamics we can coustruct a family of maguetosphneres that smoothlv transition from the Deutsch vac solution to the ideal force-free maguetosphere as the conductivity of the plasma is increased., Our finding is that using our formulation of the resistive force-free electrodynamics we can construct a family of magnetospheres that smoothly transition from the Deutsch vacuum solution to the ideal force-free magnetosphere as the conductivity of the plasma is increased.
+ Such intermediate maeguetosplheres possess interesting properties that we discuss in this paper., Such intermediate magnetospheres possess interesting properties that we discuss in this paper.
+ We study the variation of the spiu-down power with magnetic juclination angle as a function of plasma conductivity aud relate it to physical couditions such as the effective potential drop ou the open field lines., We study the variation of the spin-down power with magnetic inclination angle as a function of plasma conductivity and relate it to physical conditions such as the effective potential drop on the open field lines.
+ Iu refsecicrurentder πο discuss the derivation of the resistive force-free prescription. in refsecisinimlations we describe our code for solving the equations and preseut sample magnetospheric solutious.," In \\ref{sec:currentder} we discuss the derivation of the resistive force-free prescription, in \\ref{sec:simulations} we describe our code for solving the equations and present sample magnetospheric solutions."
+ Discussion and poteutial applicatious to pulsar plivsics are dns refsecidiscussion.., Discussion and potential applications to pulsar physics are in \\ref{sec:discussion}.
+ We describe lereder our prescription for resistive current as used iu our numerical code., We describe here our prescription for resistive current as used in our numerical code.
+ Olnuus Law is defined in the fluid rest frame., Ohm's Law is defined in the fluid rest frame.
+ In this frame the electric aud naeuctic fields are parallel: otherwise. there would be a uticle drift across iiagnetie field Hues.," In this frame the electric and magnetic fields are parallel; otherwise, there would be a particle drift across magnetic field lines."
+ The laboratory rane can be connected to the fluid rest frame through wo boosts., The laboratory frame can be connected to the fluid rest frame through two boosts.
+ One boost is in the E<B direction in the ab frame and briugs the electric field parallel to the nagnetic field., One boost is in the $\vec{E}\times\vec{B}$ direction in the lab frame and brings the electric field parallel to the magnetic field.
+ The other boost is along the parallel electric and magnetic fields and transforms the b-current o the fiuid rest frame while leaving the clectric aud naenetic fields parallel., The other boost is along the parallel electric and magnetic fields and transforms the 4-current to the fluid rest frame while leaving the electric and magnetic fields parallel.
+ We choose a simple Oluns Law to relate the current and electric field in the fluid frame: μμσημα. where σ is the plasma couductivity.," We choose a simple Ohm's Law to relate the current and electric field in the fluid frame: $\vec{j}_{\rm fluid}=\sigma \vec{E}_{\rm fluid}$ , where $\sigma$ is the plasma conductivity."
+ More geueral formulations of the Oluu’s Law that take into account time-dependent currents. inertial effects. pressure. aud the Tall effect have been derived (seec.g.Meier 2001).," More general formulations of the Ohm's Law that take into account time-dependent currents, inertial effects, pressure, and the Hall effect have been derived \citep[see
+e.g.,][]{Meier04}."
+. It is likely sate to ignore most of these effects in the bulk of the strongly maguetized cold pair plasma. as cavisioned for typical pulsar niaguetosplieres: however. singular current sheets may require a more elaborate treatiuent.," It is likely safe to ignore most of these effects in the bulk of the strongly magnetized cold pair plasma, as envisioned for typical pulsar magnetospheres; however, singular current sheets may require a more elaborate treatment."
+ Iu this work we use a coustaut nuiform conductivity throughout the domain aud restrict ourselves to the simple Oluu’s Law to elucidate the basic plavsics., In this work we use a constant uniform conductivity throughout the domain and restrict ourselves to the simple Ohm's Law to elucidate the basic physics.
+ Boosting back to the laboratory frame. the current vector can be expressed as See Appendix A for a full derivation of this expression and the subsequent Buts that we discuss below.," Boosting back to the laboratory frame, the current vector can be expressed as See Appendix A for a full derivation of this expression and the subsequent limits that we discuss below."
+ Tere. Jj is the magnitude of the boost aloug the parallel electric and iuagnetie fields to the Suid rest frame. p is the charge density iu the laboratory frame. aud Ly aud By are the magnitudes of the electric aud magnetic fields in all frames in which they are parallel.," Here, $\beta_{||}$ is the magnitude of the boost along the parallel electric and magnetic fields to the fluid rest frame, $\rho$ is the charge density in the laboratory frame, and $E_0$ and $B_0$ are the magnitudes of the electric and magnetic fields in all frames in which they are parallel."
+ £y and By are defined by the expressions (Cauzinoy2007.2008) where we allow By to be positive or negative. depending ou whether the magnetic field is aligned or antiahened with the electric field.," $E_0$ and $B_0$ are defined by the expressions \citep{Gruzinov07,Gruzinov08}
+ where we allow $B_0$ to be positive or negative, depending on whether the magnetic field is aligned or antialigned with the electric field."
+ Note that iu addition to the advection of charge in the E«B direction. there is conduction current along both the laboratory frame magnetic ficld aud the clectric field (the last term in equation 11).," Note that in addition to the advection of charge in the $\vec{E}\times\vec{B}$ direction, there is conduction current along both the laboratory frame magnetic field and the electric field (the last term in equation \ref{generalcurrent}) )."
+ There is still a principal ambiguity here im defining the fluid rest frame. as we lave not vetspecified," There is still a principal ambiguity here in defining the fluid rest frame, as we have not yetspecified"
+The Subuüllietre./ Commonu-User Bolometer Avrav (SCUBA: Tolland et 11999) on the James Cerk Maxwell Telescope (CMT) has receuth been σαι tc) detect a new population of extremely luminous disaut dusty ealaxics (Snail. Ivison. Blain J9907).,"The Submillimetre Common-User Bolometer Array (SCUBA; Holland et 1999) on the James Clerk Maxwell Telescope (JCMT) has recently been used to detect a new population of extremely luminous distant dusty galaxies (Smail, Ivison, Blain 1997)."
+ By «'xploilue the strong maguification bias due to lensing through rich clusters of ealaxics. we reach fainter detection threshok Sil a giveu observing time. while suffering ittle contaminalon from cluster ealaxics (Blain 1997).," By exploiting the strong magnification bias due to lensing through rich clusters of galaxies, we reach fainter detection thresholds in a given observing time, while suffering little contamination from cluster galaxies (Blain 1997)."
+" Cosistent imuber ""0 sources have also been detected im the field (Bar vore 11998: Eales ct 11999: Toland ct 11998: IIughes e 11998).", Consistent numbers of sources have also been detected in the field (Barger et 1998; Eales et 1999; Holland et 1998; Hughes et 1998).
+ The results from these stuies have 1uporali consequences for the history of star fiOration (Blai 1e 11999) aud may also provide subsantial insight iut) the prevalence of dust-cushrouded acive ealacic nuclei (AGNs) diving the early evolution of ealaxies (Ivison e 11998)., The results from these studies have important consequences for the history of star formation (Blain et 1999) and may also provide substantial insight into the prevalence of dust-enshrouded active galactic nuclei (AGNs) during the early evolution of galaxies (Ivison et 1998).
+ Sunudl et ((1997) estinated cotnts from six galaxies detected in maps of two clusters; 3:370 and 02. using simple models of the cUsters as sinewlar isothermal11 spheres.," Smail et (1997) estimated counts from six galaxies detected in maps of two clusters, 370 and $-$ 02, using simple models of the clusters as singular isothermal spheres."
+ The complete SCUBA lens survey sample inchicdes deeper maps of these two clusters and a further five clusters mapped to a sensitivity better than 2uuuJv (Simail et 11998)., The complete SCUBA lens survey sample includes deeper maps of these two clusters and a further five clusters mapped to a sensitivity better than mJy (Smail et 1998).
+ The analysis of the naps aud the source extraction procedures are discussed by σα et ((1997)., The analysis of the maps and the source extraction procedures are discussed by Smail et (1997).
+ Tlic| full catalog of the SCUBA and follow-up data cau be found. elsewhere (Simail et. 11999)., The full catalog of the SCUBA and follow-up data can be found elsewhere (Smail et 1999).
+ Iu this new selfcoutained analysis of the whole Ίος SHTVOV sample. we exploit. lutxlels of the cluster potentials derived from deep (LEST)) aud lugh-resolution eround-based optical imaging aud the LENSTOOL vay-tracing code (ος et 11993) to correct for the effects of lens amplification aud to calculate robust nuniber counts of faint sub-nuu galaxies below the flux limit of previously published work.," In this new self-contained analysis of the whole lens survey sample, we exploit models of the cluster potentials derived from deep ) and high-resolution ground-based optical imaging and the LENSTOOL ray-tracing code (Kneib et 1993) to correct for the effects of lens amplification and to calculate robust number counts of faint sub-mm galaxies below the flux limit of previously published work."
+ The caleulatious are carried out in an Einsteinde Sitter world model. but the results depend oulv weakly on the world model assunned.," The calculations are carried out in an Einstein–de Sitter world model, but the results depend only weakly on the world model assumed."
+ Iu a blauk field the cimative sub-unuu galaxy countsthe surface density of galaxies brighter than a eiven flux density limitcan be derived by simply dividing the mumber of detected sources by the surveyed area. although the effects of clustering may have to be considered in small fields.," In a blank field the cumulative sub-mm galaxy counts---the surface density of galaxies brighter than a given flux density limit—can be derived by simply dividing the number of detected sources by the surveyed area, although the effects of clustering may have to be considered in small fields."
+ The calculation is more complicated in the field of a eravitational lens. which magnifies aud distorts the background skythe source planeby an amount that," The calculation is more complicated in the field of a gravitational lens, which magnifies and distorts the background sky—the source plane—by an amount that"
+αμ follows plausible arguueuts. we max hope to obtain the same sign for such a dynamically deternuned a (Braudenbure & Sclunitt 1998).,"$\alpha_{\rm H}$ follows plausible arguments, we may hope to obtain the same sign for such a dynamically determined $\alpha$ (Brandenburg & Schmitt )."
+ As far as the dependence of o on rotation is concerned. the simulation nav vield more reliable results than for the other parameters.," As far as the dependence of $\alpha$ on rotation is concerned, the simulation may yield more reliable results than for the other parameters."
+ Iu the solar convection zone the Coriolis nuniberis suall near the surface. and increases to values around 1 or somewhat larger near its base.," In the solar convection zone the Coriolis number is small near the surface, and increases to values around 1 or somewhat larger near its base."
+ This is the reenne shown in Fig., This is the regime shown in Fig.
+. The rotational quenching of ay aud the saturation of απ (Fig.5," The rotational quenching of $\alpha_{\rm V}$ and the saturation of $\alpha_{\rm H}$ (Fig.,"
+. top) nav therefore just Ce iu he depth region where the dynam operates. mt in solar-type stars which are vouuger and therefore rotating faster. these effects may be important.," top) may therefore just begin in the depth region where the dynamo operates, but in solar-type stars which are younger and therefore rotating faster, these effects may be important."
+ We have observed some agrecinent between the uunerical results oro aud the predictionsof Rüddiger & Wichatinoy (1993). as far as the dependence on rotation aud on depth is concerned.," We have observed some agreement between the numerical results for $\alpha$ and the predictions of Rüddiger & Kichatinov ), as far as the dependence on rotation and on depth is concerned."
+ Iu both the uustable aud the stable laver. he approximate C'oriolis number dependence and the correct sien for oq; are reproduced.," In both the unstable and the stable layer, the approximate Coriolis number dependence and the correct sign for $\alpha_{\rm H}$ are reproduced."
+ The main discrepaucies concern the amplitude of à (the scaling factors of the heoretical curves in Fie.oO 5), The main discrepancies concern the amplitude of $\alpha$ (the scaling factors of the theoretical curves in Fig. )
+ aud the signOo of av\ in the stable laver. as well as the sign of the small-scale current relicity.," and the sign of $\alpha_{\rm V}$ in the unstable layer, as well as the sign of the small-scale current helicity."
+ Even within the lanited range accessible to αποΊσα] simulation. the parameter space is enormous.," Even within the limited range accessible to numerical simulation, the parameter space is enormous."
+" We have not vet considered the latitude dependence of a (although S1ue earlier work of Brandenburg (1991) indicates that [o] is maximum at about GO"" of latitude). and we have not vet explored the depeudeucies on the Revuolds and Ravleigh wuubers."," We have not yet considered the latitude dependence of $\alpha$ (although some earlier work of Brandenburg ) indicates that $|\alpha|$ is maximum at about $\pm 60^\circ$ of latitude), and we have not yet explored the dependencies on the Reynolds and Rayleigh numbers."
+ This must be done by further siuulations., This must be done by further simulations.
+and for a precursor to a long burst (3.5 s) from SCR Lo00|Lt (Coraletal. 20013).,and for a precursor to a long burst (3.5 s) from SGR 1900+14 \cite{ibrahim}) ).
+ The παλιο ind of spectral evolution has also been reported for a ~9 s ong burst from SCR 0526-66 (Colenctskiietal.L9S8T)): he softening trend is seen in the first three of the four spectra extracted. while the last oue is as lard as the first one.," The same kind of spectral evolution has also been reported for a $\sim$ 9 s long burst from SGR 0526-66 \cite{golenetskii}) ): the softening trend is seen in the first three of the four spectra extracted, while the last one is as hard as the first one."
+ Iu our sample we do not find evidence for this sud of evolution., In our sample we do not find evidence for this kind of evolution.
+ Ou the other hand. soft-to-hard evolution has Όσοι seen with the DATSE iustruneut for two peculiar masts very likely originating from SCR 1900]TL (Woodsetal. 1999)).," On the other hand, soft-to-hard evolution has been seen with the BATSE instrument for two peculiar bursts very likely originating from SGR 1900+14 \cite{woods}) )."
+ These two bursts were quite different from the usual bursts. both in terius of duration (lasting  1 s). aud spectral harcduess (AL of the order of LOO keV).," These two bursts were quite different from the usual bursts, both in terms of duration (lasting $\sim$ 1 s), and spectral hardness $kT$ of the order of 100 keV)."
+ Tn the framework of the magnetar ποσο] (Duncan&Thompson 1992)). short (~0.1 3) SCOR bursts are usually described as the radiation originating frou he cooling of an optically thick pair-photon plasma.," In the framework of the magnetar model \cite{dt92}) ), short $\sim$ 0.1 s) SGR bursts are usually described as the radiation originating from the cooling of an optically thick pair-photon plasma."
+ This asma is generated in the neutron star maguctosplere by an Alfvénn pulse. which is trigecred by a sudden shift iu he magnetospheric footpoiuts divenu wea fracture in the ieutron star crust (Thompson&Duncan 1995).," This plasma is generated in the neutron star magnetosphere by an Alfvénn pulse, which is triggered by a sudden shift in the magnetospheric footpoints driven by a fracture in the neutron star crust \cite{td95}) )."
+" This nodel is able to explain the time histories aud energeties of the typical SCR bursts, aud predicts that the effective eniperature of the specra should vary weakly during the masts. owing to the constancy of the photosphierie energy Hux."," This model is able to explain the time histories and energetics of the typical SGR bursts, and predicts that the effective temperature of the spectra should vary weakly during the bursts, owing to the constancy of the photospheric energy flux."
+ No detailed predictions are available concerning nore complex spectral evolution as we observe in some of the bursts., No detailed predictions are available concerning more complex spectral evolution as we observe in some of the bursts.
+ For exiiuple he model does not account or the presence of a soft precursor as observe in burst munber 2 (sec Fie. 2))., For example the model does not account for the presence of a soft precursor as observed in burst number 2 (see Fig. \ref{trehr}) ).
+ Qur results indicate that a harcuess-iutensity auti-correlation (which in παν bursts manifests itsclf as a soft-to-hard time evolution) is xresent in bursts from which are uot particularly long. nor PAρουταν hard aud not at al VOYV energetic.," Our results indicate that a hardness-intensity anti-correlation (which in many bursts manifests itself as a soft-to-hard time evolution) is present in bursts from which are not particularly long, nor spectrally hard and not at all very energetic."
+ It is interesting to note that this correlation is opposite to what found for the bursts emitted from 1E 2259|586 (αντetal. 2001)). which although have lower fiueuces than the ones we neasure.," It is interesting to note that this correlation is opposite to what found for the bursts emitted from 1E 2259+586 \cite{gavrill}) ), which although have lower fluences than the ones we measure."
+independent ways: with the MIDI Interactive Analysis (MIA) package. which uses the power spectrum method. and the Expert. Work. Station (EWS) package. which is based on à coherent. linear averaging method.,"independent ways: with the MIDI Interactive Analysis (MIA) package, which uses the power spectrum method, and the Expert Work Station (EWS) package, which is based on a coherent, linear averaging method."
+ For the extraction. of the spectra the two consider fitted and fix masks. respectively.," For the extraction of the spectra the two consider fitted and fix masks, respectively."
+ The interferometric results were then compared and showed agreement better than10%., The interferometric results were then compared and showed agreement better than.
+. All available calibrator data obtained on March 2. 2005 were reduced in order to determine the errors of the spectrally resolved visibilities.," All available calibrator data obtained on March 2, 2005 were reduced in order to determine the errors of the spectrally resolved visibilities."
+ After investigating the data we decided to apply a uniform multiplicative error of as a conservative estimate., After investigating the data we decided to apply a uniform multiplicative error of as a conservative estimate.
+ At both epochs the N-band spectra obtained by the two telescopes differ in their absolute flux levels by a factor of zz ].4. and their spectral shapes slightly deviate at A>12 μπι. The reason of this artifact is not yet clarified. thus we evaluated the acquisition Images and extracted flux density values at 8.7 jim for OOr! (for the calibrator. 337160. 11.85 Jy was assumed).," At both epochs the N-band spectra obtained by the two telescopes differ in their absolute flux levels by a factor of $\approx$ 1.4, and their spectral shapes slightly deviate at ${\lambda}{>}12\,\mu$ m. The reason of this artifact is not yet clarified, thus we evaluated the acquisition images and extracted flux density values at $8.7\,\mu$ m for Ori (for the calibrator, 37160, 11.85 Jy was assumed)."
+ Photometry taken with UT3 turned out to be more reliable. judged from the repeatibility of the standard star measurements. and we estimate its absolute accuracy to be15%.," Photometry taken with UT3 turned out to be more reliable, judged from the repeatibility of the standard star measurements, and we estimate its absolute accuracy to be."
+". The absolute calibration was then carried out by scaling the N-band spectra to the photometric points at 5.7 jm, reffig:fig] shows the calibrated visibilities as a function of wavelength."," The absolute calibration was then carried out by scaling the N-band spectra to the photometric points at $8.7\,\mu$ m. \\ref{fig:fig1} shows the calibrated visibilities as a function of wavelength."
+ The values gradually decrease from 0.87 to 0.75 between 8 and /m. and are approximately constant at longer wavelengths. with a small dip between 12 and 134m. This pattern clearly demonstrates that the emitting region Is resolved. and suggests à non-uniform temperature distribution of the emitting material (otherwise the visibilities would monotonically increase towards longer wavelengths due to the lowering spatial resolution).," The values gradually decrease from 0.87 to 0.75 between 8 and $\mu$ m, and are approximately constant at longer wavelengths, with a small dip between 12 and $\mu$ m. This pattern clearly demonstrates that the emitting region is resolved, and suggests a non-uniform temperature distribution of the emitting material (otherwise the visibilities would monotonically increase towards longer wavelengths due to the lowering spatial resolution)."
+ Adopting Eq., Adopting Eq.
+ | from Leinert et al. (, 1 from Leinert et al. (
+2004). we calculated visibility values for a series of exteded model sources of increasing sizes. represented by Gaussian brightness profiles. and compared them with our observations at different wavelengths.,"2004), we calculated visibility values for a series of extended model sources of increasing sizes, represented by Gaussian brightness profiles, and compared them with our observations at different wavelengths."
+" The derived FWHMs increase towards longer wavelengths from 6!5 mmas at ym to 16+3mmas_ at 13,m. At the distance of V1I64700ri 350 pe) these angular sizes approximately correspod to 2.7 and AAU linear sizes. respectively."," The derived FWHMs increase towards longer wavelengths from $^{+2}_{-3}$ mas at $\mu$ m to $\pm$ mas at $\mu$ m. At the distance of Ori (450 pc) these angular sizes approximately correspond to 2.7 and AU linear sizes, respectively."
+ The increasing sizes of Gaussians with increasing wavelength indicate a radially decreasing temperature profile., The increasing sizes of Gaussians with increasing wavelength indicate a radially decreasing temperature profile.
+ reffig:fig3 presents our 4m spectrum of V1647OO0r obtained with MIDI on UT3 in March 2005., \\ref{fig:fig3} presents our $\mu$ m spectrum of Ori obtained with MIDI on UT3 in March 2005.
+ The spectrum ts smoothly rising towards longer wavelengths. and no obvious spectral features (silicates. PAHs. molecular ices) can be seen.," The spectrum is smoothly rising towards longer wavelengths, and no obvious spectral features (silicates, PAHs, molecular ices) can be seen."
+ The shape of the December 2004 MIDI spectrum is identical to the plotted spectrum within the measurement uncertainties., The shape of the December 2004 MIDI spectrum is identical to the plotted spectrum within the measurement uncertainties.
+ For comparison. in Fig.," For comparison, in Fig."
+ 2) we also overplotted an earlier spectrum of 00r1i from March. 2004 (see figure caption)., \ref{fig:fig3} we also overplotted an earlier spectrum of Ori from March 2004 (see figure caption).
+ Comparison of the spectra suggests that the N-band spectrum changed dramatically (especially at longer wavelengths) between March and December 2004. but exhibited an approximately constant spectral shape afterwards. between December 2004 and March 2005.," Comparison of the spectra suggests that the N-band spectrum changed dramatically (especially at longer wavelengths) between March and December 2004, but exhibited an approximately constant spectral shape afterwards, between December 2004 and March 2005."
+ We calculated the spectrum of the correlated flux as the product of the measured N-band spectrum and the spectrally resolved visibilities. in order to study the emission of the innermost part of the circumstellar structure.," We calculated the spectrum of the correlated flux as the product of the measured N-band spectrum and the spectrally resolved visibilities, in order to study the emission of the innermost part of the circumstellar structure."
+ We found that the main fraction of the observed N-band flux (= 70%) is emitted in the inner regions., We found that the main fraction of the observed N-band flux $\approx 70\% $ ) is emitted in the inner regions.
+ The relatively flat shape of the correlated spectrum indicates that the innermost part has higher temperature., The relatively flat shape of the correlated spectrum indicates that the innermost part has higher temperature.
+ Like in the full spectrum. no spectral features can be observed in the innermost spectrum.," Like in the full spectrum, no spectral features can be observed in the innermost spectrum."
+The non-deal contrübutious to the chemical potentials of ITr aud II» were obtained through the numerical solution of the Orustein-Zeruike equation in the Percus-Yevick (PY) approxiiiatiou (AMlartvnoyv,"The non-ideal contributions to the chemical potentials of $\rm H\,\textrm{\scriptsize{I}}$ and $\rm H_2$ were obtained through the numerical solution of the Ornstein-Zernike equation in the Percus-Yevick (PY) approximation \citep{M}. ."
+ 1992).. For the IIIe interaction we use the pair potential of Shalabictal.(1998).. aud forthe ITWe interaction the pair potential eiven by Ree(1983) (Figue 1).," For the $\rm H - He$ interaction we use the pair potential of \citet{SH}, and for the $\rm H_2 - He$ interaction the pair potential given by \citet{RE} (Figure 1)."
+ Both poteutials are οι ab. nitio quautu mechanical calculations and are in good agreement with the independent calculations of Tang&Yang(1990) for IDTe and of Tao(1993) and Shafer&Cordon(1976) for Il Πο., Both potentials are from ab initio quantum mechanical calculations and are in good agreement with the independent calculations of \citet{TY} for $\rm H-He$ and of \citet{TAO} and \citet{SG} for $\rm H_2-He$ .
+" As we consider a helimm-dominated iuxture (Πο= 102). the IT.IL. Il,ΠΠ. ancl Il;IL interactions can be ueelected."," As we consider a helium-dominated mixture $\rm He/H\wig>10^{2}$ ), the $\rm H-H$, $\rm H_2 - H$, and $\rm H_2 - H_2$ interactions can be neglected."
+ Ileh pressure experiueuts have shown that ab initio pair poteutials are too repulsive to describe dense systems where N-body effects become iuportaut (Nellisetal.198I:Ross 1983).," High pressure experiments have shown that ab initio pair potentials are too repulsive to describe dense systems where N-body effects become important \citep{Nellis,Ross}."
+. The softening of the pair potentials at high densities can oulv be quantified experimentally or. alternatively. estimated with N-body quiutui mechanical calculations.," The softening of the pair potentials at high densities can only be quantified experimentally or, alternatively, estimated with N-body quantum mechanical calculations."
+ Since neither are available for nüxtures of trace hydrogen in heliuui we resort to ab initio potentials.," Since neither are available for mixtures of trace hydrogen in helium, we resort to ab initio potentials."
+ The net effect ou the dissociation equilibriun depends on the relative softeniug of the potentials (Eq., The net effect on the dissociation equilibrium depends on the relative softening of the potentials (Eq.
+ 5) aud is therefore less sensitive to N-body effects than the individual poteutials., 5) and is therefore less sensitive to N-body effects than the individual potentials.
+ We also calculated the chemical poteutiala iu the Uspeructted Chain approximation and found them to agree within 5% with the PY values up to 2e/cem’.," We also calculated the chemical potentials in the Hypernetted Chain approximation and found them to agree within $5\%$ with the PY values up to $\rm 2 \, g/cm^{3}$."
+ Since the PY approximation is better suited for short range potentials such as the ones we use here. we estimate that our PY caleulatious are reliable up to at least 2efem’.," Since the PY approximation is better suited for short range potentials such as the ones we use here, we estimate that our PY calculations are reliable up to at least $\rm 2 \, g/cm^{3}$."
+ For the internal partition functions Z;. we use expressions for the olectrouic ground state of the unperturbed lvdrogen imelecule accounting for the vibrational/rotational excitatious (προ&Herzberg 1979).. and set Zjj=2 for lvdrogen atom.," For the internal partition functions $Z_i$ , we use expressions for the electronic ground state of the unperturbed hydrogen molecule accounting for the vibrational/rotational excitations \citep{HH}, and set $Z_{\rm H \,\textrm{\tiny{I}}}=2$ for hydrogen atom."
+ This approximation is justified as the electronic excitation energies of both species are ]luee aud for temperatures of a few thousands degrees the populatious of the electronic excited levels are extremely simall., This approximation is justified as the electronic excitation energies of both species are large and for temperatures of a few thousands degrees the populations of the electronic excited levels are extremely small.
+" Towever. there dssienificant thermal excitation of the rotational audvibrational levels of Πο audthe effect of the deuse inedit ou Zip, lust be considered."," However, there issignificant thermal excitation of the rotational andvibrational levels of $\rm H_2$ andthe effect of the dense medium on $Z_{\rm H_2}$ must be considered."
+ Since, Since
+"fraction at the central redshift (xmrc), redshift range, neutral fraction range and the range of the rms variations in the 21-cm signal for the different light cone cubes from simulations L1 and L3 are given in Tables 1 and 2 respectively.","fraction at the central redshift $x_{\rm{H Ic}}$ ), redshift range, neutral fraction range and the range of the rms variations in the 21-cm signal for the different light cone cubes from simulations L1 and L3 are given in Tables \ref{tab:table_L1} and \ref{tab:table_f25} respectively."
+ The leftpanels of Figure., The leftpanels of Figure.
+ 4 show the spherically averaged 3D power spectra A3p(k) for light cone cubes at four different central redshifts., \ref{fig:ps3de_f10} show the spherically averaged 3D power spectra $\Delta^2_{\rm{3D}}(k)$ for light cone cubes at four different central redshifts.
+" For comparison it also shows the power spectra for coeval cubes at three redshifts (high, low redshift end and central redshift of the light cone cube)."," For comparison it also shows the power spectra for coeval cubes at three redshifts (high, low redshift end and central redshift of the light cone cube)."
+ The righthand panels plot the relative difference in the power spectra between the coeval cube at the central redshift and the light cone cube., The righthand panels plot the relative difference in the power spectra between the coeval cube at the central redshift and the light cone cube.
+ General features we see in all panels are that the effect is stronger on large scales and increases as we go up in scale., General features we see in all panels are that the effect is stronger on large scales and increases as we go up in scale.
+" In addition, we find that the power is enhanced (suppressed) with respect to the coeval cube at large (small) scales for neutral fractions zgr<0.5."," In addition, we find that the power is enhanced (suppressed) with respect to the coeval cube at large (small) scales for neutral fractions $x_{\rm{H I}}\lesssim 0.5$."
+ During the last phase of reionization (bottom panels) the power in the light cone cube is suppressed at all scales., During the last phase of reionization (bottom panels) the power in the light cone cube is suppressed at all scales.
+ At redshift z—11.20 we see that the power spectrum is enhanced by ~15% at modes k«0.5Mpc™' and suppressed by ~10% at small scales.," At redshift $z=11.20$ we see that the power spectrum is enhanced by $\sim 15 \%$ at modes $k<0.5 \, \rm{Mpc}^{-1}$ and suppressed by $\sim 10 \%$ at small scales."
+ At redshift z—9.94 we do not see much effect except on the largest scale where the power is larger by ~ 15%., At redshift $z=9.94$ we do not see much effect except on the largest scale where the power is larger by $\sim 15 \%$ .
+The svuchrotron eumissivitv at a eiven frequency 7 of al isotropic distribution of electrons in a ταοτι] oricuted distribution of maeuctic fields is where ο3=νοτοF(2nnne). vp.D)=BeBy (πι). and the angle averaged dinenusionuless spectral chussivity of a monoenergetic Isotropic electrou distribution (Crusius Schlickeiser 1986) is with Wage) Πο Whittalker's function (Abramowitz Steeun 19653).,"The synchrotron emissivity at a given frequency $\nu$ of an isotropic distribution of electrons in a randomly oriented distribution of magnetic fields is where $c_3 = \sqrt{3} e^3/(4 \pi \me c^2) $, $\nu_{\rm c}( p,B)=3 e B
+p^2 /(4\pi \me c)$ , and the angle averaged dimensionless spectral emissivity of a monoenergetic isotropic electron distribution (Crusius Schlickeiser 1986) is with $W_{\rm \lambda,\mu}(x)$ denoting Whittaker's function (Abramowitz Stegun \nocite{abramowitz65}) )."
+ For) can be approximated to an accuracy of a few percent by which is faster to evaluate nuinericallv., $\tilde{F}(x)$ can be approximated to an accuracy of a few percent by which is faster to evaluate numerically.
+ Fie., Fig.
+ | shows he svuchrotron enuüssoun of a cooling population of clectrous in the two models eiveu above., \ref{fig:sync} shows the synchrotron emission of a cooling population of electrons in the two models given above.
+ The imuportaut result of this simple iuhomoseneous uodel is that the electron. population im the weak field regions is still prescut iu the energv range of 110 350 MeW after 2 Car of cooling. so that it can produce 1ο EUW excess by IC scattering of CAIB photons.," The important result of this simple inhomogeneous model is that the electron population in the weak field regions is still present in the energy range of 140 – 350 MeV after 2 Gyr of cooling, so that it can produce the EUV excess by IC scattering of CMB photons."
+" But svuchrotron/IC cooling produces a sharp cutoff at Ισ] energics, so that for any given normalization of the electron population no observable radio emission reais."," But synchrotron/IC cooling produces a sharp cutoff at higher energies, so that for any given normalization of the electron population no observable radio emission remains."
+ Iu the case of homogeneous fields a considerable zaunouut of fine tuning of the time point of injection is necessi in order to allow a large electron population in the 110 350 MeV. rauge. without overproducing low frequency radio enmuüssionu. and therefore violating the observational constraints.," In the case of homogeneous fields a considerable amount of fine tuning of the time point of injection is necessary in order to allow a large electron population in the 140 – 350 MeV range, without overproducing low frequency radio emission, and therefore violating the observational constraints."
+ As long as the magnetic fields are not lower than 64€. the cooling time has to be near 0.6 Civi. otherwise either the cutoff is lower than 100 MeV. or the clectrous are visible iu the radio.," As long as the magnetic fields are not lower than $6\,\mu$ G, the cooling time has to be near 0.6 Gyr, otherwise either the cutoff is lower than 400 MeV, or the electrons are visible in the radio."
+ An clectrou differential umber index of 2.5 for the EUV producing energv range can casily be matched for sufficicutly flat injection spectra., An electron differential number index of 2.5 for the EUV producing energy range can easily be matched for sufficiently flat injection spectra.
+ Unfortunatelv. the preseut day differential ΠΙΟΣ ineex depends not only ou the injection index and the cooling time. but also sensitively on the distrinition of fiel strengh.," Unfortunately, the present day differential number index depends not only on the injection index and the cooling time, but also sensitively on the distribution of field strength."
+ For the uiforii distribution of field streneth assuiied above au injection index of 2.15 would lead to a present day iudex of zm2.5 in this enerev range., For the uniform distribution of field strength assumed above an injection index of $2.15$ would lead to a present day index of $\approx 2.5$ in this energy range.
+ If the feld distribution is more strongly weighted to lower ficld streusth. a steeper ijecjection spectral index would be required. and if fewer low inagnetie field regious exist a flater one.," If the field distribution is more strongly weighted to lower field strength, a steeper injection spectral index would be required, and if fewer low magnetic field regions exist a flatter one."
+" Wih our poor present day knowleeo about the distribution of field strengthso it is therefore impossible to derive the injection ""pectrulni.", With our poor present day knowledge about the distribution of field strengths it is therefore impossible to derive the injection spectrum.
+ The energy deusity of the relativistic electrons iu the oethomogencous niodel above after 2 Gyr is nearlv two orders of magnitude smaller than the injected energy density. assunune that the injected spectrum exteuds from D.=1 MeV/c to 10! MeVfc. and the injection ifereutial uunnber index is 2.15.," The energy density of the relativistic electrons in the inhomogeneous model above after 2 Gyr is nearly two orders of magnitude smaller than the injected energy density, assuming that the injected spectrum extends from $P_{\rm e} = 1$ MeV/c to $10^4$ MeV/c, and the injection differential number index is 2.15."
+ This energy loss depends stronely on the assumed distribution of fields: if more weak feld regions are present the euergv loss is mach less dramatic., This energy loss depends strongly on the assumed distribution of fields: if more weak field regions are present the energy loss is much less dramatic.
+ And since the present dav population of EUV excess producing electrons las an energv density which is a factor of a few πιάνο low the thermal energv density (see Fig., And since the present day population of EUV excess producing electrons has an energy density which is a factor of a few hundred below the thermal energy density (see Fig.
+ 5., 5.
+ in Bowyer Berehoer 1998). an injection cucrey deusitv considerably lower thau the preseut thermal enerey density would be suicient to explain the preseut dav EUV excess.," in Bowyer Berghöffer 1998), an injection energy density considerably lower than the present thermal energy density would be sufficient to explain the present day EUV excess."
+ The effects of radius dependent electron aud magnetic fiel istributious can iu principle be treated with a simular formali. where the volnuue average is then over elongated beams along the line of sight.," The effects of radius dependent electron and magnetic field distributions can in principle be treated with a similar formalism, where the volume average is then over elongated beams along the line of sight."
+ Iustead. of ewig a detailed model. which would rely on even more assunrptious. we briefly discuss the qualitative behaviour.," Instead of giving a detailed model, which would rely on even more assumptions, we briefly discuss the qualitative behaviour."
+ The magnetic fields streneth and the injected electron energv deusitv should decrease with radius., The magnetic fields strength and the injected electron energy density should decrease with radius.
+ Due to the lugher fraction of strong feld regions iu the ceuter. one would expect a stronger cooling of the clectrous there.," Due to the higher fraction of strong field regions in the center, one would expect a stronger cooling of the electrons there,"
+feature noar 1051. However. quautitativelv the fit is horrible: 4?=3015 for 91 degrees o: freedom.,"feature near $\mu$ m. However, quantitatively the fit is horrible: $\chi^2= 3045$ for 94 degrees of freedom."
+ This is primarily due to the fact that the short-vaveleusth peak is more prominent 1u the model than iu the «observations aud the silicate absorption and enmüssion features in the model are offset in wwleugth of the observed catures., This is primarily due to the fact that the short-wavelength peak is more prominent in the model than in the observations and the silicate absorption and emission features in the model are offset in wavelength of the observed features.
+ It is possible that optimizing the parameters will improve the fit as the eid is quite coarse aud some of tli| best-fit paraicters are at the extreme values of the exid. such as 0cones," It is possible that optimizing the parameters will improve the fit as the grid is quite coarse and some of the best-fit parameters are at the extreme values of the grid, such as $\theta_{cone}$."
+" Also. the viewing anele is 71 which is oulv “within the openiusg augle of the cone: however. viewing augles ο SL"" and 89° are very simular in shape. but only sliehtlv poorer fits. plotted as dotted lines in Figure 6.."," Also, the viewing angle is $71^\circ$ which is only $^\circ$ within the opening angle of the cone; however, viewing angles of $^\circ$ and $^\circ$ are very similar in shape, but only slightly poorer fits, plotted as dotted lines in Figure \ref{fig:q2237_fritz}."
+ Consequently we do not believe that the fitted paramicters are unique or even correct: indeed the simple geometry chosen by Fritz e nunay be svroug., Consequently we do not believe that the fitted parameters are unique or even correct; indeed the simple geometry chosen by Fritz et may be wrong.
+ The main point is that a dusv torus model can produce a fair qualitative fit to the SED o:0305: further development of theoretical models will be required. to obtain a better quantitative fit., The main point is that a dusty torus model can produce a fair qualitative fit to the SED of; further development of theoretical models will be required to obtain a better quantitative fit.
+ Figure 7 shows the ratios of the quasar image fixes as a funcion of wavelength frou: Table 2.., Figure \ref{fig:q2237ratios} shows the ratios of the quasar image fluxes as a function of wavelength from Table \ref{tabflux}.
+ The V-baud data are the data obtained from the OGLE data archive (2). taken at a time closest to our IRAC observations: 2005 116 00:18 UT CIID)., The $V$ -band data are the data obtained from the OGLE data archive \citep{Udalski2006} taken at a time closest to our IRAC observations: 2005 16 00:48 UT (HJD).
+ The 10 micron points are from ?.. aud tlas are not smnmultaueous to our Spitzer observatious: however there appears to be little variability at this waveleugth.," The 10 micron points are from \citet{Agol2000}, and thus are not simultaneous to our Spitzer observations; however there appears to be little variability at this wavelength."
+ Also plotted are the model fux ratios from ? which is the most complete mocdel of the lens ealaxy of cconstructed to date. iuclhiding the bar aud spiral aris.," Also plotted are the model flux ratios from \citet{Trott2002} which is the most complete model of the lens galaxy of constructed to date, including the bar and spiral arms."
+ The agreenent with the 1) idcron flux ratios is expected since these were used as a coustraint on the model: however. the model eljves very Slar flux ratios as an earlier model by ο which was constructed before the 10 micron observations.," The agreement with the 10 micron flux ratios is expected since these were used as a constraint on the model; however, the model gives very similar flux ratios as an earlier model by \citet{Schmidt1998} which was constructed before the 10 micron observations."
+ The general treud is obvious in Figure 7:5 for all three pairs o πμασος there is a strong nücroleusiug anomaly in the optical which gradually. disappears towards longer wavelengths., The general trend is obvious in Figure \ref{fig:q2237ratios}: for all three pairs of images there is a strong microlensing anomaly in the optical which gradually disappears towards longer wavelengths.
+ This is precisely the behavior expected from the standard model of quasars: the source should be larger at longer waveleneths ame thus less affected by icroleusiug., This is precisely the behavior expected from the standard model of quasars: the source should be larger at longer wavelengths and thus less affected by microlensing.
+A theoretical approach to this problem requires studies of the thermal. chemical. and excitation structure of disks at a level of detail that is appropriate for comparison with observations.,"A theoretical approach to this problem requires studies of the thermal, chemical, and excitation structure of disks at a level of detail that is appropriate for comparison with observations."
+ Studies of the possibility of non-LTE CO level populations in low accretion rate svstenis would be particularly welcome., Studies of the possibility of non-LTE CO level populations in low accretion rate systems would be particularly welcome.
+ For a more empirical approach to the problem. we might compare the observed line profiles of V836 Tau with other low accretion rate svstems in which the disk is expected to be racially continuous within a few AU.," For a more empirical approach to the problem, we might compare the observed line profiles of V836 Tau with other low accretion rate systems in which the disk is expected to be radially continuous within a few AU."
+ As an exanple of the latter approach. we discussed the CO fundamental line profiles of LkCal5. a T Tauri star with a low accretion rate simular to that of V836 Tau.," As an example of the latter approach, we discussed the CO fundamental line profiles of LkCa15, a T Tauri star with a low accretion rate similar to that of V836 Tau."
+ The more centrally peaked line profiles of LkCal5. if representative of other low accretion rate svstems. would suggest that the double-peaked emission profiles of V826 Tau arise Irom a physically truncated inner disk.," The more centrally peaked line profiles of LkCa15, if representative of other low accretion rate systems, would suggest that the double-peaked emission profiles of V836 Tau arise from a physically truncated inner disk."
+ such a physically truncated inner disk might arise if (he svstem has formed a Jovian mass planet (hat has cleared a gap in the disk., Such a physically truncated inner disk might arise if the system has formed a Jovian mass planet that has cleared a gap in the disk.
+ A simple fit to the SED of V336 Tan is consistent with a flared disk that has an optically thin region within ~1 AAU., A simple fit to the SED of V836 Tau is consistent with a flared disk that has an optically thin region within $\sim 1$ AU.
+ Thus. a possible interpretation of the data is Chat an orbiting companion has created a gap between a easeots inner disk within 0.4AAU and an optically thick outer disk bevond AAU.," Thus, a possible interpretation of the data is that an orbiting companion has created a gap between a gaseous inner disk within AU and an optically thick outer disk beyond AU."
+ Since (he above fit to the SED is non-unique. it would be useful to use both improved SED modeling techniques (e.g. Calvet et 22005) and infrared interferometry (e.g.. Ratzka οἱ 22007) to test the hypothesis that the dust disk has an optically thin gap or inner hole.," Since the above fit to the SED is non-unique, it would be useful to use both improved SED modeling techniques (e.g., Calvet et 2005) and infrared interferometry (e.g., Ratzka et 2007) to test the hypothesis that the dust disk has an optically thin gap or inner hole."
+ In svstenis that have formed a Jovian mass planet. small grains παν be filtered out οἱ (he inward accretion flow at the outer edge of the gap (Rice et 22006). rendering the dust distribution a poor tracer of (he physical structure of the disk al smaller radii.," In systems that have formed a Jovian mass planet, small grains may be filtered out of the inward accretion flow at the outer edge of the gap (Rice et 2006), rendering the dust distribution a poor tracer of the physical structure of the disk at smaller radii."
+ Gaseous disk tracers. like the CO fundamental emission discussed here. may then beneeded to probe disk structure al (hese smaller radii.," Gaseous disk tracers, like the CO fundamental emission discussed here, may then be to probe disk structure at these smaller radii."
+ Thus. there is considerable motivation to expand the study of gaseous disk diagnostics bevond (he present case. to understand more generally whether and how well diagnoses such as CO Dundamental emission can probe the racial structure ol gaseous disks.," Thus, there is considerable motivation to expand the study of gaseous disk diagnostics beyond the present case, to understand more generally whether and how well diagnostics such as CO fundamental emission can probe the radial structure of gaseous disks."
+ We are grateful to Steve Strom for stimulating and insightful discussions on (his topic., We are grateful to Steve Strom for stimulating and insightful discussions on this topic.
+ We also thank Lisa Prato for communicating her radial velocity results in advance of publication., We also thank Lisa Prato for communicating her radial velocity results in advance of publication.
+ Financial support for this work was provided by the NASA Origins of Solar Systems program (NNIIOTAC5II) and the NASA Astrobiology Institute under Cooperative Agreement CCAN-02-OSS-02 issued through the Office of Space Science., Financial support for this work was provided by the NASA Origins of Solar Systems program (NNH07AG51I) and the NASA Astrobiology Institute under Cooperative Agreement CAN-02-OSS-02 issued through the Office of Space Science.
+ This work was also supported bv the Life and Planets Astrobiologv Center (LAPLACT)., This work was also supported by the Life and Planets Astrobiology Center (LAPLACE).
+ Basie research in infrared astronomy al the Naval Research Laboratory is supported by 6.1 base funding., Basic research in infrared astronomy at the Naval Research Laboratory is supported by 6.1 base funding.
+ The authors wish to recognize and acknowledge the very significant cultural role and reverence that (he summit ol Mauna ea has always had within the indigenous [Hawaiian commmnitv., The authors wish to recognize and acknowledge the very significant cultural role and reverence that the summit of Mauna Kea has always had within the indigenous Hawaiian community.
+ We are most, We are most
+(Out) image (Figure 11)).,] image (Figure \ref{fig11}) ).
+ The 2-10 keV observed flux aud luminosity of this gas are τ1x10.1! ere τν 107 erg s.1!. respectively.," The 2–10 keV observed flux and luminosity of this gas are $7.4 \times 10^{-14}$ erg $^{-2}$ $^{-1}$ and $1.4 \times 10^{38}$ erg $^{-1}$, respectively."
+ The X-ray spectrum of tie NW extension above 3 keV is similar to that observed in the nucleus. aud is inarginally consistent with the putative unobscured uuclear spectrum.," The X-ray spectrum of the NW extension above 3 keV is similar to that observed in the nucleus, and is marginally consistent with the putative unobscured nuclear spectrum."
+ It is worthwhile to cousider whether this component could be electron scattered uuclear flux., It is worthwhile to consider whether this component could be electron scattered nuclear flux.
+ The scattered X-ray. luminosity. Li. is giveu as where Li is the intrinsic nuclear bluminosity. € is the solid augle subteuded by the scatterer as viewed from the nucleus. aud Feat is the optical depth through the scattering gas.," The scattered X-ray luminosity, $L_{\rm scat}$, is given as where $L_{\rm int}$ is the intrinsic nuclear luminosity, $\Omega$ is the solid angle subtended by the scatterer as viewed from the nucleus, and $\tau_{\rm scat}$ is the optical depth through the scattering gas."
+" The unabsorbect 2-10 keV nuclear luminosity is in the range 1.2x104 to L.7x10% ere ! (Mattetal.1999).. aud the 2-10 keV luminosity of the scattered emission. alter correcting for absorption aud the contribution from nuclear light scattered by the telescope mirrors (Section 3.3.1)). is c7x10°"" erg s/t."," The unabsorbed 2–10 keV nuclear luminosity is in the range $1.2
+\times 10^{41}$ to $1.7 \times 10^{42}$ erg $^{-1}$ \citep{mat99}, and the 2–10 keV luminosity of the scattered emission, after correcting for absorption and the contribution from nuclear light scattered by the telescope mirrors (Section \ref{sec-nw}) ), is $\simeq
+7 \times 10^{37}$ erg $^{-1}$."
+ Assuming a hall-opening angle of 0~15° for the ionization cone. the solid angle subtended by the scatterer is O/la=4(1—cos0)0.15. so the column density of electrons through the ⋅ ⋅⋅ ⊳∖∢∙⋜↕⊓≺↵↥⋅⋯∑≟∑∸⋜↕⊳∖↥⊳∖∐⊔∐↩↕⋅⋜↕∐∑≟↩↓≍∐⊢∎∎↕∩−≻≍⋯−↓∢∙⋯−⋅⊺↥↓⊳∖↓⊳∖⊳∖⋯⋜↕∐≺↵↓⋅↕∐⋜," Assuming a half-opening angle of $\theta \simeq 45^{\circ}$ for the ionization cone, the solid angle subtended by the scatterer is $\Omega/4\pi = \frac{1}{2} (1 - \cos
+\theta) = 0.15$, so the column density of electrons through the scattering gas is in the range $4 \times 10^{20}$ to $2 \times
+10^{21}$ $^{-2}$."
+↕∐⋜↕∐⊽∖⇁⊔∐⊔∐⊳∖∐∙⋜↕∣≻⊳∖∩↥⋅∣≻∐∑≟ ↽⋅↽≻ iaqq» 2 ⋅⋅ ⋅⋅⋅ ⋅ ∢∙∩∐⊔∐∐⋜↕⊳∖⊳∖∩∢∙↥⋜↕↕≺↵≺⇂∖∖↽∐∐↕∐≺↵↥⋜⋃⋅∑∸↩−⊳∖∢∙⋜↕↥≺↵≺↵∐∐⊳∖⊳∖↥∩∐↕∩↕∐≺↵↥∖⊽∖↖⊽↸⋮⊺⋜↕∣≻↥≺↵∶⋝⋝⋝⋜⋯≼⊔↥⋯⊳∖≺⇂∩≺↵⊳∖∐∩↕↕⋅≺↲≺↽↓⇂∐⋅≺↵ ↕∐≺↵⊳∖∢↝⋜⋯≺↵↕⋅↥∐∑∸∑∸⋜≹⊳∖↕∩∣⋈↵∖⇁≺↵↕⋅⊽∖⊽∐↕∑⇁⋡⊔⊽∖⇁↕∩∐↕∠≺↵≺⊔⋅∩↕⋅↕, This is smaller than any intrinsic absorbing column associated with the large-scale emission to the NW (Table \ref{tbl-3}) ) and thus does not require the scattering gas to be very highly ionized for the scattered X-rays to escape.
+"∐≺↵⊳∖∢↝⋜≹⊓≺↵⋅≺↵≺∟∖≶−↕⋅⋜↕⊽∖⊽⊳∖↕∩≺↵⊳∖∢↝⋜↕↥↽≻≺↵− ⊺∐≺↵↕∐≺↵⊔∐⋜↕↥∢∙∩⋯↥↽≻∩∐↩∐↕↥∐↕∐≺↵⊳∖↥↽≻≺↵∢∙⋃⋅⋃⋯∩↥∎↕∐≺↵↓∖⊽∖↾∖⊽≺↵⊸∖↕≺↵∐⊳∖↥∩∐∐⋜↕⊳∖⋜↕↕≺↲⋯↥↽≻≺↵↕⋅⋜⋯⊔⋅≺↲∩↥∎⊺≃⊤⋊⊁∪∪ Ix- ancl an emission.measure ofJ )2V.ςOrze65aE1093 2em. which. ΙΕ⋅ 0.15ye “assuming2 ""[n] a volume Vo—28x10 ci? for the emitting region."," The thermal component in the spectrum of the NW extension has a temperature of $T \simeq 7 \times 10^{6}$ K and an emissionmeasure of $n_{\rm e}^{2} V \simeq 6 \times 10^{61}$ $^{-3}$, which yields $n_{\rm e} \simeq 0.15$ $^{-3}$ assuming a volume $V = 2.8 \times
+10^{63}$ $^{3}$ for the emitting region."
+ The pressure and cooling time of the X-ray. emitting eas are then nATc1.5x10.I! erg cm? and nKTV/LscTx10* years. respectively. assuming a 0.5-10 keV luminosity. after correcting for absorption. of Ly~2.0x10 erg !.," The pressure and cooling time of the X-ray emitting gas are then $n_{\rm e} k T \simeq 1.5
+\times 10^{-10}$ erg $^{-3}$ and $n_{\rm e} k T V / L_{\rm X} \simeq
+7 \times 10^{7}$ years, respectively, assuming a 0.5–10 keV luminosity, after correcting for absorption, of $L_{\rm X} \simeq 2.0
+\times 10^{38}$ erg $^{-1}$."
+ The velocity eradients across the [Orr] emission line images are of order 100-200 kins + 1997).. indicating that the gas is too hot to have been heated in situ by shocks with velocities iu this rauge.," The velocity gradients across the ] emission line images are of order $100$ $200$ km $^{-1}$ \citep{vb97}, indicating that the gas is too hot to have been heated in situ by shocks with velocities in this range."
+ Thus. a plausible explanation is that the N-ray emitting gas is associated with a wind outflowing Grom the uucleus.," Thus, a plausible explanation is that the X-ray emitting gas is associated with a wind outflowing from the nucleus."
+ The alteruative is that tlie X-ray οπο gas is by the uucleus aud thus mostly emission lines from a much cooler plasma., The alternative is that the X-ray emitting gas is photo-ionized by the nucleus and thus mostly emission lines from a much cooler plasma.
+ Higher spectral resolution observations are needed to distinguish these two quite different moclels., Higher spectral resolution observations are needed to distinguish these two quite different models.
+ The compact X-ray sources within 10” (190 pc) of the nucleus are most likely to be X-ray binary systelus or superuova remnants within the starburst riugs of the Circinus galaxy., The compact X-ray sources within $10^{\prime\prime}$ (190 pc) of the nucleus are most likely to be X-ray binary systems or supernova remnants within the starburst rings of the Circinus galaxy.
+ The measured Colt deusities for at least tree of the closest objects (sources C. E. aud (4) are abovethe Cralactie value. presumably because of the large amount of molecular gas associated with the starburst disk.," The measured column densities for at least three of the closest objects (sources C, E, and G) are abovethe Galactic value, presumably because of the large amount of molecular gas associated with the starburst disk."
+ The association of the other sources with the Circinus galaxy is less certain., The association of the other sources with the Circinus galaxy is less certain.
+" Source F is close to. but not coincident with. a patch of Ha emission ~27"" south of the nucleus. possibly an H region iu a spiral arm of the galaxy. (Elmouttieetal. 1998).."," Source F is close to, but not coincident with, a patch of $\alpha$ emission $\simeq 27^{\prime\prime}$ south of the nucleus, possibly an H region in a spiral arm of the galaxy \citep{elm98}. ."
+ This X-ray source las a bard X-ray spectrum aud, This X-ray source has a hard X-ray spectrum and
+Figure 9aa shows two 74/330 MHz spectral index contours (traced at au=—0.05 and TO= —0.25) enclosing the east rim where we found regions with very flat spectrum 1n IC 443 superposed onto the NIR detected in the / (1.25 p n). H (1.65 jm). and Κι (2.17 um) bands as taken from the Two Micron All Sky Survey (2MASS) (?)..,"Figure \ref{ir-radio}a a shows two 74/330 MHz spectral index contours (traced at $\alpha_{74}^{330}=-0.05$ and $\alpha_{74}^{330}=-0.25$ ) enclosing the east rim where we found regions with very flat spectrum in IC 443 superposed onto the NIR detected in the J (1.25 $\mu$ m), H (1.65 $\mu$ m), and $_{s}$ (2.17 $\mu$ m) bands as taken from the Two Micron All Sky Survey (2MASS) \citep{rho01}."
+ The 2MASS image of IC 443 shows the dramatic contrast in near infrared color betwee1 the east rim and the southern porticon of the remnant., The 2MASS image of IC 443 shows the dramatic contrast in near infrared color between the east rim and the southern portion of the remnant.
+ In the color representation of the infrared emission. blue traces the J--band flux. while the infrared data in the H and Κι are shown in gree1 and red. respectively: white thus erlightens," In the color representation of the infrared emission, blue traces the -band flux, while the infrared data in the H and $_{s}$ are shown in green and red, respectively; white thus enlightens"
+Jellery&Saigo(2006) recently reviewed. the excitation of pulsations in certain low-mass stars bw so-called: “be-bump instability.,\citet{Jef06} recently reviewed the excitation of pulsations in certain low-mass stars by so-called ``Fe-bump instability”.
+ This instability is caused by a significant contribution to stellar opacity. from. Al-shell electrons. in iron-group elements at. temperatures around. Why (Rogers&Iglesias1992:TheOpacityProjectTeam1995).," This instability is caused by a significant contribution to stellar opacity from M-shell electrons in iron-group elements at temperatures around K \citep{OPAL92, OP95}."
+. lt provides the driving mechanism for radial pulsations in certain extreme helium stars (Saio1993). and. p-moce oscillations in. hot subdwart B stars. (EECT14026. variables: after prototvpe 1140262647 = να) (Charpinetetal.1996.1997:Wilkenny1997 ).," It provides the driving mechanism for radial pulsations in certain extreme helium stars \citep{Sai93} and p-mode oscillations in hot subdwarf B stars (EC14026 variables: after prototype 14026–2647 $\equiv$ Hya) \citep{Cha96,Cha97,Kil97}."
+. The Fe-bump opacity mechanism is ellective in these stars because of an increase in the contrast between opacity due to iron-group elements and opacity from other sources., The Fe-bump opacity mechanism is effective in these stars because of an increase in the contrast between opacity due to iron-group elements and opacity from other sources.
+ In the case of extreme helium. stars. the background hvdrogen opacity is suppressed.," In the case of extreme helium stars, the background hydrogen opacity is suppressed."
+ In the case of EC'14026 stars. Jellerv&Saio(2006) demonstrated that. for racial and non-racdial p-moce oscillations. the observed boundaries of the instability strip can be explained only by. increasing the iron. abundance (Yp.) without increasing the heavy element. abundance (Z) as a whole.," In the case of EC14026 stars, \citet{Jef06} demonstrated that, for radial and non-radial p-mode oscillations, the observed boundaries of the instability strip can be explained only by increasing the iron abundance $X_{\rm Fe}$ ) without increasing the heavy element abundance $Z$ ) as a whole."
+ lt was already well known that raciative forces act cülferentiallv on the ions in the stellar envelope such that substantial chemical gradients are established over a diffusion time scale 107v. (Michaud 1989).., It was already well known that radiative forces act differentially on the ions in the stellar envelope such that substantial chemical gradients are established over a diffusion time scale $\sim 10^5$ y \citep{Mic89}.
+ The consequent levitation and accumulation of iron in lavers at around: Wh (Chaver.Fontaine&Wesemael1995). enhances the Ee-bump opacity sullicientlv to excite pulsations in about of sdB stars within the EC14026 instability zone (Charpinetetal.2001)., The consequent levitation and accumulation of iron in layers at around K \citep{Cha95} enhances the Fe-bump opacity sufficiently to excite pulsations in about of sdB stars within the EC14026 instability zone \citep{Cha01}.
+. The cliscovery of oscillations with periods of a few hours in many cool πα stars (PCGITIG variables: alter 11716]426) has presented a challenge to stellar pulsation theory (Cireenctal.2003)., The discovery of oscillations with periods of a few hours in many cool sdB stars (PG1716 variables: after 1716+426) has presented a challenge to stellar pulsation theory \citep{Gre03}.
+. While the raciatively- diffusion of iron can still operate in these stars. p-modes were reported to be stable in the chemically stratified models (Charpinetctal.2001).," While the radiatively-driven diffusion of iron can still operate in these stars, p-modes were reported to be stable in the chemically stratified models \citep{Cha01}."
+. On the other hand. non-radial g-mocdes of high radial order (A2 10) and hish spherical degree (/27 3) were found to be unstable. but only in mocels of stars cooler than those in which variability has been detected (Fontaineetal.2003).," On the other hand, non-radial g-modes of high radial order $k\geq10$ ) and high spherical degree $l\geq3$ ) were found to be unstable, but only in models of stars cooler than those in which variability has been detected \citep{Fon03}."
+. Jellery&Saio(2006) found that. for homogeneous models of blue horizontal branch stars. non-raclial e-mocdoes of high radial order (A2 10) and low spherical degree (/> 2) could be excited even without any enhancement of the iron abundance. but only in a narrow range of effective temperature around. Kh. Ehe width of the g-mocde instability strip and the number of excited modes increases substantially with increasing iron abundance. but Jelfery& were unable to obtain unstable e-mocles for digl524000 KI. Even these modes were only obtained with iron enhanced by a factor of twenty over à. base," \citet{Jef06} found that, for homogeneous models of blue horizontal branch stars, non-radial g-modes of high radial order $k\geq10$ ) and low spherical degree $l\ge2$ ) could be excited even without any enhancement of the iron abundance, but only in a narrow range of effective temperature around K. The width of the g-mode instability strip and the number of excited modes increases substantially with increasing iron abundance, but \citet{Jef06} were unable to obtain unstable g-modes for $T_{\rm eff}>24\,000$ K. Even these modes were only obtained with iron enhanced by a factor of twenty over a base"
+IINC aud IICN in comet Iale-Dopp (Bockeléc-A\Lorvanetal.2005:Bockeléce-Morvau 2007).. which show that the IINC/IICN visibility ratio is almost flat as a function of baseline. aud equal to the sinegle-dish flux ratio.,"HNC and HCN in comet Hale-Bopp \citep{bockel05, bockel07}, which show that the HNC/HCN visibility ratio is almost flat as a function of baseline, and equal to the single-dish flux ratio."
+ This sugeests that. contrary to model predictious. IINC in colmct Tale-Bopp was produced. efficieutlv iu the mner conia (parent scale-leneth «2000 kin).," This suggests that, contrary to model predictions, HNC in comet Hale-Bopp was produced efficiently in the inner coma (parent scale-length $\ll 2000$ km)."
+" Comet 73P/Sclavassinanu-Wachinaii is a shiort-period (5.31. vears) comet of the Jupiter family that likely originated frou, the trans-Neptuuian population in the outer Solar Svstem.", Comet 73P/Schwassmann-Wachmann is a short-period (5.34 years) comet of the Jupiter family that likely originated from the trans-Neptunian population in the outer Solar System.
+ Narrow-hband photometry aud spectroscopic observations during the 1990 apparition showed that this comet is strouglv depleted im Co and Πο relative to CN and ΟΠ. placing it in the 21P/Caacobini-Zinner class of odd comets with extreme depletions (AIIearnetal.1995:Fink&Ticks1996).," Narrow-band photometry and spectroscopic observations during the 1990 apparition showed that this comet is strongly depleted in $_2$ and $_2$ relative to CN and OH, placing it in the 21P/Giacobini-Zinner class of odd comets with extreme depletions \citep{ahearn95,fink96}."
+ During its 1995 perihelion passage. comet το) split into several fragments (Bochuhardtetal. 2002).," During its 1995 perihelion passage, comet 73P split into several fragments \citep{boeh02}."
+. Its 2006 apparition thus offered an exceptional opportunity to measure the IINC/IICN ratio iu fragments of a Jupiter-£uuilv comet of deviant composition., Its 2006 apparition thus offered an exceptional opportunity to measure the HNC/HCN ratio in fragments of a Jupiter-family comet of deviant composition.
+ Iudecd. colct 73P made a very close approach to Earth in May 2006 (at 0.067 and 0.079. AU for Fragments B aud C. respectively). that compensated for its moderate activity (water production rate of a few < 1075 s+ at poribeliou) and allowing scusitive investigations.," Indeed, comet 73P made a very close approach to Earth in May 2006 (at 0.067 and 0.079 AU for Fragments B and C, respectively), that compensated for its moderate activity (water production rate of a few $\times$ $^{28}$ $^{-1}$ at perihelion) and allowing sensitive investigations."
+ Iu the present paper. we compile the existing measurements of the IENC'/IICN ratio in moderately active comets. incliding our sensitive upper ιτ obtained iun comet Το/Schwassimauu-Wachinaun and discuss possible nuplicatious for the origin of INC in coluctary atimospheres.," In the present paper, we compile the existing measurements of the HNC/HCN ratio in moderately active comets, including our sensitive upper limit obtained in comet 73P/Schwassmann-Wachmann and discuss possible implications for the origin of HNC in cometary atmospheres."
+ Observations of Comet 73P/Sclwassimann-Waclinan (Fragment. B) presented here were carried out on Alay 911. 2006 UT using the 10.bin Leighton telescope of Caltech Subiuillnieter: Observatory (C80) on Mauna Ίνοα in Wawaii.," Observations of Comet 73P/Schwassmann-Wachmann (Fragment B) presented here were carried out on May 9–11, 2006 UT using the 10.4-m Leighton telescope of Caltech Submillimeter Observatory (CSO) on Mauna Kea in Hawaii."
+ The heliocentric and gcocentiic distances of the comet at the time of the CSO observations were 1.02 aud 0.072 AU. respectively.," The heliocentric and geocentric distances of the comet at the time of the CSO observations were 1.02 and 0.072 AU, respectively."
+ We used the 315 GIIz acilitv SIS receiver and acousto-optical spectrometers with total bandwidths of 50 MIIZ aud 1 Giz., We used the 345 GHz facility SIS receiver and acousto-optical spectrometers with total bandwidths of 50 MHz and 1 GHz.
+ The requency scale of the spectrometers (the reference channel aud the channel width) was established. by injecting calibration signals from LO MITz aud 100 ΛΙΠΗ. requeucyv comb generators. for the 50 MITz and 1 GIIz AOS. respectively.," The frequency scale of the spectrometers (the reference channel and the channel width) was established by injecting calibration signals from 10 MHz and 100 MHz frequency comb generators, for the 50 MHz and 1 GHz AOS, respectively."
+ The effective frequency resolution of he high-resolution (50 MIIz baudwidth) spectrometer was 95 kIIz (2 channels}. correspoucding to ~0.08-," The effective frequency resolution of the high-resolution (50 MHz bandwidth) spectrometer was 95 kHz (2 channels), corresponding to 0.08."
+" Pointing of the telescope. checked by performing five-x»""t coutimmiun scaus of Jupiter iu the south. was deteriuiued to dift by oon time-scales of ~3 hows at the tine of the observations (one fifth of the FEWIIM. CSO beam at 360 1)."," Pointing of the telescope, checked by performing five-point continuum scans of Jupiter in the south, was determined to drift by on time-scales of 3 hours at the time of the observations (one fifth of the FWHM CSO beam at 360 GHz)."
+" The accuracy of the cometary eplicmers. checked by performing five-point scaus of the IICN 7-123 transition iu the comet, was determined to be9""."," The accuracy of the cometary ephemeris, checked by performing five-point scans of the HCN $J$ =4–3 transition in the comet, was determined to be."
+. The nuin beam efficieucy. determined from total-power observatious of Jupiter. Saturn aud Mars. was measured to be—TÍÁ.," The main beam efficiency, determined from total-power observations of Jupiter, Saturn and Mars, was measured to be."
+. The absolute calibration wncertaimty of the individual mcasurelents is~20%., The absolute calibration uncertainty of the individual measurements is.
+. However. since the TICN aud IINC lines were observed almost simultaneously. with the sale receiver. we estimate the calibration uncertainty of the INC/TICN line ratio to be sinaller. of order.," However, since the HCN and HNC lines were observed almost simultaneously, with the same receiver, we estimate the calibration uncertainty of the HNC/HCN line ratio to be smaller, of order."
+. IICN J=13 spectra (rest frequeney 351.50517 GIIz) taken ou Max 911. UT are shown in Figue 1 aud the line intensitics are listed a Table 1.. along with the correspoucing statistical uncertainties. as determined from the noise level measured in the spectra.," HCN $J$ =4–3 spectra (rest frequency 354.50547 GHz) taken on May 9–11 UT are shown in Figure \ref{fig_hcn} and the line intensities are listed in Table \ref{tab_int}, along with the corresponding statistical uncertainties, as determined from the noise level measured in the spectra."
+ A siguificaut decrease iu the ΠΟΝ iuteusitv (bv a factor of 1.8) is seen over the period of our observations. following au earlier outburst.," A significant decrease in the HCN intensity (by a factor of 1.8) is seen over the period of our observations, following an earlier outburst."
+ The low-level pedestal. seen most clearly im the Mav 9 spectrum. is due to the IICN lypertine structure (c.e@.. Lisetal. L997).," The low-level pedestal, seen most clearly in the May 9 spectrum, is due to the HCN hyperfine structure (e.g., \citealt{lis97}) )."
+ TING J=13 spectra Gest frequency 362.63031 GIIZ) taken on Mav 1011 UT are shown in Figure 2.., HNC $J$ =4–3 spectra (rest frequency 362.63031 GHz) taken on May 10–11 UT are shown in Figure \ref{fig_hnc}.
+ The UNC aud IICN observations on these two nights were interleaved iu time to account for the time variability of the outeassine., The HNC and HCN observations on these two nights were interleaved in time to account for the time variability of the outgassing.
+ No Ην ciission is detected (Table 13)., No HNC emission is detected (Table \ref{tab_int}) ).
+ Fragiieut D of comet 73P underwent several outbursts, Fragment B of comet 73P underwent several outbursts
+observable.,observable.
+ The contents of this section are ‘subject to change since much of our mformation is complete., The contents of this section are `subject to change' since much of our information is incomplete.
+ The reason for including this source list is to provide an overview., The reason for including this source list is to provide an overview.
+ Some aspects of the classification are subjective and we anticipate changes as new data become available., Some aspects of the classification are subjective and we anticipate changes as new data become available.
+ We /inaiutain current information at/., We maintain current information at.
+. Table 3 contains a list of bona fide N-ray jet sources kuown to us as of 2001 March., Table \ref{tab:sources} contains a list of bona fide X-ray jet sources known to us as of 2001 March.
+ Iu the last column is the classification which is described here., In the last column is the classification which is described here.
+ This category is perhaps more secure than the others because the data confirm the predictions rather well., This category is perhaps more secure than the others because the data confirm the predictions rather well.
+ All | SSC sources (sec. 2?)), All 4 SSC sources (sec. \ref{sec:ssc}) )
+ are ERII radio galaxies aud their terminal hotspots rather than brightuess cuhancements iu the jets are the SSC cmitters., are FRII radio galaxies and their terminal hotspots rather than brightness enhancements in the jets are the SSC emitters.
+" For all of these. the maenetic field streneths iu the hotspots are in the range 100 to LOO μέ, values consistent with the equipartition field for the case of little or no contribution to the particle energy density from protous."," For all of these, the magnetic field strengths in the hotspots are in the range 100 to 400 $\mu$ G, values consistent with the equipartition field for the case of little or no contribution to the particle energy density from protons."
+ Svuchirotrou DIununositfies and Lifetimes are unaffected by the IC losses and the electrous responsible for the observed X-rays are those enüttiug svuchrotrou chussion in the normal radio band., Synchrotron luminosities and lifetimes are unaffected by the IC losses and the electrons responsible for the observed X-rays are those emitting synchrotron emission in the normal radio band.
+ Alternative models for these hotspots are the geucratio- of high energev electrons via the PIC (Proton Tuduced Cascade. \Tannheiu et al.," Alternative models for these hotspots are the generation of high energy electrons via the PIC (Proton Induced Cascade, Mannheim et al."
+ 1991) or proto- svuchrotrou (Aharouian 2001) processes. both of which would require πιο strouger magnetic fields (QDx 5004).," 1991) or proton synchrotron (Aharonian 2001) processes, both of which would require much stronger magnetic fields $B\,\gg\,500\mu$ G)."
+ Aliudful of ougoiug debate about the evidence for spectral eutoffs from optical data (sec. ?7)).," Mindful of ongoing debate about the evidence for spectral cutoffs from optical data (sec. \ref{sec:sync}) ),"
+ we classify MBs? knots D. A. and D as svuchrotrou cluitters (Diretta. Stern. and Warris. 1991: Marshall et al.," we classify M87 knots D, A, and B as synchrotron emitters (Biretta, Stern, and Harris, 1991; Marshall et al.,"
+ 20015)., 2001b).
+ Additional sources in this category are 3€ 390.3 hotspot D (Uarris et al., Additional sources in this category are 3C 390.3 hotspot B (Harris et al.
+ 1998). knots Al and Bl in the 3C 273 jet (Alarshall ct al.," 1998), knots A1 and B1 in the 3C 273 jet (Marshall et al."
+ 200La). and 3€ 66D (ILudcastle. Birkinshaw. Worrall. 20015) although for some of these. there is the problem that a cooling break iu the cuissio- spectruui has not been isolated.," 2001a), and 3C 66B (Hardcastle, Birkinshaw, Worrall, 2001b) although for some of these, there is the problem that a cooling break in the emission spectrum has not been isolated."
+ The extension of the svuchnrotron spectrum from the optical to the Nav. has only a snall effect on the total cuerey and on the calculation of the equipartition magnetic field., The extension of the synchrotron spectrum from the optical to the X-ray has only a small effect on the total energy and on the calculation of the equipartition magnetic field.
+" Ποπονα, this model requires the exteusion of the clectrou spectrum from ~cLO to >107."," However, this model requires the extension of the electron spectrum from $\gamma~\approx~10^5$ to $\geq~10^7$."
+ The primary observable consequence ds that the electrous responsible for the N-ravs have lifetimes of order LO to 100 veurs., The primary observable consequence is that the electrons responsible for the X-rays have lifetimes of order 10 to 100 years.
+ If a somewhat “ad hoc additional spectral conrponeut is allowed (sec. 22:, If a somewhat 'ad hoc' additional spectral component is allowed (sec. \ref{sec:distinct};
+ ie. à high energy conrponeut of the electron energv distribution with a flatter spectra han that observed iu the radio/optical domain). then svuchrotron cCluission niodels can be devised for sources such as Pictor A (Wilson. Young. Shopbell 2001). section. [.2.2.2). 3€ 273 knot D. PISS 0637. and 3€ 120 (ILuris et al.," i.e. a high energy component of the electron energy distribution with a flatter spectrum than that observed in the radio/optical domain), then synchrotron emission models can be devised for sources such as Pictor A (Wilson, Young, Shopbell 2001b, section 4.2.2.2), 3C 273 knot D, PKS 0637, and 3C 120 (Harris et al."
+ 1999: where a verv hard spectral mdex of accelerated particles of p=1.7 would be required)., 1999; where a very hard spectral index of accelerated particles of $p$ =1.7 would be required).
+ The SSC sources (sec. ?7)), The SSC sources (sec. \ref{sec:ssc}) )
+ could also be accomodated by svuchrotron models if additional spectral components are allowed., could also be accomodated by synchrotron models if additional spectral components are allowed.
+ lu sections ?? aud ?7 we discuss nuuiv of the details aud ramifications of this model., In sections \ref{sec:beam} and \ref{sec:conflict} we discuss many of the details and ramifications of this model.
+ While there is little debate concerning the utility of the beanung model in its ability to explain a larger ratio of N-rav to svuchrotron emüssiou than other models. there is no completely comvincing evidence that kpce-cale. jets actually have bulk. velocities close to c. Beanune models have been prescuted for PIKS0637 (Celotti et al. 2001: Tavecchio et al.," While there is little debate concerning the utility of the beaming model in its ability to explain a larger ratio of X-ray to synchrotron emission than other models, there is no completely convincing evidence that kpc-scale jets actually have bulk velocities close to c. Beaming models have been presented for PKS0637 (Celotti et al, 2001; Tavecchio et al."
+ 2000) aud for 3C° 273 (Saimbruna ct al., 2000) and for 3C 273 (Sambruna et al.
+ 2001)., 2001).
+ For the knots iu these jets. viewing angles are ecucrally required to be less than 107. D values range from 5 to 20. and equipartition fields are siguificautlv sunaller than for wnbeamed svuchrotron emission.," For the knots in these jets, viewing angles are generally required to be less than $^{\circ}$, $\Gamma$ values range from 5 to 20, and equipartition fields are significantly smaller than for unbeamed synchrotron emission."
+ A very significant differcuce between leamine and svuchrotron imodels is that the beaming model posits that the electrons responsible for the observed X-rays are at the verv bottoni of the electron distribution with * iu the range 20-150., A very significant difference between beaming and synchrotron models is that the beaming model posits that the electrons responsible for the observed X-rays are at the very bottom of the electron distribution with $\gamma$ in the range 20-150.
+ This implies that no variability is expected iu the X-ray iuteusitv., This implies that no variability is expected in the X-ray intensity.
+ Because there is a siguificaut difference between the electron cucreies producing, Because there is a significant difference between the electron energies producing
+The dinensiouless parameter 9 gives the rate at which 2p states are produced by captured Lyiman-a radiation in terms of the 2p creation rate [from recombination (at low densities).,The dimensionless parameter $S$ gives the rate at which $2p$ states are produced by captured $\alpha$ radiation in terms of the $2p$ creation rate from recombination (at low densities).
+ The ratio n3p/(3105) provides a determination of the relative importance of Lyman-a pumping ol the 2p states aud whether a fiue structure line is expected to appear in absorption or (stimulated)einission., The ratio $n_{2p}/(3 n_{2s})$ provides a determination of the relative importance of $\alpha$ pumping of the $2p$ states and whether a fine structure line is expected to appear in absorption or (stimulated)emission.
+" I£ ye,/(3ies)>1. then the 0.0 GHz. 284-2py/o line will appear in emission. and the 1.1 GHz. 254/5-2p4;3 line will appear iu absorption."," If $n_{2p}/(3 n_{2s}) > 1$, then the 9.9 GHz, $2s_{1/2}$ $2p_{3/2}$ line will appear in emission, and the 1.1 GHz, $2s_{1/2}$ $2p_{1/2}$ line will appear in absorption."
+ Of course. Nap)/(Bies)1 unplies the opposite. (," Of course, $n_{2p}/(3 n_{2s}) < 1$ implies the opposite. ("
+This assumes that the two 2p states are populated accordiug to their relative statistical weights.,This assumes that the two $2p$ states are populated according to their relative statistical weights.
+ To evaluate cases in which ορ(οκ) is of order unity. it would be necessary to separately account for the rates at which 2p4;51/2 aid 2p4;53/2 states are created and destroved. iucludiugo the collisional rates Coupling these states.)," To evaluate cases in which $n_{2p}/(3 n_{2s})$ is of order unity, it would be necessary to separately account for the rates at which $2p_{1/2}$ and $2p_{3/2}$ states are created and destroyed, including the collisional rates coupling these states.)"
+ From equatious (11) and (12) it follows that 5/0324)>1$ >S554 where. (The numerical value of 5.4; Corresponds to f= 1/3.)," From equations (11) and (12) it follows that $n_{2p}/(3 n_{2s}) > 1$ if $S > S_{crit}$ where, (The numerical value of $S_{crit}$ corresponds to $f = 1/3$ .)"
+ Because the 2p states naturally decay about 105 timesfaster than the 25 states. population equality thus requires a puiplug rate some ἃ orders of magultude faster than the approximate rate at which 2s aud 2p states are formed through recombination.," Because the $2p$ states naturally decay about $10^8$ timesfaster than the $2s$ states, population equality thus requires a pumping rate some 8 orders of magnitude faster than the approximate rate at which $2s$ and $2p$ states are formed through recombination."
+ Determination of the Lyiman-a deusity in ΠΠ regions is a complicated trausfer problem., Determination of the $\alpha$ density in HII regions is a complicated transfer problem.
+ It is ikely. however. that the dominant mechauisim for removing Lyiman-a photous is quite straightforward. i.e. absorption by dust (Ixaplan&Pikeluer1970:Spitzer1978).," It is likely, however, that the dominant mechanism for removing $\alpha$ photons is quite straightforward, i.e. absorption by dust \citep{kap70,spi78}."
+". Thus. we eschew the noncoherent ""adiative transfer problem aud find the upper limits to {ρω and ο set by absorption."," Thus, we eschew the noncoherent radiative transfer problem and find the upper limits to $t_{Ly\alpha}$ and $S$ set by absorption."
+ Other competing removal processes would reduce the lifetime. aud therefore density. of Lyiman-a photons. 'esultiug in a lower value for 5.," Other competing removal processes would reduce the lifetime, and therefore density, of $\alpha$ photons, resulting in a lower value for $S$."
+ Using a silicate-graphite model for dust in HII regions (Aannestad1989).. with a cust-to-gas atio of 0.009. the extinction at Lyman-a can be shown to be [Nj/0.ExLOcm2) mag. where Vy) is the columu deusity of hydrogeu.," Using a silicate-graphite model for dust in HII regions \citep{aan89}, with a dust-to-gas ratio of 0.009, the extinction at $\alpha$ can be shown to be $N_H/(5.4 \times
+10^{20} {\rm cm}^{-2})$ mag, where $N_H$ is the column density of hydrogen."
+ Thealbedo lor this mixture is about 0.1 at Lyman-a [).., Thealbedo for this mixture is about 0.4 at $\alpha$ \citep{dra84}. .
+ The lifetime of Lyinan-a photous against absorption by dust cau then be calculated as py=(3.3x10!em3 s)/ng.," The lifetime of $\alpha$ photons against absorption by dust can then be calculated as $t_{Ly\alpha} = (3.3 \times 10^{10}\ 
+{\rm cm}^{-3}\ {\rm s}) / n_H$ ."
+ We then fiud that Since X<<I throughout most of the volume of au HII region (Osterbrock LIS9).. we couclude," We then find that Since $\chi << 1$ throughout most of the volume of an HII region \citep{ost89}, , we conclude"
+ol thickness negligible in respect to the main dimension.,of thickness negligible in respect to the main dimension.
+ A tvpical dimension of the laver can be found as follows., A typical dimension of the layer can be found as follows.
+ The averaged observed diameter of the galaxies is where De?=45004 corresponds to the extension of the maximum void visible on ihe CFA® slices., The averaged observed diameter of the galaxies is where $D_{max}^{obs}=4500~\frac{Km}{sec}$ corresponds to the extension of the maximum void visible on the CFA2 slices.
+" In the framework of the theory of the primordial explosions .see Charlton& and Zanine(ti&Ferraro(1990).. this means that the mean observed area of a bubble οὐ, is The averaged area of a face of Voronoi polvhedra. Ay. is where Vo is (he averaged number of regular [aces of the Voronoi polvhedra. i.e. Np =16 see Okabeetal.(1992):Zaninetti(2006)."," In the framework of the theory of the primordial explosions ,see \citet{charlton} and \citet{ferraro}, this means that the mean observed area of a bubble $\overline{A^{obs}} $, is The averaged area of a face of Voronoi polyhedra, $\overline{A_V} $, is where ${\overline{N_F}}$ is the averaged number of irregular faces of the Voronoi polyhedra, i.e. ${\overline{N_F}}$ =16, see \citet{okabe,Zaninetti2006}."
+". The averaged side of a [ace of a irregular polvhedron . £4:. is The thickness of the laver. ὃ can be derived [rom the shock theory. see (1984).. and is 1/12 of the radius of the advancing shock The number of galaxies in (his typical laver. ""Nc; can be Found by multiplving n,e0.1 . the densily of galaxies. for the volume of the cube of side 12 Ape: i.e. Ne;2112."," The averaged side of a face of a irregular polyhedron , $L_V$, is The thickness of the layer, $\delta$, can be derived from the shock theory, see \cite{deeming}, and is 1/12 of the radius of the advancing shock, The number of galaxies in this typical layer, $N_G$, can be found by multiplying $n_*\approx 0.1$ , the density of galaxies, for the volume of the cube of side 12 $Mpc$: i.e. $N_G \approx 172 $."
+ The more common way to insert the seeds of the Voronoi polvgons is a random secquence in the X and Y. direction. see Figure 1..," The more common way to insert the seeds of the Voronoi polygons is a random sequence in the $X$ and $Y$ direction, see Figure \ref{area2D_random}."
+ The distribution of the area of the irregular Voronoi polvgons is fitted. with a Αίας [unetion. see formula - (A5)) in Áppendix A..: and the captionsof Figure 1.also report the value of e as deduced [rom the parameters ol the sample.," The distribution of the area of the irregular Voronoi polygons is fitted with a Kiang function, see formula \ref{kiang}) ) in Appendix \ref{appendixa}, ,: and the captionsof Figure \ref{area2D_random} also report the value of $c$ as deduced from the parameters of the sample."
+(he analvsis. and 4 explains the implications of (hese restuts.,"the analysis, and 4 explains the implications of these results."
+ Finally. 5 summarizes our conclusions.," Finally, 5 summarizes our conclusions."
+ We obtained observations of the M31 bulge with the ACIS-I camera on 2004+ (OhslID 4681). 2004-Mav-23. (ObsID 4632). and 2004-Julv-1T. (ΟΡΙΟ 4719).," We obtained observations of the M31 bulge with the ACIS-I camera on 2004-January-31 (ObsID 4681), 2004-May-23 (ObsID 4682), and 2004-July-17 (ObsID 4719)."
+ Observations 4681 and 4682 were Guaranteed Time Observations supplied by ο. \hirray., Observations 4681 and 4682 were Guaranteed Time Observations supplied by S. Murray.
+" All of the data were obtained in ""alternating readout mode"" which reduces event pileup for bright sources but lowers the effective exposure by ~20%.", All of the data were obtained in “alternating readout mode” which reduces event pileup for bright sources but lowers the effective exposure by $\sim20\%$.
+ Observation 4681 was observed centered. on R.A.=00:42:44.4. Dec.—41:16:08.3. with a roll angle of 305.6 degrees for an effective exposure Gme of 4.09 ks.," Observation 4681 was observed centered on R.A.=00:42:44.4, Dec.=41:16:08.3 with a roll angle of 305.6 degrees for an effective exposure time of 4.09 ks."
+ Observation 4682 was centered on the same coordinates wilh a roll angle of 80.0 degrees for an effective exposure time of 3.95 ks. aud observation 4719 was observed centered on R.À.—00:42:44.3. Dec.—41:16:08.4 with a roll angle of 116.8 degrees lor an elleclive exposure time of 4.12 ks.," Observation 4682 was centered on the same coordinates with a roll angle of 80.0 degrees for an effective exposure time of 3.95 ks, and observation 4719 was observed centered on R.A.=00:42:44.3, Dec.=41:16:08.4 with a roll angle of 116.8 degrees for an effective exposure time of 4.12 ks."
+ These observations were all reduced. using the software package CIAO v3.1 with the CALDD v2.28., These observations were all reduced using the software package CIAO v3.1 with the CALDB v2.28.
+" We created exposure maps lor (the images using (the task and we found ancl measured positions aud fluxes of the sources in the image using the CIAO task Each data set detected sources down to (0.3.LO keV) fluxes of 6x10. "" photons ? Lor 0.310 keV luminosities of ~10™ erg ! for a twpical X-ray. binary svstem in M3.", We created exposure maps for the images using the task and we found and measured positions and fluxes of the sources in the image using the CIAO task Each data set detected sources down to (0.3–10 keV) fluxes of $\sim$ $\times$ $^{-6}$ photons $^{-2}$ $^{-1}$ or 0.3–10 keV luminosities of $\sim$ $^{36}$ erg $^{-1}$ for a typical X-ray binary system in M31.
+ We aligned the coordinate svstem of the N-rav images with the optical images of the Local Group Survey (LGS: Masseyetal. 2001))., We aligned the coordinate system of the X-ray images with the optical images of the Local Group Survey (LGS; \citealp{massey2001}) ).
+" These images have an assigned J2000 (FINS) world coordinate svstem accurate to 0.25"". and they provided the standard coordinate svslem (o which we aligned all of our data for this project."," These images have an assigned J2000 (FK5) world coordinate system accurate to $\sim$ $''$, and they provided the standard coordinate system to which we aligned all of our data for this project."
+ The positions of X-ray sources with known globular cluster counterparts were aligned with the elobular cluster centers in the LGS images using the taskeemap. allowing adjustments of the pixel scale as well," The positions of X-ray sources with known globular cluster counterparts were aligned with the globular cluster centers in the LGS images using the task, allowing adjustments of the pixel scale as well"
+The mean value ancl the one sigma error on the on the retrieved. values P are shown in Fig δ..,The mean value and the one sigma error on the on the retrieved values $\Gamma$ are shown in Fig \ref{recsph}.
+ This figure shows that even with only ten or so galaxies for which the polarization hase been measured. it is possible to fairly accurately determine the function P(r).," This figure shows that even with only ten or so galaxies for which the polarization hase been measured, it is possible to fairly accurately determine the function $\Gamma(r)$."
+ IH is also important to point out that if one uses two galaxies at similar radial distances from the centre. it is possible to test the assumption of spherical svnimictry.," It is also important to point out that if one uses two galaxies at similar radial distances from the centre, it is possible to test the assumption of spherical symmetry."
+ In this section we cliscuss the problem of determining the lens parameters when no assumptions are made about the lens., In this section we discuss the problem of determining the lens parameters when no assumptions are made about the lens.
+ Lf we no longer assume that the lens is spherical. then the angle 8 is not known a priori.," If we no longer assume that the lens is spherical, then the angle $\theta$ is not known a priori."
+" The svstem of equations (10)).(11)) and (13)) then consists of 3 equations for 4 unknowns (/,.δές. 0). and so it is impossible to solve for the unknowns."," The system of equations \ref{sr1}) \ref{sr2}) ) and \ref{sr4}) ) then consists of $3$ equations for 4 unknowns $l_s,\Delta l_s, \Gamma, \theta$ ), and so it is impossible to solve for the unknowns."
+ However if we assume that there are two galaxies in the sample that canellectively be associated with the same values of E and 8. then the previous svstem written for both galaxies will contain 6 unknowns. GL.BR.INTELANDp 8). and six equations. which means that it can be solved exactly for E. ϐ and the source galaxies’ parameters.," However if we assume that there are two galaxies in the sample that caneffectively be associated with the same values of $\Gamma$ and $\theta$ , then the previous system written for both galaxies will contain $6$ unknowns, $(l_{s}^{1}, l_s^{2}, \Delta
+l_{s}^{1}, \Delta l_{s}^{2}, \Gamma, \theta)$ , and six equations, which means that it can be solved exactly for $\Gamma$, $\theta$ and the source galaxies' parameters."
+ One should note. however. that if the orientation of the source galaxies is not known. adding more galaxies in this way cannot close the system. as there will always be more Unknowns than independent equations: cach new galaxy. brings three unknowns and three equations. ancl so the equations are always uncderdetermined.," One should note, however, that if the orientation of the source galaxies is not known, adding more galaxies in this way cannot close the system, as there will always be more unknowns than independent equations: each new galaxy brings three unknowns and three equations, and so the equations are always underdetermined."
+ Phere are a number of methods for finding galaxies associated with the same pairs of values of (L.A).," There are a number of methods for finding galaxies associated with the same pairs of values of $(\Gamma, \theta )$."
+ Lowe make no assumptions about the lens. then one has to take two galaxies that have an angular separation much smaller than the angular diameter of the lensing cluster.," If we make no assumptions about the lens, then one has to take two galaxies that have an angular separation much smaller than the angular diameter of the lensing cluster."
+ One can then make the approximation that the potential is the same for both ogalaxies., One can then make the approximation that the potential is the same for both galaxies.
+ I£ on the other hand. we assume that the potential is spherically symmetric. then one can take ogalaxies which are well separated but at the same distance from the centre. so they have the same E and ϐ.," If on the other hand we assume that the potential is spherically symmetric, then one can take galaxies which are well separated but at the same distance from the centre, so they have the same $\Gamma$ and $\theta$."
+ Phere could of course be other svmmetries that allow one to do this., There could of course be other symmetries that allow one to do this.
+ The equation (18)) for cachgalaxy gives the following system: ∖∖⊽⇂↥∢⋅↓⋅⋖⋅∙↙∣⊥↿∖↓, The equation \ref{solr}) ) for eachgalaxy gives the following system: where $g_1$ (resp.
+⋅⋖⋅≱∖↓≻⋡∙↙∣⊐⊐∠⇂⋖⋅⊔∪↿∢⊾⊏↥⇂⋯⊔⇂⊲↓↿⊲⊓⋅⊳∖↓⋅⋖⋅↓⋜⋯⋅∠⇂↿∪↿↓↕⋖⋅∐, $g_2$ ) denote quantities related to the first (resp.
+↓⋅⊳∖↿↿∖↓⋅⋖⊾⊳∖↓≻⊳, second) galaxy.
+⊳∖∢⋅≼⇍∪⊔∠⊔⋏∙≟⋜↧↓⋜∟∖∙∖⇁⊳↾∐∐⋅↓≻⋜⊔⋅⋜⋯↓⋖⋅∩⋅↓⋅≱∖↙⊲∕↙↴∕⋰⋜⋯∠⇂↔⋤∕↿∖∕⋅⋅∶∣⊽⋡⋅↔⊽∃ ∠⇂∢⊾↓≻⋖⊾⊔∠⇂∪⊔," The parameters $C_{k}^{g_l}$ and $S_{k}^{g_l}$ $k=i,s$ ) depend only on $\theta$ ."
+↓∙∖⇁∪⊔↙−∕⋡↾↓∖⇂↥⋠↓≱∖≱∖∙∖⇁≱∖↿⋖⋅⊔⊔∙⋜⋯∣⋊⋅≱∖∪⇂∖⇁⋯⇂⇂⋅∪↓⋅↓↸, Thissystem canbe solved for $\Gamma$ and $\theta$ .
+⋜⋯∠⊓⇥⊳⇀∖⊳∖∣⋡∢⋅⇂∎∪↓⋅∢⋅⋡∖∖⊽∢⋅⊓⋅≱∖⊓⋅∠⇂↥↓↥∢⋅↓≻∪≱∖≱∖⊀↓∣⋡⊲↓↓⋠↓↿∙∖⇁∪⇂∎↓⋅⋖⋅≼∙∪⊔," As before, we tested the possibility of reconstructing lens parameters"
+⋜⋯∠⊓⇥⊳⇀∖⊳∖∣⋡∢⋅⇂∎∪↓⋅∢⋅⋡∖∖⊽∢⋅⊓⋅≱∖⊓⋅∠⇂↥↓↥∢⋅↓≻∪≱∖≱∖⊀↓∣⋡⊲↓↓⋠↓↿∙∖⇁∪⇂∎↓⋅⋖⋅≼∙∪⊔≱," As before, we tested the possibility of reconstructing lens parameters"
+⋜⋯∠⊓⇥⊳⇀∖⊳∖∣⋡∢⋅⇂∎∪↓⋅∢⋅⋡∖∖⊽∢⋅⊓⋅≱∖⊓⋅∠⇂↥↓↥∢⋅↓≻∪≱∖≱∖⊀↓∣⋡⊲↓↓⋠↓↿∙∖⇁∪⇂∎↓⋅⋖⋅≼∙∪⊔≱∖," As before, we tested the possibility of reconstructing lens parameters"
+⋜⋯∠⊓⇥⊳⇀∖⊳∖∣⋡∢⋅⇂∎∪↓⋅∢⋅⋡∖∖⊽∢⋅⊓⋅≱∖⊓⋅∠⇂↥↓↥∢⋅↓≻∪≱∖≱∖⊀↓∣⋡⊲↓↓⋠↓↿∙∖⇁∪⇂∎↓⋅⋖⋅≼∙∪⊔≱∖⇂," As before, we tested the possibility of reconstructing lens parameters"
+⋜⋯∠⊓⇥⊳⇀∖⊳∖∣⋡∢⋅⇂∎∪↓⋅∢⋅⋡∖∖⊽∢⋅⊓⋅≱∖⊓⋅∠⇂↥↓↥∢⋅↓≻∪≱∖≱∖⊀↓∣⋡⊲↓↓⋠↓↿∙∖⇁∪⇂∎↓⋅⋖⋅≼∙∪⊔≱∖⇂↓," As before, we tested the possibility of reconstructing lens parameters"
+⋜⋯∠⊓⇥⊳⇀∖⊳∖∣⋡∢⋅⇂∎∪↓⋅∢⋅⋡∖∖⊽∢⋅⊓⋅≱∖⊓⋅∠⇂↥↓↥∢⋅↓≻∪≱∖≱∖⊀↓∣⋡⊲↓↓⋠↓↿∙∖⇁∪⇂∎↓⋅⋖⋅≼∙∪⊔≱∖⇂↓⋅," As before, we tested the possibility of reconstructing lens parameters"
+⋜⋯∠⊓⇥⊳⇀∖⊳∖∣⋡∢⋅⇂∎∪↓⋅∢⋅⋡∖∖⊽∢⋅⊓⋅≱∖⊓⋅∠⇂↥↓↥∢⋅↓≻∪≱∖≱∖⊀↓∣⋡⊲↓↓⋠↓↿∙∖⇁∪⇂∎↓⋅⋖⋅≼∙∪⊔≱∖⇂↓⋅⋯," As before, we tested the possibility of reconstructing lens parameters"
+⋜⋯∠⊓⇥⊳⇀∖⊳∖∣⋡∢⋅⇂∎∪↓⋅∢⋅⋡∖∖⊽∢⋅⊓⋅≱∖⊓⋅∠⇂↥↓↥∢⋅↓≻∪≱∖≱∖⊀↓∣⋡⊲↓↓⋠↓↿∙∖⇁∪⇂∎↓⋅⋖⋅≼∙∪⊔≱∖⇂↓⋅⋯∼," As before, we tested the possibility of reconstructing lens parameters"
+⋜⋯∠⊓⇥⊳⇀∖⊳∖∣⋡∢⋅⇂∎∪↓⋅∢⋅⋡∖∖⊽∢⋅⊓⋅≱∖⊓⋅∠⇂↥↓↥∢⋅↓≻∪≱∖≱∖⊀↓∣⋡⊲↓↓⋠↓↿∙∖⇁∪⇂∎↓⋅⋖⋅≼∙∪⊔≱∖⇂↓⋅⋯∼⇂," As before, we tested the possibility of reconstructing lens parameters"
+⋜⋯∠⊓⇥⊳⇀∖⊳∖∣⋡∢⋅⇂∎∪↓⋅∢⋅⋡∖∖⊽∢⋅⊓⋅≱∖⊓⋅∠⇂↥↓↥∢⋅↓≻∪≱∖≱∖⊀↓∣⋡⊲↓↓⋠↓↿∙∖⇁∪⇂∎↓⋅⋖⋅≼∙∪⊔≱∖⇂↓⋅⋯∼⇂⊀," As before, we tested the possibility of reconstructing lens parameters"
+⋜⋯∠⊓⇥⊳⇀∖⊳∖∣⋡∢⋅⇂∎∪↓⋅∢⋅⋡∖∖⊽∢⋅⊓⋅≱∖⊓⋅∠⇂↥↓↥∢⋅↓≻∪≱∖≱∖⊀↓∣⋡⊲↓↓⋠↓↿∙∖⇁∪⇂∎↓⋅⋖⋅≼∙∪⊔≱∖⇂↓⋅⋯∼⇂⊀↓," As before, we tested the possibility of reconstructing lens parameters"
+⋜⋯∠⊓⇥⊳⇀∖⊳∖∣⋡∢⋅⇂∎∪↓⋅∢⋅⋡∖∖⊽∢⋅⊓⋅≱∖⊓⋅∠⇂↥↓↥∢⋅↓≻∪≱∖≱∖⊀↓∣⋡⊲↓↓⋠↓↿∙∖⇁∪⇂∎↓⋅⋖⋅≼∙∪⊔≱∖⇂↓⋅⋯∼⇂⊀↓⊔," As before, we tested the possibility of reconstructing lens parameters"
+⋜⋯∠⊓⇥⊳⇀∖⊳∖∣⋡∢⋅⇂∎∪↓⋅∢⋅⋡∖∖⊽∢⋅⊓⋅≱∖⊓⋅∠⇂↥↓↥∢⋅↓≻∪≱∖≱∖⊀↓∣⋡⊲↓↓⋠↓↿∙∖⇁∪⇂∎↓⋅⋖⋅≼∙∪⊔≱∖⇂↓⋅⋯∼⇂⊀↓⊔⋏," As before, we tested the possibility of reconstructing lens parameters"
+⋜⋯∠⊓⇥⊳⇀∖⊳∖∣⋡∢⋅⇂∎∪↓⋅∢⋅⋡∖∖⊽∢⋅⊓⋅≱∖⊓⋅∠⇂↥↓↥∢⋅↓≻∪≱∖≱∖⊀↓∣⋡⊲↓↓⋠↓↿∙∖⇁∪⇂∎↓⋅⋖⋅≼∙∪⊔≱∖⇂↓⋅⋯∼⇂⊀↓⊔⋏∙," As before, we tested the possibility of reconstructing lens parameters"
+⋜⋯∠⊓⇥⊳⇀∖⊳∖∣⋡∢⋅⇂∎∪↓⋅∢⋅⋡∖∖⊽∢⋅⊓⋅≱∖⊓⋅∠⇂↥↓↥∢⋅↓≻∪≱∖≱∖⊀↓∣⋡⊲↓↓⋠↓↿∙∖⇁∪⇂∎↓⋅⋖⋅≼∙∪⊔≱∖⇂↓⋅⋯∼⇂⊀↓⊔⋏∙≟," As before, we tested the possibility of reconstructing lens parameters"
+⋜⋯∠⊓⇥⊳⇀∖⊳∖��⋡∢⋅⇂∎∪↓⋅∢⋅⋡∖∖⊽∢⋅⊓⋅≱∖⊓⋅∠⇂↥↓↥∢⋅↓≻∪≱∖≱∖⊀↓∣⋡⊲↓↓⋠↓↿∙∖⇁∪⇂∎↓⋅⋖⋅≼∙∪⊔≱∖⇂↓⋅⋯∼⇂⊀↓⊔⋏∙≟↓," As before, we tested the possibility of reconstructing lens parameters"
+⋜⋯∠⊓⇥⊳⇀∖⊳∖∣⋡∢⋅⇂∎∪↓⋅∢⋅⋡∖∖⊽∢⋅⊓⋅≱∖⊓⋅∠⇂↥↓↥∢⋅↓≻∪≱∖≱∖⊀↓∣⋡⊲↓↓⋠↓↿∙∖⇁∪⇂∎↓⋅⋖⋅≼∙∪⊔≱∖⇂↓⋅⋯∼⇂⊀↓⊔⋏∙≟↓⋖," As before, we tested the possibility of reconstructing lens parameters"
+⋜⋯∠⊓⇥⊳⇀∖⊳∖∣⋡∢⋅⇂∎∪↓⋅∢⋅⋡∖∖⊽∢⋅⊓⋅≱∖⊓⋅∠⇂↥↓↥∢⋅↓≻∪≱∖≱∖⊀↓∣⋡⊲↓↓⋠↓↿∙∖⇁∪⇂∎↓⋅⋖⋅≼∙∪⊔≱∖⇂↓⋅⋯∼⇂⊀↓⊔⋏∙≟↓⋖⊾," As before, we tested the possibility of reconstructing lens parameters"
+⋜⋯∠⊓⇥⊳⇀∖⊳∖∣⋡∢⋅⇂∎∪↓⋅∢⋅⋡∖∖⊽∢⋅⊓⋅≱∖⊓⋅∠⇂↥↓↥∢⋅↓≻∪≱∖≱∖⊀↓∣⋡⊲↓↓⋠↓↿∙∖⇁∪⇂∎↓⋅⋖⋅≼∙∪⊔≱∖⇂↓⋅⋯∼⇂⊀↓⊔⋏∙≟↓⋖⊾⊔," As before, we tested the possibility of reconstructing lens parameters"
+⋜⋯∠⊓⇥⊳⇀∖⊳∖∣⋡∢⋅⇂∎∪↓⋅∢⋅⋡∖∖⊽∢⋅⊓⋅≱∖⊓⋅∠⇂↥↓↥∢⋅↓≻∪≱∖≱∖⊀↓∣⋡⊲↓↓⋠↓↿∙∖⇁∪⇂∎↓⋅⋖⋅≼∙∪⊔≱∖⇂↓⋅⋯∼⇂⊀↓⊔⋏∙≟↓⋖⊾⊔⊳," As before, we tested the possibility of reconstructing lens parameters"
+⋜⋯∠⊓⇥⊳⇀∖⊳∖∣⋡∢⋅⇂∎∪↓⋅∢⋅⋡∖∖⊽∢⋅⊓⋅≱∖⊓⋅∠⇂↥↓↥∢⋅↓≻∪≱∖≱∖⊀↓∣⋡⊲↓↓⋠↓↿∙∖⇁∪⇂∎↓⋅⋖⋅≼∙∪⊔≱∖⇂↓⋅⋯∼⇂⊀↓⊔⋏∙≟↓⋖⊾⊔⊳∖," As before, we tested the possibility of reconstructing lens parameters"
+⋜⋯∠⊓⇥⊳⇀∖⊳∖∣⋡∢⋅⇂∎∪↓⋅∢⋅⋡∖∖⊽∢⋅⊓⋅≱∖⊓⋅∠⇂↥↓↥∢⋅↓≻∪≱∖≱∖⊀↓∣⋡⊲↓↓⋠↓↿∙∖⇁∪⇂∎↓⋅⋖⋅≼∙∪⊔≱∖⇂↓⋅⋯∼⇂⊀↓⊔⋏∙≟↓⋖⊾⊔⊳∖↓," As before, we tested the possibility of reconstructing lens parameters"
+⋜⋯∠⊓⇥⊳⇀∖⊳∖∣⋡∢⋅⇂∎∪↓⋅∢⋅⋡∖∖⊽∢⋅⊓⋅≱∖⊓⋅∠⇂↥↓↥∢⋅↓≻∪≱∖≱∖⊀↓∣⋡⊲↓↓⋠↓↿∙∖⇁∪⇂∎↓⋅⋖⋅≼∙∪⊔≱∖⇂↓⋅⋯∼⇂⊀↓⊔⋏∙≟↓⋖⊾⊔⊳∖↓≻," As before, we tested the possibility of reconstructing lens parameters"
+⋜⋯∠⊓⇥⊳⇀∖⊳∖∣⋡∢⋅⇂∎∪↓⋅∢⋅⋡∖∖⊽∢⋅⊓⋅≱∖⊓⋅∠⇂↥↓↥∢⋅↓≻∪≱∖≱∖⊀↓∣⋡⊲↓↓⋠↓↿∙∖⇁∪⇂∎↓⋅⋖⋅≼∙∪⊔≱∖⇂↓⋅⋯∼⇂⊀↓⊔⋏∙≟↓⋖⊾⊔⊳∖↓≻⋜," As before, we tested the possibility of reconstructing lens parameters"
+⋜⋯∠⊓⇥⊳⇀∖⊳∖∣⋡∢⋅⇂∎∪↓⋅∢⋅⋡∖∖⊽∢⋅⊓⋅≱∖⊓⋅∠⇂↥↓↥∢⋅↓≻∪≱∖≱∖⊀↓∣⋡⊲↓↓⋠↓↿∙∖⇁∪⇂∎↓⋅⋖⋅≼∙∪⊔≱∖⇂↓⋅⋯∼⇂⊀↓⊔⋏∙≟↓⋖⊾⊔⊳∖↓≻⋜⊔," As before, we tested the possibility of reconstructing lens parameters"
+⋜⋯∠⊓⇥⊳⇀∖⊳∖∣⋡∢⋅⇂∎∪↓⋅∢⋅⋡∖∖⊽∢⋅⊓⋅≱∖⊓⋅∠⇂↥↓↥∢⋅↓≻∪≱∖≱∖⊀↓∣⋡⊲↓↓⋠↓↿∙∖⇁∪⇂∎↓⋅⋖⋅≼∙∪⊔≱∖⇂↓⋅⋯∼⇂⊀↓⊔⋏∙≟↓⋖⊾⊔⊳∖↓≻⋜⊔⋅," As before, we tested the possibility of reconstructing lens parameters"
+⋜⋯∠⊓⇥⊳⇀∖⊳∖∣⋡∢⋅⇂∎∪↓⋅∢⋅⋡∖∖⊽∢⋅⊓⋅≱∖⊓⋅∠⇂↥↓↥∢⋅↓≻∪≱∖≱∖⊀↓∣⋡⊲↓↓⋠↓↿∙∖⇁∪⇂∎↓⋅⋖⋅≼∙∪⊔≱∖⇂↓⋅⋯∼⇂⊀↓⊔⋏∙≟↓⋖⊾⊔⊳∖↓≻⋜⊔⋅⋜," As before, we tested the possibility of reconstructing lens parameters"
+⋜⋯∠⊓⇥⊳⇀∖⊳∖∣⋡∢⋅⇂∎∪↓⋅∢⋅⋡∖∖⊽∢⋅⊓⋅≱∖⊓⋅∠⇂↥↓↥∢⋅↓≻∪≱∖≱∖⊀↓∣⋡⊲↓↓⋠↓↿∙∖⇁∪⇂∎↓⋅⋖⋅≼∙∪⊔≱∖⇂↓⋅⋯∼⇂⊀↓⊔⋏∙≟↓⋖⊾⊔⊳∖↓≻⋜⊔⋅⋜⋯," As before, we tested the possibility of reconstructing lens parameters"
+⋜⋯∠⊓⇥⊳⇀∖⊳∖∣⋡∢⋅⇂∎∪↓⋅∢⋅⋡∖∖⊽∢⋅⊓⋅≱∖⊓⋅∠⇂↥↓↥∢⋅↓≻∪≱∖≱∖⊀↓∣⋡⊲↓↓⋠↓↿∙∖⇁∪⇂∎↓⋅⋖⋅≼∙∪⊔≱∖⇂↓⋅⋯∼⇂⊀↓⊔⋏∙≟↓⋖⊾⊔⊳∖↓≻⋜⊔⋅⋜⋯↓," As before, we tested the possibility of reconstructing lens parameters"
+⋜⋯∠⊓⇥⊳⇀∖⊳∖∣⋡∢⋅⇂∎∪↓⋅∢⋅⋡∖∖⊽∢⋅⊓⋅≱∖⊓⋅∠⇂↥↓↥∢⋅↓≻∪≱∖≱∖⊀↓∣⋡⊲↓↓⋠↓↿∙∖⇁∪⇂∎↓⋅⋖⋅��∙∪⊔≱∖⇂↓⋅⋯∼⇂⊀↓⊔⋏∙≟↓⋖⊾⊔⊳∖↓≻⋜⊔⋅⋜⋯↓∢," As before, we tested the possibility of reconstructing lens parameters"
+⋜⋯∠⊓⇥⊳⇀∖⊳∖∣⋡∢⋅⇂∎∪↓⋅∢⋅⋡∖∖⊽∢⋅⊓⋅≱∖⊓⋅∠⇂↥↓↥∢⋅↓≻∪≱∖≱∖⊀↓∣⋡⊲↓↓⋠↓↿∙∖⇁∪⇂∎↓⋅⋖⋅≼∙∪⊔≱∖⇂↓⋅⋯∼⇂⊀↓⊔⋏∙≟↓⋖⊾⊔⊳∖↓≻⋜⊔⋅⋜⋯↓∢⋅," As before, we tested the possibility of reconstructing lens parameters"
+⋜⋯∠⊓⇥⊳⇀∖⊳∖∣⋡∢⋅⇂∎∪↓⋅∢⋅⋡∖∖⊽∢⋅⊓⋅≱∖⊓⋅∠⇂↥↓↥∢⋅↓≻∪≱∖≱∖⊀↓∣⋡⊲↓↓⋠↓↿∙∖⇁∪⇂∎↓⋅⋖⋅≼∙∪⊔≱∖⇂↓⋅⋯∼⇂⊀↓⊔⋏∙≟↓⋖⊾⊔⊳∖↓≻⋜⊔⋅⋜⋯↓∢⋅∩," As before, we tested the possibility of reconstructing lens parameters"
+⋜⋯∠⊓⇥⊳⇀∖⊳∖∣⋡∢⋅⇂∎∪↓⋅∢⋅⋡∖∖⊽∢⋅⊓⋅≱∖⊓⋅∠⇂↥↓↥∢⋅↓≻∪≱∖≱∖⊀↓∣⋡⊲↓↓⋠↓↿∙∖⇁∪⇂∎↓⋅⋖⋅≼∙∪⊔≱∖⇂↓⋅⋯∼⇂⊀↓⊔⋏∙≟↓⋖⊾⊔⊳∖↓≻⋜⊔⋅⋜⋯↓∢⋅∩⊾," As before, we tested the possibility of reconstructing lens parameters"
+⋜⋯∠⊓⇥⊳⇀∖⊳∖∣⋡∢⋅⇂∎∪↓⋅∢⋅⋡∖∖⊽∢⋅⊓⋅≱∖⊓⋅∠⇂↥↓↥∢⋅↓≻∪≱∖≱∖⊀↓∣⋡⊲↓↓⋠↓↿∙∖⇁∪⇂∎↓⋅⋖⋅≼∙∪⊔≱∖⇂↓⋅⋯∼⇂⊀↓⊔⋏∙≟↓⋖⊾⊔⊳∖↓≻⋜⊔⋅⋜⋯↓∢⋅∩⊾↓," As before, we tested the possibility of reconstructing lens parameters"
+⋜⋯∠⊓⇥⊳⇀∖⊳∖∣⋡∢⋅⇂∎∪↓⋅∢⋅⋡∖∖⊽∢⋅⊓⋅≱∖⊓⋅∠⇂↥↓↥∢⋅↓≻∪≱∖≱∖⊀↓∣⋡⊲↓↓⋠↓↿∙∖⇁∪⇂∎↓⋅⋖⋅≼∙∪⊔≱∖⇂↓⋅⋯∼⇂⊀↓⊔⋏∙≟↓⋖⊾⊔⊳∖↓≻⋜⊔⋅⋜⋯↓∢⋅∩⊾↓⋅," As before, we tested the possibility of reconstructing lens parameters"
+⋜⋯∠⊓⇥⊳⇀∖⊳∖∣⋡∢⋅⇂∎∪↓⋅∢⋅⋡∖∖⊽∢⋅⊓⋅≱∖⊓⋅∠⇂↥↓↥∢⋅↓≻∪≱∖≱∖⊀↓∣⋡⊲↓↓⋠↓↿∙∖⇁∪⇂∎↓⋅⋖⋅≼∙∪⊔≱∖⇂↓⋅⋯∼⇂⊀↓⊔⋏∙≟↓⋖⊾⊔⊳∖↓≻⋜⊔⋅⋜⋯↓∢⋅∩⊾↓⋅⊳," As before, we tested the possibility of reconstructing lens parameters"
+⋜⋯∠⊓⇥⊳⇀∖⊳∖∣⋡∢⋅⇂∎∪↓⋅∢⋅⋡∖∖⊽∢⋅⊓⋅≱∖⊓⋅∠⇂↥↓↥∢⋅↓≻∪≱∖≱∖⊀↓∣⋡⊲↓↓⋠↓↿∙∖⇁∪⇂∎↓⋅⋖⋅≼∙∪⊔≱∖⇂↓⋅⋯∼⇂⊀↓⊔⋏∙≟↓⋖⊾⊔⊳∖↓≻⋜⊔⋅⋜⋯↓∢⋅∩⊾↓⋅⊳∖," As before, we tested the possibility of reconstructing lens parameters"
+pressure £54. ἐνοι. From Equations (4)) and (5)) we point out tha neglecting £4 (Le. setting 2.= Lin) systematically. biases low the determination of cluster masses. roughly by a factor of £.,"pressure $P_{\rm tot}$, i.e., From Equations \ref{aa4}) ) and \ref{pnt12}) ) we point out that neglecting $P_{\rm nt}$ (i.e., setting $P_{\rm tot} = P_{\rm th}$ ) systematically biases low the determination of cluster masses, roughly by a factor of $\xi$."
+ Note that X-ray data probe only the therma component of the gas Py=nm.Κω.," Note that X-ray data probe only the thermal component of the gas $P_{\rm th}=n_e\, {\bf k} T_{\rm gas}$."
+ The work of Lee&Suto(2003). showed that the IC eas and DM halos are well approximated by a sequence of concentric triaxial clistributions with cillerent cecentricity ratio., The work of \cite{lee2003} showed that the IC gas and DM halos are well approximated by a sequence of concentric triaxial distributions with different eccentricity ratio.
+" We define ο (ey) and es (6,7) as the cccentricity of DM UC gas) on the plane of the sky and along the line of sight. respectively."," We define $e_{b'}$ $\epsilon_{b'}$ ) and $e_{c'}$ $\epsilon_{c'}$ ) as the eccentricity of DM (IC gas) on the plane of the sky and along the line of sight, respectively."
+ Ehe iso-potential surfaces of the triaxia dark halo coincide also with the iso-densitv (pressure. temperature) surfaces of the intracluster gas.," The iso-potential surfaces of the triaxial dark halo coincide also with the iso-density (pressure, temperature) surfaces of the intracluster gas."
+ This is simply a direct consequence of the (Buote&Canizares 1994): the hyedrostatie equilibrium equation (4)) vields Note that ο=eyt(eyino) amd e;‘=clea). with u—r/r.. unlike the constant e.c; Lor the adopted DM halo profile.," This is simply a direct consequence of the \citep{buote1994}; the hydrostatic equilibrium equation \ref{aa4}) ) yields Note that $\epsilon_{b'}=\epsilon_{b'}(e_{b'},u,\alpha)$ and $\epsilon_{c'}=\epsilon_{c'}(e_{c'},u,\alpha)$ , with $\bu \equiv
+\br/r_{\rm s}$, unlike the constant $e_{b'},e_{c'}$ for the adopted DM halo profile."
+ In the whole range of i. esfen (cof oor) is less than unity (~0.7 at the center). ie. the intracluster eas is altoeether more spherical than the underlying DM halo.," In the whole range of $u$, $\epsilon_{b'}/e_{b'}$ $\epsilon_{c'}/e_{c'}$ ) is less than unity $\sim 0.7$ at the center), i.e., the intracluster gas is altogether more spherical than the underlying DM halo."
+ We construct the likelihood performing a joint analysis for lensing and X-ray data. in order to constrain the properties of the model parameters q representing the concentration parameter. scale radius. inner slope of the DM. cecentricity of the DAL along the line of sight ancl fractional non-thermal pressure of the eas. respectively.," We construct the likelihood performing a joint analysis for lensing and X-ray data, in order to constrain the properties of the model parameters ${\bf q}$ representing the concentration parameter, scale radius, inner slope of the DM, eccentricity of the DM along the line of sight and fractional non-thermal pressure of the gas, respectively."
+ Therefore. the system. of equation we simultaneously rely on in our joint N-ray|Lensing analysis is: where the three-dimensional model temperature. Lo. is recovered. by solving equation (4)) and constrained by the observed. temperature profile. and. the model. surface mass prolile X is recovered by projection of the triaxial NEW DAL moclel ancl constrained by weak/strong lensing e@measurements.," Therefore, the system of equation we simultaneously rely on in our joint X-ray+Lensing analysis is: where the three-dimensional model temperature $T_{\rm gas}$ is recovered by solving equation \ref{aa4}) ) and constrained by the observed temperature profile, and the model surface mass profile $\Sigma$ is recovered by projection of the triaxial gNFW DM model and constrained by weak/strong lensing measurements."
+ The method works by constructing. a joint rav|Lensing likelihood: Lx and Lyon. being the likelihoods coming from the and lensing (WLISL) data. respectively.," The method works by constructing a joint X-ray+Lensing likelihood: $\mathcal{L}_{\rm X}$ and ${\mathcal{L}}_{\rm lens}$ being the likelihoods coming from the X-ray and lensing (WL+SL) data, respectively."
+ Now. £xxAx 2). where being the observed. projected temperature profile in the Doi7th ring and. μη) the convenient projection (FollowingAlazzottactal.2004). of the theoretical three-dimensional temperature {μμ(ο): the latter is the result of solving the hyvdrostatic equilibrium equation. with the gas density ορτο) inferred from the X-ray surface brightness. and assuming à eNEW parametrization for the DM Ὁeruso0).," Now, ${\mathcal{L}}_{\rm X}\propto {\exp ( -\chi_{\rm
+X}^2/2)}$ , where $T^*_{\rm proj,i}$ being the observed projected temperature profile in the $i$ 'th ring and $T_{\rm proj,i}({\bf q})$ the convenient projection \citep[following][]{mazzotta2004} of the theoretical three-dimensional temperature $T_{\rm gas}({\bf q})$; the latter is the result of solving the hydrostatic equilibrium equation, with the gas density $n_e(r; {\epsilon_{c'}})$ inferred from the X-ray surface brightness, and assuming a gNFW parametrization for the DM $\rho_{\rm
+DM}=\rho_{\rm DM}({\bf r}; c,r_{\rm
+s},\alpha,e_{c'})$ ."
+ The lensing contribution is where C is the covariance matrix of the projected mass profile from strong (C) and weak (C4): lensing data. |C indicates the determinant of C. NM=(Vi.N35...Ni) are the observed. measurements of the projected. mass profile in the m elliptical annuli. and Xe.rz0.607) is the theoretical projected mass profile within our triaxialDAL model.," The lensing contribution is where $\mathbfit{C}$ is the covariance matrix of the projected mass profile from strong $\mathbfit{C}_{\rm sl}$ ) and weak $\mathbfit{C}_{\rm wl}$ ) lensing data, $|\mathbfit{C}|$ indicates the determinant of $\mathbfit{C}$, $\Sigma^*=(\Sigma_1^*,\Sigma_2^*,...,\Sigma_{m^*}^*)$ are the observed measurements of the projected mass profile in the $m^*$ elliptical annuli, and $\Sigma(c,r_{\rm s},\alpha,e_{c'})$ is the theoretical projected mass profile within our triaxialDM model."
+ For the covariance matrix of the strong lensing nmieasurements. we also account for systematic errors by replacing the observed covariance matrix with the following expression: where Z is the identity matrix.," For the covariance matrix of the strong lensing measurements, we also account for systematic errors by replacing the observed covariance matrix with the following expression: where ${\mathcal{I}}$ is the identity matrix."
+ Since we are using the surface mass density as output. [rom the parametric SL analysis of Limousinetal.(2007).. we introduce the parameter a. to allow some extra freedom due to possible systematic bias resulting from their analysis.," Since we are using the surface mass density as output from the parametric SL analysis of \cite{Limousin2007}, we introduce the parameter $\sigma_{\rm sys}$ to allow some extra freedom due to possible systematic bias resulting from their analysis."
+ La order to caleulate oy. we assumed that systematic errors can be described. as Gaussian. errors. via a diagonal matrix On.5F with the same value in each of the diagonal elements., In order to calculate $\sigma_{\rm sys}$ we assumed that systematic errors can be described as Gaussian errors via a diagonal matrix $\sigma^2_{\rm sys}\; {\mathcal{I}}$ with the same value in each of the diagonal elements.
+ We checked that this simplified assumption does not significantly allect. the average value of the physical parameters. while it slightly increases 20%)) their errors with respect to the case where we neglect a...," We checked that this simplified assumption does not significantly affect the average value of the physical parameters, while it slightly increases ) their errors with respect to the case where we neglect $\sigma_{\rm sys}$."
+ We evaluated he probability distribution function of model parameters £=£(q.05.) via the Markoy. Chain Alonte Carlo (AICAIC) algorithm.," We evaluated the probability distribution function of model parameters ${\mathcal{L}}={\mathcal{L}}({\bf q}, \sigma_{\rm sys})$ via the Markov Chain Monte Carlo (MCMC) algorithm."
+" Thus we can determine our best-fit. model parameters and their errors. and. derive various properties of the cluster,"," Thus we can determine our best-fit model parameters and their errors, and derive various properties of the cluster."
+ Inthe previous section we showed how we can determine the physical parameters of the cluster by fitting the available data based on the hyelrostatic equilibrium equation and on a DM moclel that is based on robust results of hvedrodynamical cluster simulations., In the previous section we showed how we can determine the physical parameters of the cluster by fitting the available data based on the hydrostatic equilibrium equation and on a DM model that is based on robust results of hydrodynamical cluster simulations.
+ In this section we present our results and discuss their main implications., In this section we present our results and discuss their main implications.
+ We particularly focus on the implications of our analysis for the viabilitv of the CDM scenario. the discrepancy between X-ray and lensing masses in Abell 1689. and the presence of non-thermal pressure.," We particularly focus on the implications of our analysis for the viability of the CDM scenario, the discrepancy between X-ray and lensing masses in Abell 1689, and the presence of non-thermal pressure."
+ In Figure 1. we present the results of our joint analysis for Abell 1689., In Figure \ref{entps332333} we present the results of our joint analysis for Abell 1689.
+" Both the X-ray and lensing data are reasonably fitted by our model. witha total V,=A&LASLZB88 (with 20 degrees of freedom)."," Both the X-ray and lensing data are reasonably fitted by our model, witha total $\chi^2_{\rm tot}=\chi^2_{\rm
+X}+\chi^2_{\rm sl}+\chi^2_{\rm wl}=33.8$ (with 20 degrees of freedom)."
+ The separate contributions are o124.6.3.15.1 for X-ray. SL and WL data. respectively. where TNx2log( £a).," The separate contributions are $\chi^2=12.4,6.3,15.1$ for X-ray, SL and WL data, respectively, where $\chi^2_{\rm lens}\propto
+-2\log({\mathcal{L}}_{\rm lens})$ ."
+ Note that in Figure the WL model predictions are slightlv in disagreement with1. the data. overpredicting the surface mass profile in the range 500-2000 kpe at. around," Note that in Figure \ref{entps332333} the WL model predictions are slightly in disagreement with the data, overpredicting the surface mass profile in the range 500-2000 kpc at around"
+llere we present hydrodynamical simulations of supernova ejecta interaction with an AM modilied by presupernova wind.,Here we present hydrodynamical simulations of supernova ejecta interaction with an AM modified by presupernova wind.
+ We have worked out simple moclels of time dependent wind. devised to reproduce (he history of wind given by pre-SNla binary system evolutionary models.," We have worked out simple models of time dependent wind, devised to reproduce the history of wind given by pre-SNIa binary system evolutionary models."
+ The main results of the simulations are presented and (he compatibility of each wind model with known SNR from SNla is discussed. based on the different outcomes of the caleulations (mainly. expansion parameters. velocity of the ejecta at (he reverse shock. and radius of the SNR).," The main results of the simulations are presented and the compatibility of each wind model with known SNR from SNIa is discussed, based on the different outcomes of the calculations (mainly, expansion parameters, velocity of the ejecta at the reverse shock, and radius of the SNR)."
+ Current models of presupernova binary evolution describe wind ejection as a subproduct of matter accretion on (he primary star. once it has become a white dwarl.," Current models of presupernova binary evolution describe wind ejection as a subproduct of matter accretion on the primary star, once it has become a white dwarf."
+ In the last vears. (here has been a wide interest on the possibility (Ilachisu.Kato.&Nomoto1996) that an optically thick wind stabilizes the mass (rausler in (he binary. svstem. (hus opening a new channel for the white dwail to grow up to the Chandrasekhar mass ancl explode as a SNla (Li&vandenHeuvel1997;NingTeeseling 1999).," In the last years, there has been a wide interest on the possibility \citep{hkn96} that an optically thick wind stabilizes the mass transfer in the binary system, thus opening a new channel for the white dwarf to grow up to the Chandrasekhar mass and explode as a SNIa \citep{lv97,kv98,hkn99}."
+. The time history of wind is complex. ancl depends on orbital parameters as well as on (he capability of the white dwarf to quietly accrete (he matter transferred [rom the secondary star.," The time history of wind is complex, and depends on orbital parameters as well as on the capability of the white dwarf to quietly accrete the matter transferred from the secondary star."
+ An important characteristic of the wind history is whether the wind is aclive up to the time of SN explosion or. on the contrary. the binary has a time extended phase of conservative evolution prior to white dwarf ignition.," An important characteristic of the wind history is whether the wind is active up to the time of SN explosion or, on the contrary, the binary has a time extended phase of conservative evolution prior to white dwarf ignition."
+ Another important parameter is the velocity of the wind., Another important parameter is the velocity of the wind.
+ Iere we explore the effects of the wind on the SNR. evolution. and compare wilh the results obtained on the assumption that there is no significant presupernova wind.," Here we explore the effects of the wind on the SNR evolution, and compare with the results obtained on the assumption that there is no significant presupernova wind."
+For cach model. we begin with a column density. Doppler paraicter. and covering factor forAleul.,"For each model, we begin with a column density, Doppler parameter, and covering factor for."
+ For Cloud 1. we consider a rauge of N(ALei). b(ALe) pairs that are consistent with the data. inchicling fits affected by unresolved saturation.," For Cloud 1, we consider a range of $N({\MgII})$ , $b({\rm Mg})$ pairs that are consistent with the data, including fits affected by unresolved saturation."
+ Starting with the column density as a constraimt. we used the code Cloudy (versioLE 07.02.01) to conduct photoionization models (Ferlandeal.1998).," Starting with the column density as a constraint, we used the code Cloudy (version 07.02.01) to conduct photoionization models \citep{Ferland98}."
+. For cach cloud. we assume a plane-parallel slab with the eiven IN(Mgi) and --ilunünuate it with a Taardt-\Macdan extragalactic backerouncd spectrin. including quasars and star-forming ealaxics (with au escape fraction of 0.1) at 2=0.1521 (aardt&Madan1996. 2001).," For each cloud, we assume a plane-parallel slab with the given $N({\MgII})$ and illuminate it with a Haardt-Madau extragalactic background spectrum, including quasars and star-forming galaxies (with an escape fraction of 0.1) at $z=0.4524$ \citep{HM96,
+HM01}."
+. Given that the absorption ds so weak inMeu. it seems unlikely that this absorber is housed i1 the mudst of a ealaxy where the local stellar radiatio1 field would be siguificant.," Given that the absorption is so weak in, it seems unlikely that this absorber is housed in the midst of a galaxy where the local stellar radiation field would be significant."
+ We initially run a grid of models for ionization paramcters logel=Ἱοουη.S5 to. 2.5. aud metallicities logZ/Z..=39 to 1.0.," We initially run a grid of models for ionization parameters $\log U=\log n_e/n_\gamma=-8.5$ to $-2.5$, and metallicities $\log Z/Z_{\odot}= -3.0$ to $1.0$."
+ All models assuue a solu abundance pattern. uuless otherwise stated.," All models assume a solar abundance pattern, unless otherwise stated."
+ The Cloudy output inchicdes model column deusities for all transitions. as well as au equilibrimu temperature. T. for the cloud.," The Cloudy output includes model column densities for all transitions, as well as an equilibrium temperature, $T$, for the cloud."
+" The turbulent component of the Doppler parameter IDA=Diner(AleP )-: where the thermal cuponeut b,,,,2= 2hTy is calculated from the equilibria temperature aud observed ó(Meu).", The turbulent component of the Doppler parameter $b^2_{turb}=b({\MgII})^2-b_{therm}(\mathrm{Mg})^2$ ; where the thermal component $b^2_{therm}=\frac{2kT}{m}$ ) is calculated from the equilibrium temperature and observed $b({\MgII})$.
+" Given Pay, aud bi. the Doppler parameter can be computed for all other elements."," Given $b_{turb}$ and $b_{therm}$, the Doppler parameter can be computed for all other elements."
+ After running Cloudy for all commponcuts from the Voigt profile fit. the b paraiucters and Cy values are combined with the model cohunu deusities to create a svuthetic spectrin.," After running Cloudy for all components from the Voigt profile fit, the $b$ parameters and $C_f$ values are combined with the model column densities to create a synthetic spectrum."
+ This model spectruu is compared to the data in order to constrain the log aud logZ/Z.. values., This model spectrum is compared to the data in order to constrain the $\log U$ and $\log Z/Z_{\odot}$ values.
+ The modeling method is the same as that used by Dine.etal...(2005). Lyuch&Charlton (2007).. aud \lisawaetal.(2007).," The modeling method is the same as that used by \citet{Ding05}, \citet{Lynch07}, and \citet{Misawa07}."
+. Cloud 2 cau be fit with 2.9<log— 3X lower values overproduceFerr while higher values uuderproduceFel.," Cloud 2 can be fit with $-2.9<\log U<-3.3$ ; lower values overproduce, while higher values underproduce."
+ Although there are no να series lines to use for direct coustraiuts on the metallicity. we find that if loyZ/Z..«0.3. the is overproduced.," Although there are no Lyman series lines to use for direct constraints on the metallicity, we find that if $Z/Z_{\odot}< -0.3$, the is overproduced."
+ The temperature of the cloud is found to be ~ 90001& for logZ/Z..—0.3. with a size of teus of parsecs. aud a density ~ Ix107.," The temperature of the cloud is found to be $\sim 9000$ K for $\log Z/Z_{\odot}= -0.3$, with a size of tens of parsecs, and a density $\sim 4$ $10^{-3}$."
+ Model parameters are listed in Table {ιν where Clouds 2.1 and 2.2 represent the two components of Cloud 2. while the best model is shown iu Figure 2..," Model parameters are listed in Table \ref{tab:models}, where Clouds $2.1$ and $2.2$ represent the two components of Cloud 2, while the best model is shown in Figure \ref{fig:system_plot_unresolved}."
+ Clouds 3 and lL are quite weak. and are detected in the VET spectrum only because it has such a high S/N. For Clouds 3 and £. a wide range of metalliciticsμα and ionization parameters provide an adequate fit to he transitions covered by the spectra.," Clouds 3 and 4 are quite weak, and are detected in the VLT spectrum only because it has such a high S/N. For Clouds 3 and 4, a wide range of metallicities and ionization parameters provide an adequate fit to the transitions covered by the spectrum."
+ Low ionization species other than are not detected at this velocity. mit because of the weakness of the lines. are also not medicted by models with any reasonable iouization xranmeter. logC7.5.," Low ionization species other than are not detected at this velocity, but because of the weakness of the lines, are also not predicted by models with any reasonable ionization parameter, $\log U
+> -7.5$."
+ Coverage of higher ionization states would be needed to further constrain the ionization xuiuueter of these clouds., Coverage of higher ionization states would be needed to further constrain the ionization parameter of these clouds.
+" Coustraits for these clouds"" are sununarized in Table L. where piriuneters are give1 or two acceptable values of logC aud logZ/Z.. for cac1 Cloud."," Constraints for these clouds are summarized in Table \ref{tab:models}, where parameters are given for two acceptable values of $\log U$ and $Z/{Z_{\odot}}$ for each Cloud."
+ We run a grid of models with varvine «οανα! densities. nietallicities. aud ionization paramicters. as deseribed iu section 3.. for ΕΝ+. the Doppler paralucter given by a Voigt profile ft. which assumes the line is resolved. uuisaturated. and fully covered.," We run a grid of models with varying column densities, metallicities, and ionization parameters, as described in section \ref{sec:3}, , for $b=3.1$, the Doppler parameter given by a Voigt profile fit, which assumes the line is resolved, unsaturated, and fully covered."
+ We find that detectable amounts of are not produced in anv model aud that is always uuderproduced.," We find that detectable amounts of are not produced in any model, and that is always underproduced."
+ For sinaller b paraiucters we would expect that Fel/Ale. and Feiu/Meth would be larger because aud would be ou the linear parts of their curves of erowtl where the correspouding larger ;N would affect their equivalent widths., For smaller $b$ parameters we would expect that ${\FeI}/{\MgII}$ and ${\FeII}/{\MgII}$ would be larger because and would be on the linear parts of their curves of growth where the corresponding larger $N$ would affect their equivalent widths.
+ Tn contrast. theMei would be on the flat part of its curve of growth so that the increased AN would have little effect on its equivalent width.," In contrast, the would be on the flat part of its curve of growth so that the increased $N$ would have little effect on its equivalent width."
+" We therefore considered IN(Mgir). b{(Ale) pairswith >=OlO5kuns + (giving logN(Mgl)~11 7) and logU«τη,"," We therefore considered $N({\MgII})$, $b({\rm
+Mg})$ pairswith $b=0.1-0.5$ (giving $\log N({\MgII}) \sim 14$ ) and $\log U < -7.5$."
+" An example model. with lozZ/Z.=1, 8.2, and 6—0.2 is shown in Fig. 2.."," An example model, with $Z/{Z_{\odot}} =-1.0$, $\log U=-8.2$ , and $b=0.2$ is shown in Fig. \ref{fig:system_plot_unresolved}."
+ Due to the secnunely cold nature of this cloud. we opt to add dust erains to the Cloudy (Ferlandetal.1998). simulations.," Due to the seemingly cold nature of this cloud, we opt to add dust grains to the Cloudy \citep{Ferland98} simulations."
+ The primary flaw of this model is the over-production of A2853 bv ~1 dex iu cohuun deusitv: the apparent nuderproduction of some aud transitionsnav be attributed to forest contamination of the observed profiles., The primary flaw of this model is the over-production of $\lambda$ 2853 by $\sim 1$ dex in column density; the apparent underproduction of some and transitionsmay be attributed to forest contamination of the observed profiles.
+ Caveater mectallicities further overproduceMet. and a small Fe abundance cuhancement of ~0.2 dex is needed to account for observed and profiles.," Greater metallicities further overproduce, and a small Fe abundance enhancement of $\sim 0.2$ dex is needed to account for observed and profiles."
+" Constraints for this model are eiven in Table 1.. nuder Cloud 1""."," Constraints for this model are given in Table \ref{tab:models}, under Cloud $1^{a}$."
+ We fud a size of 0.01-0.6 pe. T«100 Ik. aud 500«ay<1100 3).," We find a size of 0.01-0.6 pc, $<100$ K, and $500<n_H<1100$ ]."
+ With a neutral hydrogen column deusitv of LOY[anδα (sub-DLA to DLA range}.uneeh we find⋅ 0.72-2o3Mint222:2N(M3S)HH3)," With a neutral hydrogen column density of $10^{18.7-20.8}$ (sub-DLA to DLA range), we find $0.72<
+\frac{2N(H_2)}{2N(H_2)+N(HI)} < 0.91$."
+vain0s7qe40.91. ~10 ντ ου., $\rm{log}\frac{N(MgII)}{N(MgI)}\sim 1.0-1.3$ and $\rm{log}\frac{N(FeII)}{N(FeI)} \sim 1.6-2.2$.
+ Wo similarly find that a small b (<0.5kins 31) is necessary fo reproduce the observed with partial coverage models., We similarly find that a small $b$ $<0.5~{\kms}$ ) is necessary to reproduce the observed with partial coverage models.
+ An example of such a model. with b(Mg)=O.llans blog!= 75:xdlosZ/Z.- 1.0 is given in Fig. L.," An example of such a model, with $b({\rm Mg}) = 0.4$, $\log U =
+-7.5$, and $\log Z/Z_{\odot} = -1.0$ is given in Fig. \ref{fig:system_plot_partial}."
+ For this example model.the observedForFeu.Car Cain. Muir aud are adequatelv reproduced within the uucertainties.," For this example model,the observed, and are adequately reproduced within the uncertainties."
+ The absorption at he position of À2[8[ and A252 Lis not fully produced. rowever the location of these trausitfious iu the forest nakes contamination fairly likely.," The absorption at the position of $\lambda$ 2484 and $\lambda$ 2524 is not fully produced, however the location of these transitions in the forest makes contamination fairly likely."
+ The ouly discrepancy of this mocel is the over-production of A2853 bv 1 dex in coluun density., The only discrepancy of this model is the over-production of $\lambda$ 2853 by $\sim 1$ dex in column density.
+ Ii addition to this sample nodel. a auge of ionization parameters aud iioetallicities xovide a simulafit.," In addition to this sample model, a range of ionization parameters and metallicities provide a similarfit."
+ For logZ/Z..=—1.0 aud log--7.5 the cloud size/thickuess is0.2 pc. a size comparable o that ofthe quasar BELR.," For $\log Z/Z_{\odot} = -1.0$ and $\log U =-7.5$ the cloud size/thickness is0.2 pc, a size comparable to that ofthe quasar BELR."
+ The range of modelsthat provide au adequate fit o this cloud have extreme properties., The range of modelsthat provide an adequate fit to this cloud have extreme properties.
+The ionization varaineters for successful models rauge frou. 8.0) to 7.0. nuplvius densitiesof 30<ny1100 cm”.,"The ionization parameters for successful models range from $-8.0$ to $-7.0$ , implying densitiesof $30 < n_H <
+1100$ ."
+" For the rauge of possible metallicities. the equilibria cluperaturesare low. <50 Ik. The neutral lvdrogen coluun deusities.eye log.N(IT1)>=yoo15.+19.9 7). are in: he msub-DLA- ranec.""uvas while-:B 0.76=).a<“eeTWh2N(OJBIAh6 0.91."," For the range of possible metallicities, the equilibrium temperaturesare low, $< 50$ K. The neutral hydrogen column densities, $\log N({\HI})=18.8-19.9$ ), are in the sub-DLA range, while $0.76<
+\frac{2N(H_2)}{2N(H_2)+N(HI)} < 0.94$ ."
+ Cloud properties for acceptable models are sunuiuiairnized iu iu Table Lo uuder Cloud 1. while a sample model is, Cloud properties for acceptable models are summarized in in Table \ref{tab:models} under Cloud $1^b$ while a sample model is
+(the radius where the source profile merees with the local backeround taken as source extent) aud bv comparison with the profile expected from the PSF of the detector (for a more detailed discussion of 3383 sce Sect.,(the radius where the source profile merges with the local background taken as source extent) and by comparison with the profile expected from the PSF of the detector (for a more detailed discussion of 383 see Sect.
+ 1.1.2)., 4.1.2).
+" This resulted in source extraction radii o~~ L100"" for the brightest ceutral sources.", This resulted in source extraction radii of $\sim$ for the brightest central sources.
+ The detected sources were then removed within circular regions centered on their X-ray positions., The detected sources were then removed within circular regions centered on their X-ray positions.
+ As backeround we selected a source-free circular region in the outer part of the fov., As background we selected a source-free circular region in the outer part of the fov.
+ Since the spectra below 0.5 keV is strongly dominated by backerouud. we uxuallv excluded photous below this euerev from the spectral fit.," Since the spectrum below 0.5 keV is strongly dominated by background, we usually excluded photons below this energy from the spectral fit."
+ Treating the cold absorption as free parameter always resulted in a value (uearly) consistent with the Galactic absorption towards 3383 (Nu;—0.523107! UC. Dickey Lockinan 1990) within the errors.," Treating the cold absorption as free parameter always resulted in a value (nearly) consistent with the Galactic absorption towards 383 $N_{\rm gal} = 0.523\,10^{21}$ $^{-3}$, Dickey Lockman 1990) within the errors."
+ We therefore fixed Nyp = Nea., We therefore fixed $N_{\rm H}$ = $N_{\rm gal}$.
+ For the total Cluission within we find a temperature AL = 1.50.1 keV. The temperature values derived for the three separate regions are consistent with this value and with each other within the errors (Table 1))., For the total emission within we find a temperature $kT$ = $\pm{0.1}$ keV. The temperature values derived for the three separate regions are consistent with this value and with each other within the errors (Table \ref{fitres}) ).
+ This value of T is somewhat lower than the one given in Moregauti ct al. (, This value of $T$ is somewhat lower than the one given in Morganti et al. (
+1988: their Tab.,1988; their Tab.
+ 2) ou the basis of an IPC observation: they estimate At5 3 keV. To check the robustuess of the obtained temperature. we performed a few tests: If we donot yoniove the point sources and re-fit he total xpectium we ect AT = 10.1. keV: if we fit a hermal bremsstralliuue model we obtain AT = 1.10.2 seV: if an rs plasina with solar iusteac of depleted abundances is assumed. the quality of the fit slightlv improves without affecting the value of the temperature within the errors aud we derive 1 = L601 keV. T97 obtain AT = 1.5 keV for abundauces of «solar.," 2) on the basis of an IPC observation; they estimate $kT \approx$ 3 keV. To check the robustness of the obtained temperature, we performed a few tests: If we do remove the point sources and re-fit the total spectrum we get $kT$ = $\pm{0.1}$ keV; if we fit a thermal bremsstrahlung model we obtain $kT$ = $\pm{0.2}$ keV; if an rs plasma with solar instead of depleted abundances is assumed, the quality of the fit slightly improves without affecting the value of the temperature within the errors and we derive $kT$ = $\pm{0.1}$ keV. T97 obtain $kT$ = 1.5 keV for abundances of $\times$ solar."
+ Results of our spectral fits are summarized iu Table |, Results of our spectral fits are summarized in Table \ref{fitres}.
+ Using AL = 1.5 keV aud Ay=Nea. the total (0.1 keV) luminosity within is L4=1.510% ere/s (= 1.311055 eresS if the abundances are fixed to the solar value).," Using $kT$ = 1.5 keV and $N_{\rm H} = N_{\rm gal}$, the total (0.1--2.4 keV) luminosity within is $L_{\rm x} = 1.5\,10^{43}$ erg/s (= $^{43}$ erg/s if the abundances are fixed to the solar value)."
+ All briel ellipticals of the optical chain 3331 are individually detected ii X-ravs., All bright ellipticals of the optical chain 331 are individually detected in X-rays.
+" The ceuter of the N-ray enission imaximnua from 3383 at a = 1725.9. 6 = VLD 22000) coincides well with the position of the optical nucleus at a = 127™25.0. 6 = ""ss."," The center of the X-ray emission maximum from 383 at $\alpha$ = $^{\rm h}$ $^{\rm m}$ 25.9, $\delta$ = 5 2000) coincides well with the position of the optical nucleus at $\alpha$ = $^{\rm h}$ $^{\rm m}$ 25.0, $\delta$ = 8."
+ For au overlay of the PSPC N-ray contours on an optical nuage of the 3383 eroup see T97 (their Fig., For an overlay of the PSPC X-ray contours on an optical image of the 383 group see T97 (their Fig.
+ 1)., 1).
+ To derive physical propertics of the eroup. Wwe fist asmue spherical svaunietiv of the extended Nav ClUSSION. the chussion to be centered on 3383. and rough isothermality to hold.," To derive physical properties of the group, we first assume spherical symmetry of the extended X-ray emission, the emission to be centered on 383, and rough isothermality to hold."
+" For critical conuueuts on and justification of ""standard! assunrptions sec. e.g. Bohhyinecr et al."," For critical comments on and justification of `standard' assumptions see, e.g., Böhhringer et al."
+ 1998 (their Sects., 1998 (their Sects.
+ 2. 3).," 2, 3)."
+" A 1nodel (ο,οι, Cavaliere Fusco-Fenuano 1976. Corenstein ct al."," A $\beta$ -model (e.g., Cavaliere Fusco-Femiano 1976, Gorenstein et al."
+ 1975. Jones Forman 1981) of the form was fit to the azinuthally averaged surface briehtuess profile of the PSPC observation. (detected point sources were. again. removed except cussion from 3383. inner bius were then excluded from the fit since they are dominated by cussion from 3383 itself).," 1978, Jones Forman 1984) of the form was fit to the azimuthally averaged surface brightness profile of the PSPC observation (detected point sources were, again, removed except emission from 383, inner bins were then excluded from the fit since they are dominated by emission from 383 itself)."
+ This vields a central surface brightucss Sy=2.7910.? cts/s/aremin?. a slope parameter J = 0.3 and a core racius ο = 61kpc.," This yields a central surface brightness $S_{\rm 0} = 2.79\,10^{-3}$ $^2$, a slope parameter $\beta$ = 0.34 and a core radius $r_{\rm c}$ = 64 kpc."
+" The eas mass euclosed inside 1 Mpc amounts to Ma, = 1.511015 AL...", The gas mass enclosed inside 1 Mpc amounts to $M_{\rm gas}$ = $^{13}$ $_{\odot}$.
+ Tuspection of Fig., Inspection of Fig.
+ 1 shows that 3383 is not located perfectly in the ceuter of the large scale X-rav cussion. but offset bv about to the North-East.," 1 shows that 383 is not located perfectly in the center of the large scale X-ray emission, but offset by about to the North-East."
+" In fact. using the best-fit beta-model to construct a svuthetic model inageex""p aud subtracting the model image frou the observed oue leaves some residual extended cussion to the SW aud some ""negative. emissiou levels to the NE."," In fact, using the best-fit beta-model to construct a synthetic model image and subtracting the model image from the observed one leaves some residual extended emission to the SW and some `negative' emission levels to the NE."
+" Therefore. in a second step. the surface brielitucss profile was ceutere at a = ]hTUOSS, 6 = INIT and re-fit after cussion from 3383 was removed. within a segment."," Therefore, in a second step, the surface brightness profile was centered at $\alpha$ = $^{\rm h}$ $^{\rm m}$ 22.8, $\delta$ = 7 and re-fit after emission from 383 was removed within a segment."
+" In this case. we obtain Sy=2.67105m 202 3=Q38235003.""TES. aud à,=—.73!12SN kpe (errors are lo)."," In this case, we obtain $S_{\rm 0} = 2.67\,10^{-3}$ $^2$, $\beta = 0.38\pm{0.03}$, and $r_{\rm c} = 73^{+18}_{-15}$ kpc (errors are $\sigma$ )."
+ The only slieht change of the fit parameters uuderlines the robustuess of the fit., The only slight change of the fit parameters underlines the robustness of the fit.
+ The beta-model fit to the surface xiehtness profile is shown in Fie. 2.., The beta-model fit to the surface brightness profile is shown in Fig. \ref{profile}.
+ Again constructing a model image aud subtracting this from the observed one now does not show auv residual cussion except the poiut-like sources., Again constructing a model image and subtracting this from the observed one now does not show any residual emission except the point-like sources.
+ This model is also used for the gas and total mass estimate below., This model is also used for the gas and total mass estimate below.
+ These results differ from the previous analysis of T97 in vieldiug a smaller core radius (they derived re = 230 kpe) and shallower slope (they found ο) = 0.6) for the data prescutec here., These results differ from the previous analysis of T97 in yielding a smaller core radius (they derived $r_{\rm c}$ = 230 kpc) and shallower slope (they found $\beta$ = 0.6) for the data presented here.
+ We have checked that our results are robust (9 varving between  0.3 and 0.15 in extreme cases) against nou-removal of individual sources (up to all). shifts of the central position. aud to svsteniatic exclusion of individual bius (6.2.. the iuuer or outer ones) before fitting the surface briglhtuess profile.," We have checked that our results are robust $\beta$ varying between $\sim$ 0.3 and 0.45 in extreme cases) against non-removal of individual sources (up to all), shifts of the central position, and to systematic exclusion of individual bins (e.g., the inner or outer ones) before fitting the surface brightness profile."
+ Our best fit leads to a arecr ratio of gas to total mass at large τας]., Our best fit leads to a larger ratio of gas to total mass at large radii.
+" It also has consequences when comparing the pressure of the IGM with the nou-thermal pressure of the radio jet (Sect,", It also has consequences when comparing the pressure of the IGM with the non-thermal pressure of the radio jet (Sect.
+ 5.1.3)., 5.1.3).
+reaction with rate coefficients of 7 is replaced by. the following five reactions with rates 7. ///2. Μο./t /kf. and &K/2. respectively: In (his way. statistical expansion is (he main cause of increasing the imunber of species and reactions. “,"reaction with rate coefficients of $k$ is replaced by the following five reactions with rates $k$ , $k/2$, $k/2$, $k/2$, and $k/2$, respectively: In this way, statistical expansion is the main cause of increasing the number of species and reactions. """
+Statistically” means no fractionation occurs.,"Statistically"" means no fractionation occurs."
+ In other swords. all the isotope ratios. are equal to the elemental r 5 [C/C— ratioB and all the isotopomerB. ratios. are unit.v.," In other words, all the isotope ratios are equal to the elemental $^{12}$ $^{13}$ C] ratio and all the isotopomer ratios are unity."
+ secondly. we include the gas phase reactions in Table 1 referring to the discussion in section 2.2.," Secondly, we include the gas phase reactions in Table 1 referring to the discussion in Section 2.2."
+ The rate coefficients of these reactions are not measured in the laboratory., The rate coefficients of these reactions are not measured in the laboratory.
+ So we assume (hat (he rate coelficients are 1 x 1 cg? 1 for neutral-neutral reactions and I x ? em* 1! for ion-neutral reactions., So we assume that the rate coefficients are 1 $\times$ $^{-10}$ $^3$ $^{-1}$ for neutral-neutral reactions and 1 $\times$ $^{-9}$ $^3$ $^{-1}$ for ion-neutral reactions.
+ The dependenceof our results on these rate coefficients is discussed in section 4.3., The dependenceof our results on these rate coefficients is discussed in section 4.3.
+All three FIR maps clearly show the spiral arm structure of the galaxy and a large number of distinct sources.,All three FIR maps clearly show the spiral arm structure of the galaxy and a large number of distinct sources.
+ In addition to the spiral structure. an extended underlying component can be recognized. which becomes more significant with increasing wavelength.," In addition to the spiral structure, an extended underlying component can be recognized, which becomes more significant with increasing wavelength."
+ This is illustrated by the intensity profiles in Fig., This is illustrated by the intensity profiles in Fig.
+ 2. showing cuts through the three FIR maps at a position angle of1157.. along the minor axis.," 2, showing cuts through the three FIR maps at a position angle of, along the minor axis."
+ The structure of the profiles is identical for the three filters. except that at the edge of the galaxy the 170m profile seems to decay slower than the others.," The structure of the profiles is identical for the three filters, except that at the edge of the galaxy the $\mu$ m profile seems to decay slower than the others."
+ The total flux density mEof M33 is ol560Jy. 1250Jy.- and 2200]y at 60. 100. and 170m. respectively.," The total flux density of M33 is 560Jy, 1250Jy, and 2200Jy at 60, 100, and $\mu$ m, respectively."
+"change in spectral shape is seen as the source goes from the NB to the FB, where the disk temperature and luminosity increase significantly when modeled by the model.","change in spectral shape is seen as the source goes from the NB to the FB, where the disk temperature and luminosity increase significantly when modeled by the model."
+ The observations of DiSalvoetal.(2002) cover mostly the HB and the NB., The observations of \citet{dfbfk02} cover mostly the HB and the NB.
+ These authors suggest that the inner rim of the accretion disk the neutron star surface as the source moves from approachesthe HB to the NB., These authors suggest that the inner rim of the accretion disk approaches the neutron star surface as the source moves from the HB to the NB.
+" In an attempt to further understand the continuum spectral evolution ofX-2,, we examined the Swift//XRT spectra from our monitoring campaign."," In an attempt to further understand the continuum spectral evolution of, we examined the /XRT spectra from our monitoring campaign."
+" In the 0.6-10 keV band, a good fit to the XRT data can be achieved using a two thermal component model (disk blackbody and a blackbody)."," In the 0.6-10 keV band, a good fit to the XRT data can be achieved using a two thermal component model (disk blackbody and a blackbody)."
+" For this energy range, no additional power-law component is required."," For this energy range, no additional power-law component is required."
+" For the photoelectric absorption, we fix the column density to Ny=1.9x10?! cm”."," For the photoelectric absorption, we fix the column density to $N_H = 1.9\times10^{21}$ $^{-2}$."
+" This is consistent with values determined from both HIobservations), and from previous fits to the (e.g.DiSalvoetal.2002)."," This is consistent with values determined from both HI, and from previous fits to the X-ray spectra \citep[e.g.][]{dfbfk02}."
+". With this model, we X-rayconsistentlyspectra find temperatures for the disk blackbody and blackbody components of around ~0.5 and ~1.0—1.5 keV respectively."," With this model, we consistently find temperatures for the disk blackbody and blackbody components of around $\sim0.5$ and $\sim1.0-1.5$ keV respectively."
+ Note that these temperatures are lower than seen by the previous spectral studies of significantlydiscussed above., Note that these temperatures are significantly lower than seen by the previous spectral studies of discussed above.
+" In order to address this issue, we searched for any observations that were simultaneous with any of our XRT observations."," In order to address this issue, we searched for any observations that were simultaneous with any of our XRT observations."
+" We found that our observation 002, orbit 1 was overlapping with observation 93443-01-01-15."," We found that our observation 002, orbit 1 was overlapping with observation 93443-01-01-15."
+" The XRT observation was performed on 2008-07-02 from 23:26 to 23:46, whereas the observation ran from 23:00 to 23:38 on the same date."," The XRT observation was performed on 2008-07-02 from 23:26 to 23:46, whereas the observation ran from 23:00 to 23:38 on the same date."
+" We extracted the RXTE//PCA spectrum from PCU only (the most reliable PCU), using the standard goodtime filtering2 and deadtime corrections."," We extracted the /PCA spectrum from PCU 2 only (the most reliable PCU), using the standard goodtime filtering and deadtime corrections."
+" We use the Standard 2 mode data, applying a systematic error of to each channel of the spectrum (wefollowthesamemethodasCackettetal.2009b,fortheRXTE datareduction)."," We use the Standard 2 mode data, applying a systematic error of to each channel of the spectrum \citep[we follow the same method as][for the \rxte{}
+ data."
+" First, we have fit the two thermal component model to the XRT data in the 0.6-10 keV "," First, we have fit the two thermal component model to the XRT data in the 0.6-10 keV energy range."
+"We find an inner disk temperature of kTi;=0.35energy+0.01 range.keV and blackbody temperature of kTyy,=1.0020.01 keV (all errors at the lo level).", We find an inner disk temperature of $kT_{\mathrm{in}} = 0.35\pm0.01$ keV and blackbody temperature of $kT_{\mathrm{bb}} = 1.00\pm0.01$ keV (all errors at the $1\sigma$ level).
+" However, when fitting the same to the RXTEspectrum in the 3-23 keV band, we find KT,=1.28+ keV and kT,=2.08£0.03 keV, similar to what has been Observed in previous studies ofX-2.."," However, when fitting the same to the spectrum in the 3-23 keV band, we find $kT_{\mathrm{in}}=1.28\pm0.01$ keV and $kT_{\mathrm{bb}} = 2.08\pm0.03$ keV, similar to what has been observed in previous studies of."
+ The higher energy coverage of RXTE//PCA is much better suited to constrain these thermal components which have a peak energy of 4keV and ~6keV respectively.," The higher energy coverage of /PCA is much better suited to constrain these thermal components which have a peak energy of $\sim4\,\mathrm{keV}$ and $\sim6\,\mathrm{keV}$ respectively."
+" To test the source of this discrepancy, we have fit the XRT and spectra jointly in the 0.6-23 keV "," To test the source of this discrepancy, we have fit the XRT and spectra jointly in the 0.6-23 keV range."
+We tie all model parameters in the absorbed two thermal range.component model between the two data sets., We tie all model parameters in the absorbed two thermal component model between the two data sets.
+ We also add a constantfactor to allow for any absolute flux calibration mismatch., We also add a constantfactor to allow for any absolute flux calibration mismatch.
+" The resulting spectral fit is shown in Figure 6,, and is clearly a bad fit, with a reduced X? of 3.5."," The resulting spectral fit is shown in Figure \ref{fig:jointspec}, and is clearly a bad fit, with a reduced $\chi^2$ of 3.5."
+" The best-fitting model returns temperatures consistent with the parameters we find when fitting the data alone, and therange."," The best-fitting model returns temperatures consistent with the parameters we find when fitting the data alone, and the."
+". However, there are large residuals present below 3 keV. After trying a wide range of other models (including combinations of disk blackbody, and Comptonization) we were unable to find a blackbody,good fit."," However, there are large residuals present below 3 keV. After trying a wide range of other models (including combinations of disk blackbody, blackbody, power-law and Comptonization) we were unable to find a good fit."
+" power-lawAllow Nz to float as a free parameter did improve the fit, but not to an acceptable level, and also returned a very low value (~0.9x10?! cm”)."," Allow $N_H$ to float as a free parameter did improve the fit, but not to an acceptable level, and also returned a very low value $\sim
+0.9\times10^{21}\,\mathrm{cm}^{-2}$ )."
+" By fitting the XRT data alone we are able to minimize the residuals below 3 keV. However, an extrapolation of the XRT spectrum to higher energies significantly underpredicts the hard X-rays, as shown in Figure 7.."," By fitting the XRT data alone we are able to minimize the residuals below 3 keV. However, an extrapolation of the XRT spectrum to higher energies significantly underpredicts the hard X-rays, as shown in Figure \ref{fig:xtrapol}."
+" In general, the 10-20 keV spectrum of NSXRBs can be well described by a 2-3 keV blackbody (Cackettetal.2009a)."," In general, the 10-20 keV spectrum of NSXRBs can be well described by a 2-3 keV blackbody \citep{cmbbb09}."
+". The XRT fit is dominated by the spectral shape of the soft X-rays, where the response function "," The XRT fit is dominated by the spectral shape of the soft X-rays, where the response function peaks."
+"Therefore, the fit to the XRT spectrum alone, in which the peaks.temperature of the blackbody component is grossly underestimated, cannot be properly extrapolated to the 10-20 keV range."," Therefore, the fit to the XRT spectrum alone, in which the temperature of the blackbody component is grossly underestimated, cannot be properly extrapolated to the 10-20 keV range."
+ We have checked that the difference between the XRT and spectra does not appear to be a cross-calibration issue between the two instruments., We have checked that the difference between the XRT and spectra does not appear to be a cross-calibration issue between the two instruments.
+" As a confirmation, we analyzed several near-simultaneous XTE and observations of the black hole candidate LMXB during its 2006 outburst (seeRykoff observations). The Swift"," As a confirmation, we analyzed several near-simultaneous XTE and observations of the black hole candidate LMXB during its 2006 outburst \citep[see][for details on all \swift{}
+ ."
+//XRT of this object can be fit by a disk blackbody plus a spectrapower-law over a wide range in simpleluminosity (Rykoffetal. 2007).., The /XRT spectra of this object can be fit by a simple disk blackbody plus a power-law over a wide range in luminosity \citep{rmst07}. .
+" When looking at near-simultaneous observations, we find that apart from a slight offset"," When looking at near-simultaneous observations, we find that apart from a slight offset"
+depends on resolution. spectral tvpe of the star and. the continuum fitting.,"depends on resolution, spectral type of the star and the continuum fitting."
+ For stars of spectral tvpe late A- E. the equivalent widths could be measured with an accuracy of ~ 5 -8 in absence of line asvmametries.," For stars of spectral type late A- F, the equivalent widths could be measured with an accuracy of $\sim$ 5 -8 in absence of line asymmetries."
+ With the above mentioned accuracy in. measurecl equivalent widths. the microturbulence velocity could: be measured with an accuracy of £0.25 km temperature with Ε 1501 ancl of £0.25 em s," With the above mentioned accuracy in measured equivalent widths, the microturbulence velocity could be measured with an accuracy of $\pm$ 0.25 km $^{-1}$ , temperature with $\pm$ 150K and of $\pm$ 0.25 cm $^{-2}$."
+ Dut for cooler stars. line strengths. could. only be measured with accuracy of 8-10%.," But for cooler stars, line strengths could only be measured with accuracy of 8-10."
+. Vhe microturbulence velocity could be measured. with an accuracy. of 40.5 km 5 temperature with Επ] ancg of 40.5 em s," The microturbulence velocity could be measured with an accuracy of $\pm$ 0.5 km $^{-1}$, temperature with $\pm$ 250K and of $\pm$ 0.5 cm $^{-2}$."
+ The derived. atmospheric parameters and heliocentric racial velocities for the epoch of observations for the program stars are presented in Table 2., The derived atmospheric parameters and heliocentric radial velocities for the epoch of observations for the program stars are presented in Table 2.
+" The sensitivity of the derived. abuncances to. the uncertainties of atmospheric parametersg. and £ are summarized in ""Table 3."," The sensitivity of the derived abundances to the uncertainties of atmospheric parameters, and $\xi$ are summarized in Table 3."
+ For four stars representing the full temperature range of our sample. we present changes in XfFeo] caused by varving atmospheric parameters by 20019. 0.25 cm s and 0.5 km (average accuracies of these parameters) with respect to the chosen model for cach star.," For four stars representing the full temperature range of our sample, we present changes in [X/Fe] caused by varying atmospheric parameters by 200K, 0.25 cm $^{-2}$ and 0.5 km $^{-1}$ (average accuracies of these parameters) with respect to the chosen model for each star."
+ The total stellar. parameter related: error. is. estimated by taking the square root of the sum of the square of the svstematic errors (individual errors associated with uncertainties in temperatures and eravities) and bars in the abundance plot correspond to the systematic error., The total stellar parameter related error is estimated by taking the square root of the sum of the square of the systematic errors (individual errors associated with uncertainties in temperatures and gravities) and bars in the abundance plot correspond to the systematic error.
+ llaving determined the atmospheric parameters. the abundances of dillerent. elements were derived. using the available lines.," Having determined the atmospheric parameters, the abundances of different elements were derived using the available lines."
+ The derived abuncances relative to the solar abundances are presented in their respective tables., The derived abundances relative to the solar abundances are presented in their respective tables.
+ The solar photospherie abundances given by Xsplund. Cirevesse Sauval (2005) have been used as reference values.," The solar photospheric abundances given by Asplund, Grevesse Sauval (2005) have been used as reference values."
+ The sources of thegf values for dillerent elements used in our abundance analysis have been listed in Table 4., The sources of the values for different elements used in our abundance analysis have been listed in Table 4.
+ We have investigated possible svstematic ellects caused by the adoptedgf values from cülferent. sources as follows., We have investigated possible systematic effects caused by the adopted values from different sources as follows.
+ We have measured. solar equivalent widths for our lines on Solar Flux Atlas by Ixurucz et al. (, We have measured solar equivalent widths for our lines on Solar Flux Atlas by Kurucz et al. (
+"1984) and estimated the abuncdances using model atmosphere appropriate for the sun Taj10.g o£ 44 and £j, o£ 0.9 1.","1984) and estimated the abundances using model atmosphere appropriate for the sun =5770, of 4.4 and $\xi_{t}$ of 0.9 $^{-1}$."
+" ""Phe agreement for most elements is within 0.04 to 0.09 dex although the linelist is different for cilferent stars.", The agreement for most elements is within 0.04 to 0.09 dex although the linelist is different for different stars.
+ The lines of certain elements are alfected by hvperfine splitting., The lines of certain elements are affected by hyperfine splitting.
+ We have included. the hvperline structure components in our line list while synthesizing the spectral feature of these elements for deriving the abundances., We have included the hyperfine structure components in our line list while synthesizing the spectral feature of these elements for deriving the abundances.
+ For elements Sc and Mp. we have used hls component list and theirgf given by Prochaska & MeWilliam (2000). for Cu (Allen & Porto de Mello 2011). for Eu (Mucciarelli οἱ al.," For elements Sc and Mn, we have used hfs component list and their given by Prochaska $\&$ McWilliam (2000), for Cu (Allen $\&$ Porto de Mello 2011), for Eu (Mucciarelli et al."
+ 2008) and for Ba (AleWilliam 1998)., 2008) and for Ba (McWilliam 1998).
+ “Phe effect. of isotopic components were considered. for Eu Lh (2 isotopes) anc Ba LL (5 isotopes) lines., The effect of isotopic components were considered for Eu II (2 isotopes) and Ba II (5 isotopes) lines.
+ on-LTE corrections lor CNO elements. show strong emperature dependence (particularly for N see e.g. Lyubimkov et al., Non-LTE corrections for CNO elements show strong temperature dependence (particularly for N see e.g. Lyubimkov et al.
+ 2011. Schiller Przvbilla 2008) and they also vary from multiplet to multiplet.," 2011, Schiller Przybilla 2008) and they also vary from multiplet to multiplet."
+ Venn (1995) have abulated these corrections for € and N for a range of stellar emperatures for lines belonging to dillerent multiplets., Venn (1995) have tabulated these corrections for C and N for a range of stellar temperatures for lines belonging to different multiplets.
+ For B-l stars the correction for N varies from — 0.3 dex to 1.0 dex while for € itis OL to 0.5 dex., For B-F stars the correction for N varies from $-$ 0.3 dex to $-$ 1.0 dex while for C it is $-$ 0.1 to $-$ 0.5 dex.
+ ‘Takeda and “Takacla-Lliclai (1998) have calculated Non-LTE corrections or oxvgen abundance using lines for A-E stars and correction varies between 0.1 to 0.4 dex., Takeda and Takada-Hidai (1998) have calculated Non-LTE corrections for oxygen abundance using lines for A-F stars and correction varies between $-$ 0.1 to $-$ 0.4 dex.
+ The neglect. of departure from. LYE also. introduces errors in the estimated abundances of heavier elements., The neglect of departure from LTE also introduces errors in the estimated abundances of heavier elements.
+ These errors for a given clement vary with the. stellar temperatures and also on metallicities., These errors for a given element vary with the stellar temperatures and also on metallicities.
+ In. very metal-poor stars the over-ionisation and in some cases usage of resonance lines for abundance determination results in errors as large as [0.5 dex for elements like Na., In very metal-poor stars the over-ionisation and in some cases usage of resonance lines for abundance determination results in errors as large as $+0.5$ dex for elements like Na.
+ We have only used: subordinates lines for sodium. believed to be formed in deeper lavers and for them Non-ETI corrections of about 0.10 dex are reported (Lind et al.," We have only used subordinates lines for sodium, believed to be formed in deeper layers and for them Non-LTE corrections of about $-0.10$ dex are reported (Lind et al."
+ 2011. Gehren et al.," 2011, Gehren et al."
+ 2004)., 2004).
+ Gohren et al have estimated. Non-LTI corrections lor Ale and. Al and for the lines used in our analysis for à sample of CG stars covering a range in metallicities., Gehren et al have estimated Non-LTE corrections for Mg and Al and for the lines used in our analysis for a sample of G stars covering a range in metallicities.
+ For thick disc metallicities Non-L'EE corrections are. |0.05 for Mg/Ee] and [0.2 for ΑΙΓΟ is reported., For thick disc metallicities Non-LTE corrections are $+0.05$ for [Mg/Fe] and $+$ 0.2 for [Al/Fe] is reported.
+ The Non-LTE corrections for lines have been estimated by IxXorotin (2009)., The Non-LTE corrections for lines have been estimated by Korotin (2009).
+ For the lines used in our work (6042 6056. 6743 6758 X.) the elfect is negligible.," For the lines used in our work (6042 $-$ 6056, 6743 $-$ 6758 ), the effect is negligible."
+ Wecemever (2001) has done Non-LYE calculation for lines and Non-L'TIS correction in range | 0.01 to 0.05 have been estimated for the Sun and Vega respectively., Wedemeyer (2001) has done Non-LTE calculation for lines and Non-LTE correction in range $-$ 0.01 to $-$ 0.05 have been estimated for the Sun and Vega respectively.
+ Mashonkina. korn and. Przvbilla (2007) have studied: departure from LPE over a range of stellar parameter and. meallicities for a large number of and lines.," Mashonkina, Korn and Przybilla (2007) have studied departure from LTE over a range of stellar parameter and meallicities for a large number of and lines."
+ Atsolar ormoderately deficient. metallicities the Non-LTI correction is smaller than | 0.1 dex for lines used in our analysis for deriving Ca abundance., Atsolar ormoderately deficient metallicities the Non-LTE correction is smaller than $+$ 0.1 dex for lines used in our analysis for deriving Ca abundance.
+ Also the correction is negligible for very weal features., Also the correction is negligible for very weak features.
+is the Alfvén velocity square.,is the $\acute{e}$ n velocity square.
+ If the angular momentum is transported by an Alfvén wave. and the magnetic field B dominates as an azimuthal field. so that the larger the radio VES. the more efficient the angular momentum transport.," If the angular momentum is transported by an $\acute{e}$ n wave, and the magnetic field B dominates as an azimuthal field, so that the larger the radio $v^{2}_{rA}$ $v^{2}_{A}$, the more efficient the angular momentum transport."
+ The code originally written by Paezyriski (1969:; 1970)) was updated by Sienkiewiez in 1995 and Yang et al. (2001))., The code originally written by $\acute{n}$ ski \cite{Paczynski69}; \cite{Paczynski70}) ) was updated by Sienkiewicz in 1995 and Yang et al. \cite{Yang}) ).
+ We modified it to incorporate the hydrostatic effects of rotation on the equations of stellar structure. using the method of Kippenhahn-Meynet (Kippenhahn Thomas 1970:: Meynet Maeder 1997).," We modified it to incorporate the hydrostatic effects of rotation on the equations of stellar structure, using the method of Kippenhahn-Meynet (Kippenhahn Thomas \cite{Kippenhahn70}; Meynet Maeder \cite{Meynet}) )."
+ The models are calculated using the OPAL equation of state (Rogers et al. 1996)).," The models are calculated using the OPAL equation of state (Rogers et al. \cite{Rogers}) ),"
+ OPAL opacity (Iglesias Rogers 1996)). and the Alexander Ferguson (1994)) opacity table for low temperature.," OPAL opacity (Iglesias Rogers \cite{Iglesias}) ), and the Alexander Ferguson \cite{Alexander}) ) opacity table for low temperature."
+ Element diffusion is incorporated for helium and metals (Thoul et al. 1994))., Element diffusion is incorporated for helium and metals (Thoul et al. \cite{Thoul}) ).
+ The nuclear reaction rates have been updated according to Baheall Pinsonneault (1995))., The nuclear reaction rates have been updated according to Bahcall Pinsonneault \cite{Bahcall}) ).
+ Energy transfer by convection is treated according to the standard mixing length theory. and the boundaries of the convection zones are determined by the Schwarzschild criterion.," Energy transfer by convection is treated according to the standard mixing length theory, and the boundaries of the convection zones are determined by the Schwarzschild criterion."
+" We adopt the solar age as 46x10? year. luminosity Lo=3.844«10? ergs7!. radius R-=6.96x10!"" em. and the ratio of heavy elements to hydrogen by mass Z/X=0.023 (Grevesse Sauval 1998))."," We adopt the solar age as $\times 10^{9}$ year, luminosity $L_{\odot}$ $\times10^{33}$ $s^{-1}$, radius $R_{\odot}$ $\times 10^{10}$ cm, and the ratio of heavy elements to hydrogen by mass $Z/X$ =0.023 (Grevesse Sauval \cite{Grevesse}) )."
+ As mentioned above. we assume that solid body rotation was eforced in the convective region.," As mentioned above, we assume that solid body rotation was enforced in the convective region."
+ This assumption has been used by Pinsonneault et al (1989)). Chaboyer et al (1995)). and Huang (2004b))," This assumption has been used by Pinsonneault et al \cite{Pinsonneault}) ), Chaboyer et al \cite{Chaboyer95}) ), and Huang \cite{Huang04b}) )."
+ The total initial angular momentum is a free parameter that is adjusted until the surface velocities of the solar-age models are near the solar surface velocity., The total initial angular momentum is a free parameter that is adjusted until the surface velocities of the solar-age models are near the solar surface velocity.
+" We take the rotation rate of the Sun as a solid-body rotation, about 2.72 x107° rad/s (Komm et al. 2003))."," We take the rotation rate of the Sun as a solid-body rotation, about 2.72 $\times10^{-6}$ rad/s (Komm et al. \cite{Komm}) ),"
+ as the rotation rate of convective zone of our model., as the rotation rate of convective zone of our model.
+ The strength and spatiotemporal distribution of magnetic fields inside the star are poorly known., The strength and spatiotemporal distribution of magnetic fields inside the star are poorly known.
+ Classical dynamo models predict toroidal fields that are not stronger than about 10 Gauss., Classical dynamo models predict toroidal fields that are not stronger than about $10^{4}$ Gauss.
+ But Choudhuri Gilman (1987)). D'Silva Choudhur (1993)). and Caligart et al. (1995))," But Choudhuri Gilman \cite{Choudhuri}) ), D'Silva Choudhuri \cite{D'Silva}) ), and Caligari et al. \cite{Caligari}) )"
+ have pointed out the value of the magnetic field at the bottom of convective zone as around 10° Gauss., have pointed out the value of the magnetic field at the bottom of convective zone as around $10^{5}$ Gauss.
+ The virial theorem (Parker 1979)) sets an upper limit to the magnitude of the average solar magnetic field: (B)<105 Gauss., The virial theorem (Parker \cite{Parker}) ) sets an upper limit to the magnitude of the average solar magnetic field: $\langle B\rangle \leq 10^{8}$ Gauss.
+ Dudorov et al. (1989)), Dudorov et al. \cite{Dudorov}) )
+ estimated the value of the poloidal magnetic field in the solar radiative zone to be about the order of unity., estimated the value of the poloidal magnetic field in the solar radiative zone to be about the order of unity.
+" Using a kinematic model with prescribed internal rotation and a standard solar model. Fox and Bernstein (1987)) investigated the existence of large-scale magnetic fields in the Sun. and found that the ratio of poloidal to toroidal components of the magnetic field is about 10110,"," Using a kinematic model with prescribed internal rotation and a standard solar model, Fox and Bernstein \cite{Fox}) ) investigated the existence of large-scale magnetic fields in the Sun, and found that the ratio of poloidal to toroidal components of the magnetic field is about $10^{-10}$."
+" Although the ratio of B, to B may change with the radial coordinate r. the rotation rate Q. and the time ¢. the relation of 5 tor. Q. and ris unclear."," Although the ratio of $B_{r}$ to $B$ may change with the radial coordinate $r$, the rotation rate $\Omega$, and the time $t$, the relation of $\frac{B_{r}}{B}$ to $r$, $\Omega$, and $t$ is unclear."
+" As a first test. we take B,/B to be à constant. and take the order of the ratio of B, to Bas 107 in the radiative region."," As a first test, we take $B_{r}/B$ to be a constant, and take the order of the ratio of $B_{r}$ to $B$ as $10^{-5}$ in the radiative region."
+ In order to study the effects of magnetic fields. we construct three different types of solar models.," In order to study the effects of magnetic fields, we construct three different types of solar models."
+ These models are labelled as follows: Some parameters are summarized in Table 1.., These models are labelled as follows: Some parameters are summarized in Table \ref{Parat}.
+ The length parameters a. Zo. Xo. fo. and f. are free parameters.," The mixing-length parameters $\alpha$ , $Z_{0}$, $X_{0}$, $f_{\Omega}$, and $f_{c}$ are free parameters."
+" The Χο. Z) and « are adjusted until the solar-age model has the values of the solar luminosity. radius. and (Z/X),."," The $X_{0}$ , $Z_{0}$ and $\alpha$ are adjusted until the solar-age model has the values of the solar luminosity, radius, and $(Z/X)_{s}$."
+ Then Yo is determined by Yo=|—Xo-Za., Then $Y_{0}$ is determined by $Y_{0}=1-X_{0}-Z_{0}$.
+ The parameter fo was adjusted to fit the radial profiles of the angular velocity. which were got from helioseismology in the solar interior.," The parameter $f_{\Omega}$ was adjusted to fit the radial profiles of the angular velocity, which were got from helioseismology in the solar interior."
+ Parameter Ff. was adjusted to get the best radial profile of the sound speed., Parameter $f_{c}$ was adjusted to get the best radial profile of the sound speed.
+" Finally. Ἐν. (Z/X), and Rj,- are the results of calculation."," Finally, $Y_{s}$, $(Z/X)_{s}$ and $R_{bcz}$ are the results of calculation."
+ All models are constructed by evolving a fully convective. pre-main-sequence. one solar mass model to the age of the present Sun.," All models are constructed by evolving a fully convective, pre-main-sequence, one solar mass model to the age of the present Sun."
+ We define the angular momentum density J(r. €) and total angular momentum A(t) as where fis time.," We define the angular momentum density J(r, t) and total angular momentum A(t) as where $t$ is time."
+ In Figs., In Figs.
+ laa and Ibb. we present the rotation rates and their differentiation with respect to radius r as a function r at the age of 4.6 G years.," \ref{omj}a a and \ref{omj}b b, we present the rotation rates and their differentiation with respect to radius $r$ as a function $r$ at the age of 4.6 G years."
+ The radial profile ofthe rotationrate of model M3 is flatter than of model M2., The radial profile ofthe rotationrate of model M3 is flatter than of model M2.
+ In our models. the radial gradient of angular velocity is negative.," In our models, the radial gradient of angular velocity is negative."
+ But the helioseismology has revealed that the gradient Is positive at, But the helioseismology has revealed that the gradient is positive at
+260 IX temperature (Gehrz&Smith1999).. the energy densitv in LR photons which serve as target radiation field is μμ=8x10!eVem. several orders of magnitude higher than the density of the cosmic microwave background (0.27eVem.I) or typical ISM photon fields.,"260 K temperature \citep{EtaCar:Gehrz99}, the energy density in IR photons which serve as target radiation field is $U_{\rm rad} = 8\times10^4~{\rm eV~cm}^{-3}$, several orders of magnitude higher than the density of the cosmic microwave background $0.27~{\rm eV~cm}^{-3}$ ) or typical ISM photon fields."
+ For infrared target photons the Ixlein-Nishina (XN) suppression of the IC cross-section [or electron energies £x100 GeV is a less than effect. ancl is not considered in the following., For infrared target photons the Klein-Nishina (KN) suppression of the IC cross-section for electron energies $E \leq 100$ GeV is a less than effect and is not considered in the following.
+ The relevant IC cooling time for electrons in the blast wave region is therefore Thom334000vrsElo., The relevant IC cooling time for electrons in the blast wave region is therefore $\tau_{\rm loss} \approx 4000~{\rm yrs}~E_{\rm GeV}^{-1}$.
+ In a 104G field. the svnchrotron-Ioss timescale for 100 GeV. electrons is of O((109) vears: IC losses will dominate unless B>2mCG.," In a $10\,\mu{\rm G}$ field, the synchrotron-loss timescale for 100 GeV electrons is of $10^6$ ) years: IC losses will dominate unless $B>2{\rm mG}$."
+ In dense media. Dremsstrahlung may also contribute significantly to the energy-loss rate of electrons.," In dense media, Bremsstrahlung may also contribute significantly to the energy-loss rate of electrons."
+" The estimated mean density of 400—500 particles per  implies a typical Bremsstrvahling loss time of νους, ", The estimated mean density of $400-500$ particles per $^{-3}$ implies a typical Bremsstrahlung loss time of $\tau_{\rm Brems} \approx 8-11\times 10^4$ years.
+Therelore. at > GeV energies. IC emission is verv likely to be the dominant enission process for electrons in this svstenm.," Therefore, at $>$ GeV energies, IC emission is very likely to be the dominant emission process for electrons in this system."
+ The maximum electron energy is determined (lor cooling-limitecd acceleration) by the balance between radiative energve. losses and the rate of energvONS gain via DSA. when Thlom=Tree!acc This maximum energv for the measured blast wave speed is consistent. with the HE ατα emission found by for moderate magnetic fields of 10iQ and diffusion close to the Bohm limit.," The maximum electron energy is determined (for cooling-limited acceleration) by the balance between radiative energy losses and the rate of energy gain via DSA, when $\tau_{\rm loss} = \tau_{\rm acc}$: This maximum energy for the measured blast wave speed is consistent with the HE -ray emission found by for moderate magnetic fields of $10\,\mu{\rm G}$ and diffusion close to the Bohm limit."
+ Moreover. the enerev-loss timescale is much shorter than the age of the remnant. implving that acceleration of electrons is indeed cooling limited in (his svstem.," Moreover, the energy-loss timescale is much shorter than the age of the remnant, implying that acceleration of electrons is indeed cooling limited in this system."
+ In (his scenario most of (he energy injected into the acceleration process emerges as IC radiation in (he GeV domain., In this scenario most of the energy injected into the acceleration process emerges as IC radiation in the GeV domain.
+" For such an efficient process (he required energv input in electrons is modest: within an order of magnitude of Ll,~10 erg.", For such an efficient process the required energy input in electrons is modest: within an order of magnitude of $L_{\gamma} t_{\rm sys} \sim 10^{45}$ erg.
+" IC. emission is expected to dominate over p-p emission in this svstem as long as E,>10οι. Van", IC emission is expected to dominate over p-p emission in this system as long as $E_{e}>10^{-3} E_{\rm CR}$.
+nonietal.(2009). performed numerical caleulations where (μον studied acceleration and radiation of electrons in an radiation-dominated environment such as the one investigated in (his work., \citet{Vannoni09} performed numerical calculations where they studied acceleration and radiation of electrons in an radiation-dominated environment such as the one investigated in this work.
+ Compared to the magnetic field of 107G used here. the authors obtain for magnetic fields of 100CG maximum particle energies in the TeV regime. where INN effects become verv significant.," Compared to the magnetic field of $10\,{\mu \rm G}$ used here, the authors obtain for magnetic fields of $100\,{\mu \rm G}$ maximum particle energies in the TeV regime, where KN effects become very significant."
+The Sun's meridional circulation has been a part of solar magnetic dvnamo theory. [or half à century.,The Sun's meridional circulation has been a part of solar magnetic dynamo theory for half a century.
+ A poleward meridional flow from the Sun's mid latitudes was invoked in the earliest models (long belore the How was actually measured) to transport magnetic elements from decaving sunspot regions to the poles where (μον would erode the opposite polarity magnetic field from the old sunspot evele and build up the polar fields of the new sunspot evele (Babcock1961)., A poleward meridional flow from the Sun's mid latitudes was invoked in the earliest models (long before the flow was actually measured) to transport magnetic elements from decaying sunspot regions to the poles where they would erode the opposite polarity magnetic field from the old sunspot cycle and build up the polar fields of the new sunspot cycle \citep{Babcock61}.
+.. This surface meridional How. along with its latitudinal structure ancl variation in (ime. is now well observed (Topkaetal.1932:Konan1993:Lathawayelal.1996:Hathaway&Rightmire2010.2011) and its role in the surface magnetic [lux (ransport is well established (DeVore&SheeleyLOST:vanDallegooijenscehrijver&Tide 2001).," This surface meridional flow, along with its latitudinal structure and variation in time, is now well observed \citep{Topka_etal82, Komm_etal93, Hathaway_etal96, HathawayRightmire10, HathawayRightmire11} and its role in the surface magnetic flux transport is well established \citep{DeVoreSheeley87, vanBallegooijen_etal98, SchrijverTitle01}."
+ Dynamo theories over (he last decade and a half have assumed that the mass traveling poleward in the surface lavers sinks inward at (he poles and returns to the equator along the base of the Sun's convection zone al a depth of 200 Mm., Dynamo theories over the last decade and a half have assumed that the mass traveling poleward in the surface layers sinks inward at the poles and returns to the equator along the base of the Sun's convection zone at a depth of 200 Mm.
+ In these theories this slow. dense. equatorward flow is responsible for the equatorward drilt of sunspot," In these theories this slow, dense, equatorward flow is responsible for the equatorward drift of sunspot"
+from its fitted range of (to cover its complete bandpass.,from its fitted range of to cover its complete bandpass.
+ Hence there is some extra uncertainty in the intensities. of lines with wavelengths &reater than numeasured from this light., Hence there is some extra uncertainty in the intensities of lines with wavelengths greater than measured from this flight.
+ However. no such problem exists with the SIZIUES.89 measurements of these lines.," However, no such problem exists with the SERTS–89 measurements of these lines."
+ We have searched. for emission lines. in. the SERTSS9 and SERTS95 spectra using the detections of Thomas Neupert (1994) anc Brosius οἱ al. (, We have searched for emission lines in the SERTS–89 and SERTS–95 spectra using the detections of Thomas Neupert (1994) and Brosius et al. (
+1998). supplemented with. those [rom other sources. including the NIST. the latest version (V5.2) of the database (Dore ct al.,"1998b), supplemented with those from other sources, including the NIST the latest version (V5.2) of the database (Dere et al."
+ LOOT: Landi ct al., 1997; Landi et al.
+ 2006). the Atomic Line List of van and in particular the excellent summary of line identifications by Del Zanna οἱ al. (," 2006), the Atomic Line List of van and in particular the excellent summary of line identifications by Del Zanna et al. ("
+2004).,2004).
+ Phe latter provides not only a comprehensive list of wavelengths. for wellobserved transitions. but also indicates alternative wavelengths where previous identilications are not consistent with their conclusions.," The latter provides not only a comprehensive list of wavelengths for well-observed transitions, but also indicates alternative wavelengths where previous identifications are not consistent with their conclusions."
+ In ‘Table 1 we list the transitions found in the SERESSO and SERTS spectra. alone with their measured. wavelengths.," In Table 1 we list the transitions found in the SERTS–89 and SERTS--95 spectra, along with their measured wavelengths."
+ We also. indicate possible blending features. or alternative identifications as suggested. by Thomas Neupert or 3rosius et al., We also indicate possible blending features or alternative identifications as suggested by Thomas Neupert or Brosius et al.
+ in their original line lists., in their original line lists.
+ lntensities and line widths (EWIIM) of the features are given in Tables 2 and 3 for the SERTS89 and SERS active regions. respectively. alone with the associated lo errors.," Intensities and line widths (FWHM) of the features are given in Tables 2 and 3 for the SERTS–89 and SERTS--95 active regions, respectively, along with the associated $\sigma$ errors."
+ These were measured using moclilied versions of the Gaussian fitting routines emploved by Thomas reupert (1994). as discussed by Ixeenan οἱ al. (," These were measured using modified versions of the Gaussian fitting routines employed by Thomas Neupert (1994), as discussed by Keenan et al. ("
+2007).,2007).
+ As a consequence. the intensities. PWIA values and. their uncertainties listed in Tables 2 and 3 are somewhat cillerent rom those originally reported in Thomas Neupert and Drosius et al. (," As a consequence, the intensities, FWHM values and their uncertainties listed in Tables 2 and 3 are somewhat different from those originally reported in Thomas Neupert and Brosius et al. ("
+19950).,1998b).
+ Also. a uniform [actor of 1.24 has »en applied here to all SERSSO intensities. rellecting a more recent re-evaluation of its absolute calibration scale.," Also, a uniform factor of 1.24 has been applied here to all SERTS–89 intensities, reflecting a more recent re-evaluation of its absolute calibration scale."
+ Even so. in all clirectly comparable cases. the resulting ine intensity values usually diller only slightly. from those oeviouslv obtained.," Even so, in all directly comparable cases, the resulting line intensity values usually differ only slightly from those previously obtained."
+ For the SERTS95 data set. several of he stronger first-order lines could also be detected. and heir measurements are therefore included in Table 3 along with the seconc-order results.," For the SERTS–95 data set, several of the stronger first-order lines could also be detected, and their measurements are therefore included in Table 3 along with the second-order results."
+ However. for the SERTS89 spectrum only the first-order features could be reliably identified.," However, for the SERTS–89 spectrum only the first-order features could be reliably identified."
+ In Figs 15 we plot portions of the SERTS89 and SERTS95 spectra containing transitions. focusing on emission lines which have not previously been identified in SERTS data sets.," In Figs 1–5 we plot portions of the SERTS–89 and SERTS–95 spectra containing transitions, focusing on emission lines which have not previously been identified in SERTS data sets."
+ We note that several of these features have line intensities ancl widths comparable to the noise Iluctuations., We note that several of these features have line intensities and widths comparable to the noise fluctuations.
+ In these instances. the reality of the feature was confirmed by a visual inspection of the original SIZICES film.," In these instances, the reality of the feature was confirmed by a visual inspection of the original SERTS film."
+ llowever. the lines are weak aad clearly further observations to strengthen these identifications would be desirable.," However, the lines are weak and clearly further observations to strengthen these identifications would be desirable."
+ The model ion for consisted. of the 90 energeticallv lowest. fine-structure levels., The model ion for consisted of the 90 energetically lowest fine-structure levels.
+ These are listed in table 1 of Agearwal Keenan (2005). ancl comprise levels. arising from the Αρ”. 3s3p. δρα. Bs8p73el and 373p Sel configurations.," These are listed in table 1 of Aggarwal Keenan (2005), and comprise levels arising from the $^{2}$ $^{5}$, $^{6}$, $^{2}$ $^{4}$ 3d, $^{5}$ 3d and $^{2}$ $^{3}$ $^{2}$ configurations."
+ Where available. energies for these levels," Where available, energies for these levels"
+Now that we have established the empirical. modified-wind momentum luminosity relation for the three galaxies. we can determine the M(Z) relation.,"Now that we have established the empirical modified-wind momentum luminosity relation for the three galaxies, we can determine the $\mdot(Z)$ relation."
+ Before doing so. we first inter compare the results for these three environments. and confront them with predictions.," Before doing so, we first inter compare the results for these three environments, and confront them with predictions."
+ In Fig., In Fig.
+ 4 the empirical modified-wind momentum relations determined for the total observed samples (solid lines) are shown alongside the predicted relations of ?? (dotted lines).," \ref{fig:wlr-glob} the empirical modified-wind momentum relations determined for the total observed samples (solid lines) are shown alongside the predicted relations of \cite{vink00, vink01}
+ (dotted lines)."
+ The top. middle and bottom lines of each line style. respectively. correspond to the Galactic. LMC and SMC observed and predicted relations.," The top, middle and bottom lines of each line style, respectively, correspond to the Galactic, LMC and SMC observed and predicted relations."
+ To facilitate a meaningful comparison one sigma confidence intervals for the observed WLRs are shown as grey areas., To facilitate a meaningful comparison one sigma confidence intervals for the observed WLRs are shown as grey areas.
+ First focusing on the left-hand side. showing the relations without a clumping correction. we see that the empirical relations are clearly separated from each other beyond the fitting uncertainties.," First focusing on the left-hand side, showing the relations without a clumping correction, we see that the empirical relations are clearly separated from each other beyond the fitting uncertainties."
+ For the entire investigated luminosity range. the relative separations of the empirical WLRs agree well with the separations predicted by ?.. although the empirical WLRs do show a systematic offset compared to the theoretical results.," For the entire investigated luminosity range, the relative separations of the empirical WLRs agree well with the separations predicted by \cite{vink01}, although the empirical WLRs do show a systematic offset compared to the theoretical results."
+ Focusing for a moment on the relative separations only: measured at logLy/L=5.75. which coincides with the region where the fit uncertainties are the smallest. the offsets between the empirical and theoretical relations relative to the Galaxy are: 0.28 and 0.25 for the LMC. and 0.62 and 0.57 for the SMC.," Focusing for a moment on the relative separations only: measured at $\log \lstar/\lsun = 5.75$, which coincides with the region where the fit uncertainties are the smallest, the offsets between the empirical and theoretical relations relative to the Galaxy are: $0.28$ and $0.25$ for the LMC, and $0.62$ and $0.57$ for the SMC."
+ The systematic offset between observations and theory 1s of the order of 0.2 dex., The systematic offset between observations and theory is of the order of 0.2 dex.
+ This suggests thatpredictions. the empirical mass-loss rates would be overestimated by at most a factor of ~ two.," This suggests that, the empirical mass-loss rates would be overestimated by at most a factor of $\sim$ two."
+ Indeed. when turning to the clumping corrected empirical relations. the relative offset almost disappears.," Indeed, when turning to the clumping corrected empirical relations, the relative offset almost disappears."
+ This is shown on the right-hand side of Fig. 4..," This is shown on the right-hand side of Fig. \ref{fig:wlr-glob},"
+ where we show the wind momentum relations obtained from the total observed samples accounting for this correction., where we show the wind momentum relations obtained from the total observed samples accounting for this correction.
+ Note that. strictly speaking. this agreement pertains only to stars showing ΠΠ emission. as the clumping corrections ts designed such as to scale the ones that do show eemission to those that do not (see Sect. 2?)).," Note that, strictly speaking, this agreement pertains only to stars showing in emission, as the clumping corrections is designed such as to scale the ones that do show emission to those that do not (see Sect. \ref{sec:gal}) )."
+ The slopes of the empirical relations. however. are slightly flatter. as the clumping correction preferentially affects the high-luminosity objects.," The slopes of the empirical relations, however, are slightly flatter, as the clumping correction preferentially affects the high-luminosity objects."
+" To calculate the mass loss metallicity dependence we use the relative separation between the empirical modified-wind momentum relations at logL,/L.«=5.75. for which the fit uncertainty in all three relations is at its minimum."," To calculate the mass loss metallicity dependence we use the relative separation between the empirical modified-wind momentum relations at $\log \lstar/\lsun = 5.75$, for which the fit uncertainty in all three relations is at its minimum."
+ Note that the differential slopes also allow for a luminosity dependent Z-dependence. however. based on the uncertainties we doubt whether this would be meaningful.," Note that the differential slopes also allow for a luminosity dependent $Z$ -dependence, however, based on the uncertainties we doubt whether this would be meaningful."
+ Moreover. for the SMC it was shown that this slope is very sensitive to the rather uncertain low luminosity domain.," Moreover, for the SMC it was shown that this slope is very sensitive to the rather uncertain low luminosity domain."
+ In Tab., In Tab.
+ 3. the power law indices for the observed M(Z) dependence for the LMC and SMC relative to the Galaxy are listed., \ref{tab:wlr-par} the power law indices for the observed $\mdot(Z)$ dependence for the LMC and SMC relative to the Galaxy are listed.
+ These were calculated, These were calculated
+We conclude that among all various correlations. only the FIR-RC and FIB-IICN correlations stand out clistinguishably better than the rest.,"We conclude that among all various correlations, only the FIR-RC and FIR-HCN correlations stand out distinguishably better than the rest."
+ Here. we plot the well-known FIR-RC correlation in Figure 6 for our HCN sample.," Here, we plot the well-known FIR-RC correlation in Figure 6 for our HCN sample."
+ Similar to Figure l(a). AGNs are labeled in order to see whether there is anv obvious trend in the difference between. AGNs and other galaxies.," Similar to Figure 1(a), AGNs are labeled in order to see whether there is any obvious trend in the difference between AGNs and other galaxies."
+ In particular. we plot other directly related correlations such as FIRRC versus HCN and. FIR/RC versus HCN/CO (Figure 7) in order to examine any svstematic contribution to the possible seatters in the tightest FIR-RC correlation.," In particular, we plot other directly related correlations such as FIR/RC versus HCN and FIR/RC versus HCN/CO (Figure 7) in order to examine any systematic contribution to the possible scatters in the tightest FIR-RC correlation."
+ The obvious trend in Figure 7(a). in which FIR/RC versus WON emission is plotted. is (hat smaller scatter in the FIR/RC ratio (thus a bit tighter correlation) can be observed ab the lower ICN luminosity end. ie. the normal galaxies and less luminous LIRGs since all ULIRGs have large ICN luminosity.," The obvious trend in Figure 7(a), in which FIR/RC versus HCN emission is plotted, is that smaller scatter in the FIR/RC ratio (thus a bit tighter correlation) can be observed at the lower HCN luminosity end, i.e., the normal galaxies and less luminous LIRGs since all ULIRGs have large HCN luminosity."
+ However. for ULIRGs and starburst galaxies (i.e.. ealaxies wilh 0.06. as claimed by GS04a). there is an obvious (rend of larger scatters in (he FIR/RC ratio lor higher ION/CO ratio (Fieure 1(b)) or high luminosities in ICN or even CO.," However, for ULIRGs and starburst galaxies (i.e., galaxies with $>0.06$, as claimed by GS04a), there is an obvious trend of larger scatters in the FIR/RC ratio for higher HCN/CO ratio (Figure 7(b)) or high luminosities in HCN or even CO."
+ The situation is quite similar if we show the IR. and RC! luminosities in the v-anis., The situation is quite similar if we show the IR and RC luminosities in the $x$ -axis.
+ In fact. the FIR/RC ratio is not entirely constant and seems to weakly correlate with the IICN/BRC (Figure 4(b)). vet is only marginally dependent on the CO/RC ratio (Figure 5(a)).," In fact, the FIR/RC ratio is not entirely constant and seems to weakly correlate with the HCN/RC (Figure 4(b)), yet is only marginally dependent on the CO/RC ratio (Figure 5(a))."
+ Again. this appears to suggest that the extremely large scatter in the FIR/RC ratio only exisis in galaxies wilh either (he highest or smallest ICN/RC ratios.," Again, this appears to suggest that the extremely large scatter in the FIR/RC ratio only exists in galaxies with either the highest or smallest HCN/RC ratios."
+ In Section 3.3.1. we only listed the three-parzmeter (FIR. HCN. RC) fits.," In Section 3.3.1, we only listed the three-parameter (FIR, HCN, RC) fits."
+ The details of the (hree-parameter (FIR. ILCN. CO) fits can be found in GSO4a. which manifest a much lighter linear FIR-IICN correlation versus the nonlinear less tight FIR-CO correlation.," The details of the three-parameter (FIR, HCN, CO) fits can be found in GS04a, which manifest a much tighter linear FIR-HCN correlation versus the nonlinear less tight FIR-CO correlation."
+ For the sake of completeness. we list here (wo other combinations of three-parameter fits involving RC. HCN. and CO and FIR. RC. and CO.," For the sake of completeness, we list here two other combinations of three-parameter fits involving RC, HCN, and CO and FIR, RC, and CO."
+ First. RC. ICN. and CO: The contributing [actor from HCN (0.72) to RC is nearly (vice as that [rom CO (0.40).," First, RC, HCN, and CO: The contributing factor from HCN (0.72) to RC is nearly twice as that from CO (0.40)."
+ This is much less than the factor of 25 difference between HCN and CO in the fit to FIR (GS04a)., This is much less than the factor of $>$ 5 difference between HCN and CO in the fit to FIR (GS04a).
+Z=(0.000886 was run.,$Z=0.000886$ was run.
+ In this model the star reached. the horizontal branch but the entire envelope was lost by the ene of the E-AGB., In this model the star reached the horizontal branch but the entire envelope was lost by the end of the E-AGB.
+ The mass-loss on the E-ACGB increases as the value of Y goes up because the core mass &oes up., The mass-loss on the E-AGB increases as the value of $Y$ goes up because the core mass goes up.
+ The luminosity of an AGB star will depend stronely on the core mass and this increased bIuminosity drives the increased niass-loss., The luminosity of an AGB star will depend strongly on the core mass and this increased luminosity drives the increased mass-loss.
+ This model produces a core mass of 0.60NL;.," This model produces a core mass of $\,{\rm
+ M}_{\sun}$."
+ 1 the mass is lowered to 0.75AL. then the model once again produces a LB star but all the remaining envelope mass is lost on the INCID.," If the mass is lowered to $0.75\,{\rm M}_{\sun}$ then the model once again produces a HB star but all the remaining envelope mass is lost on the E-AGB."
+ This lower mass produces a 0.58AL. remnant.," This lower mass produces a $0.58\,{\rm
+ M}_{\sun}$ remnant."
+ Why is this significant?, Why is this significant?
+ This is close to the inferred mass of CJJC-I's CSPN and production of a hvdrogen free nebula requires getting rid of the initial hydrogen before the PN phase., This is close to the inferred mass of GJJC-1's CSPN and production of a hydrogen free nebula requires getting rid of the initial hydrogen before the PN phase.
+ Getting rid of all the envelope. means if this star experiences a thermal pulse then there would be little to no hydrogen ejectecl ancl the star would. eject helium rich material and it would look much like GJIC-1., Getting rid of all the envelope means if this star experiences a thermal pulse then there would be little to no hydrogen ejected and the star would eject helium rich material and it would look much like GJJC-1.
+ This suggestion should be regarded: as speculative ancl the positive evidence for it is thin but the idea seems to be possible and merits additional study., This suggestion should be regarded as speculative and the positive evidence for it is thin but the idea seems to be possible and merits additional study.
+ Without a detailed. population study whieh is beyond. the scope of this paper it is only possible to show the numbers work out approximately correct., Without a detailed population study which is beyond the scope of this paper it is only possible to show the numbers work out approximately correct.
+ L start by assuming the Jacobyetal.(L997) statement that given the total luminosity of the GC system then the number of PNe should be 16., I start by assuming the \citet{jac} statement that given the total luminosity of the GC system then the number of PNe should be 16.
+ Wit is assumed that 70. percent of globular cluster stars are primordial and have a turn-olf in the O.80-0.90AL range then this reduces the number to about 4-5 since these stars produce none.," If it is assumed that 70 percent of globular cluster stars are primordial and have a turn-off in the $\,{\rm M}_{\sun}$ range then this reduces the number to about 4-5 since these stars produce none."
+ The number from this appears to be 3 which is in rough agreement., The number from this appears to be 3 which is in rough agreement.
+ Further. assuming all blue stragelers produce a visible PN. we assume about. 1 from this part which matches the designation of Ps 1 (IN648) as being produced by this channel.," Further, assuming all blue stragglers produce a visible PN, we assume about 1 from this part which matches the designation of Ps 1 (K648) as being produced by this channel."
+ The rest will be produced by second generation stars (which may or may not all produce a visible PN)., The rest will be produced by second generation stars (which may or may not all produce a visible PN).
+ Pherefore the numbers are roughly consistent with the number of known Όλα, Therefore the numbers are roughly consistent with the number of known PNe.
+ What do Jalu 1 and Jabu 2 tell us about second generation GC stars?, What do JaFu 1 and JaFu 2 tell us about second generation GC stars?
+ They confirm. independently of the use of colour magnitude diagrams. stars exist in globular clusters with )—0280.33 and it suggests these second. generation stars are a significant fraction of the number of GC stars.," They confirm, independently of the use of colour magnitude diagrams, stars exist in globular clusters with $Y=0.28-0.33$ and it suggests these second generation stars are a significant fraction of the number of GC stars."
+ These two Όλο could turn out to be very important. since they allow the direct observation of elements which can not be observed directly. using stellar spectroscopy., These two PNe could turn out to be very important since they allow the direct observation of elements which can not be observed directly using stellar spectroscopy.
+ For Jaku 1 logN/O=0.52 (Jacoby οἱ al.), For JaFu 1 $\log{\rm N/O}=-0.52$ (Jacoby et al.)
+ which is slightly higher than logN/O.;=0.88 CAsplund.Cirevesse.&Sauval 2005)., which is slightly higher than $\log{\rm N/O}_{\sun}=-0.88$ \citep{asp05}.
+. This ratio is consistent with both a first dredge-up event ancl possibly a small amount of nitrogen enrichment or oxvecn depletion., This ratio is consistent with both a first dredge-up event and possibly a small amount of nitrogen enrichment or oxygen depletion.
+ For Jaku l logS/O=1.35 which is consistent with the solar logS/O=1.50.," For JaFu 1 $\log{\rm S/O}=-1.35$ which is consistent with the solar $\log{\rm
+ S/O}=-1.50$."
+ From these abundance ratios either all these have been enriched bv the same relative amount or the star has not. been enriched relative to a first. generation star in Pal 6., From these abundance ratios either all these have been enriched by the same relative amount or the star has not been enriched relative to a first generation star in Pal 6.
+ The sulfer abundance of Jabu Lis given by0., The sulfer abundance of JaFu 1 is given by.
+"55. ""Phis is consistent with a cluster with metallicity between 1/3 and 1/2 solar.", This is consistent with a cluster with metallicity between 1/3 and 1/2 solar.
+ Wo the clusters metallicity is on the lower end of its range then this would. indicate all elements have been enhanced in this cluster., If the cluster's metallicity is on the lower end of its range then this would indicate all elements have been enhanced in this cluster.
+ Jabu 2 max indicate a depletion of oxvgen in NGC 6441., JaFu 2 may indicate a depletion of oxygen in NGC 6441.
+ Phe value of οσοlogO/1L;=0.98 and the value of logAr/llολες=—0.72., The value of $\log{\rm O/H}-\log{\rm O/H}_{\sun}=-0.93$ and the value of $\log{\rm Ar/H}-\log{\rm Ar/H}_{\sun}=-0.72$.
+ Since this is a lower metallicity cluster we would expect. oxygen as an alpha clement to be enhanced. however it appears to. be less enhanced. relative to argon., Since this is a lower metallicity cluster we would expect oxygen as an alpha element to be enhanced however it appears to be less enhanced relative to argon.
+ TFhis would. be consistent with the enhanced material having been processed by CNO cvcling., This would be consistent with the enhanced material having been processed by CNO cycling.
+ As GCs age they leak stars into the field via collisions and also by tidal stripping., As GCs age they leak stars into the field via collisions and also by tidal stripping.
+ In fact it is estimated. that of halo stars are 2nd. generation. elobular cluster stars (Vesperiniet.al.2010)., In fact it is estimated that of halo stars are 2nd generation globular cluster stars \citep{ves10}.
+. If second. generation. halo. stars become halo stars they could. produce PN similar to «αι 1 and Jabu 2 and may be €JJC€-1., If second generation halo stars become halo stars they could produce PN similar to JaFu 1 and JaFu 2 and may be GJJC-1.
+ However. this is a small fraction of the halo stars and it is quite possible no such star has been observed.," However, this is a small fraction of the halo stars and it is quite possible no such star has been observed."
+ Other places to look for similar PN might be the svstem. of satellite galaxies., Other places to look for similar PN might be the system of satellite galaxies.
+ The GCs w Centauri. M54 ancl M22 are all possible captured. satellite galaxies ancl all of them have multiple populations (see Piotto 2009).," The GCs $\omega$ Centauri, M54 and M22 are all possible captured satellite galaxies and all of them have multiple populations (see Piotto 2009)."
+ This suggests multiple populations might be common in satellite galaxies and this model predicts there may be PNe in these satellite ealaxies similar to the GC PNe., This suggests multiple populations might be common in satellite galaxies and this model predicts there may be PNe in these satellite galaxies similar to the GC PNe.
+" This paper presents à series of models for the expected TI""-AGB stars in globular clusters.", This paper presents a series of models for the expected TP-AGB stars in globular clusters.
+ The results of these models are compared. to the observed. abundances of the elobular cluster PNe and the measured masses of WDs., The results of these models are compared to the observed abundances of the globular cluster PNe and the measured masses of WDs.
+Intriguingly. we note that all the runs at very high resolution (VR) even show a small but systematie decrease of the cold gas fraction.,"Intriguingly, we note that all the runs at very high resolution (VR) even show a small but systematic decrease of the cold gas fraction."
+ Including feedback from winds has the twofold effect of suppressing the resolution-dependence of the cumulative efficiency and to make the SFR history almost independent of resolution at 2X 2-3 (see the right panel of Fig. 35., Including feedback from winds has the twofold effect of suppressing the resolution–dependence of the cumulative efficiency and to make the SFR history almost independent of resolution at $z\mincir 2$ –3 (see the right panel of Fig. \ref{fi:sfr_profs_res}) ).
+ The enhanced star formation activity at high-redshift of the high resolution runs produces more efficient gas pre-heating. which has the effect of inhibiting star ormation at later times.," The enhanced star formation activity at high–redshift of the high resolution runs produces more efficient gas pre–heating, which has the effect of inhibiting star formation at later times."
+ This presumably explains that the star 'ormation rate of the VR run at 252 is even below those of the ower-resolution runs., This presumably explains that the star formation rate of the VR run at $z\mincir 2$ is even below those of the lower–resolution runs.
+ Assuming that the smallest resolved halos where cooling can take place contain  100 DM particles. we tind their escape velocity to be of order kis.+ for the resolution (LR) runs. and a factor of about three smaller for the VR runs of£1.," Assuming that the smallest resolved halos where cooling can take place contain $\sim$ 100 DM particles, we find their escape velocity to be of order $\vel$ for the low--resolution (LR) runs, and a factor of about three smaller for the VR runs of."
+. Therefore. the wind velocity is always larger than he escape valocity of the smallest halos which are first resolved at high redshift. thereby implying that our feedback scheme is efficient in preventing star formation already in the first generation of resolved galaxies.," Therefore, the wind velocity is always larger than the escape valocity of the smallest halos which are first resolved at high redshift, thereby implying that our feedback scheme is efficient in preventing star formation already in the first generation of resolved galaxies."
+ An alternative explanation for the decrease of the cold fraction in the VR runs could be that the suppression of star formation in the highest resolution run is related to some numerical effect., An alternative explanation for the decrease of the cold fraction in the VR runs could be that the suppression of star formation in the highest resolution run is related to some numerical effect.
+ For instance. one may argue that an improved resolution provides a more accurate description of the gas behaviour at cooling interface.," For instance, one may argue that an improved resolution provides a more accurate description of the gas behaviour at cooling interface."
+ Indeed. a coarse description of this interface is expected to cause spurious gas cooling(2).. a feature that should however be weak in the entropy conserving formulation of SPH implemented in our code.," Indeed, a coarse description of this interface is expected to cause spurious gas cooling, a feature that should however be weak in the entropy conserving formulation of SPH implemented in our code."
+ However. if this was really the case. one would expect the same effect to appear also in runs without galactic winds (NW runs).," However, if this was really the case, one would expect the same effect to appear also in runs without galactic winds (NW runs)."
+ But the left panel of Fig., But the left panel of Fig.
+ 3. clearly demonstrates that this is not the case., \ref{fi:sfr_profs_res} clearly demonstrates that this is not the case.
+ Hence. the stable behaviour of star formation with increasing resolution is more likely related to the inhibiting effect of more efficient high-z: feedback onto later generations of gulaxies.," Hence, the stable behaviour of star formation with increasing resolution is more likely related to the inhibiting effect of more efficient $z$ feedback onto later generations of galaxies."
+ The number of identified galaxies within each cluster increases with better resolution and with the cluster müss. as expected.," The number of identified galaxies within each cluster increases with better resolution and with the cluster mass, as expected."
+ As shown in the right panel of Figure 2.. the number of identified galaxies grows by more than one order of magnitude when passing from the LR to the VR runs.," As shown in the right panel of Figure \ref{fi:fstar_res}, the number of identified galaxies grows by more than one order of magnitude when passing from the LR to the VR runs."
+ In line with the behaviour of the stellar fraction. the number of galaxies for the CL4 cluster increases by a factor of about three when the wind feedback is switched off. almost independent of the resolution.," In line with the behaviour of the stellar fraction, the number of galaxies for the CL4 cluster increases by a factor of about three when the wind feedback is switched off, almost independent of the resolution."
+ We defer a detailed analysis of the stellar mass function of the identitied galaxies and of the diffuse stellar component as a function of resolution to forthcoming work (Murante et al..," We defer a detailed analysis of the stellar mass function of the identified galaxies and of the diffuse stellar component as a function of resolution to forthcoming work (Murante et al.,"
+ in preparation)., in preparation).
+ In Figure 4.. we summarize the results on the star fraction by plotting it as a function of the cluster virial mass. for the reference runs of both (filled circles) and of all the 20 clusters of (squares).," In Figure \ref{fi:fstar_mvir}, we summarize the results on the star fraction by plotting it as a function of the cluster virial mass, for the reference runs of both (filled circles) and of all the 20 clusters of (squares)."
+ Both simulation sets confirm a trend of decreasing star fraction as a function of the cluster mass., Both simulation sets confirm a trend of decreasing star fraction as a function of the cluster mass.
+ However. clusters belonging to show a stellar fraction which is systematically higher than that of#1.. the difference being larger than any possible object-to-object intrinsic scatter induced by the varying dynamical histories of different clusters.," However, clusters belonging to show a stellar fraction which is systematically higher than that of, the difference being larger than any possible object-to-object intrinsic scatter induced by the varying dynamical histories of different clusters."
+" However. we recall that the two sets of cluster simulations differ in the strength of the adopted feedback and in the normalization of the power spectrum,"," However, we recall that the two sets of cluster simulations differ in the strength of the adopted feedback and in the normalization of the power spectrum."
+ To investigate the effect of the different feedback. we have run additional simulations of CLS and CL6 of by increasing the wind speed to the same value as used for (SW runs).," To investigate the effect of the different feedback, we have run additional simulations of CL5 and CL6 of by increasing the wind speed to the same value as used for (SW runs)."
+ The results for the stellar fraction are shown with tilled triangles in Fig. 4.., The results for the stellar fraction are shown with filled triangles in Fig. \ref{fi:fstar_mvir}.
+ Although increasing the feedback efficiency produces a significant suppression of f.. the effect is still not large enough to fully account for the difference between the two simulation sets.," Although increasing the feedback efficiency produces a significant suppression of $f_*$, the effect is still not large enough to fully account for the difference between the two simulation sets,"
+evervihing (he SDSS classifies as a galaxy: spectroscopy. shows that this selection includes some stars.,everything the SDSS classifies as a galaxy; spectroscopy shows that this selection includes some stars.
+ Among the remaining 822 objects without spectroscopy. we expect ~5% stars.," Among the remaining 822 objects without spectroscopy, we expect $\sim 5\%$ stars."
+ The integral completeness is thus e79%. Figure 1 shows that the differential completeness drops to e40% αἱ r=21.3., The integral completeness is thus $\sim 79\%$ Figure \ref{fig:completeness.subaru.histo.ps} shows that the differential completeness drops to $\sim 40\%$ at $ r = 21.3$.
+ In Figure 2 the open circles show the positions of galaxies within the primary sample with (upper panel) and without (lower panel) a redshilt., In Figure \ref{fig:nozandz.ps} the open circles show the positions of galaxies within the primary sample with (upper panel) and without (lower panel) a redshift.
+ The pattern of incompleteness is a natural consequence of the fiber positioning constraints., The pattern of incompleteness is a natural consequence of the fiber positioning constraints.
+ Figure 2. also shows (he positions of the 6 robust convergence nap peaks (see Section 3.1. and Table 3))., Figure \ref{fig:nozandz.ps} also shows the positions of the 6 robust convergence map peaks (see Section \ref{newmap} and Table \ref{tbl:VDisp}) ).
+ Note that the numerical designations of the robust peaks are not sequential because (here are peaks of intervening rank that are not secure (Section 3.1))., Note that the numerical designations of the robust peaks are not sequential because there are peaks of intervening rank that are not secure (Section \ref{newmap}) ).
+ Our secondary (Gargets ave (1) galaxies wilh r>21.3. 6—22»0.4. g—r>1.0 ranked by central surface brightness and (ii) bluer objects with r<21.3.," Our secondary targets are (i) galaxies with $r > 21.3$, $r-i > 0.4$, $g-r > 1.0$ ranked by central surface brightness and (ii) bluer objects with $r < 21.3$."
+ We observe the objects with >21.3 beginning with the highest surface brightness objects., We observe the objects with $r > 21.3$ beginning with the highest surface brightness objects.
+ The lower panel of Figure 3 shows the distribution on the skv of the 546 faint red galaxies with redshifts., The lower panel of Figure \ref{fig:faintandblue.ps} shows the distribution on the sky of the 546 faint red galaxies with redshifts.
+ The upper panel shows the distribution on the skv of the 974 bluer objects with a redshift., The upper panel shows the distribution on the sky of the 974 bluer objects with a redshift.
+ The faint red objects are reasonably uniform across the field: the blue objects are concentrated toward the center of the field as a result of (he positioning of the IHectospec fields., The faint red objects are reasonably uniform across the field; the blue objects are concentrated toward the center of the field as a result of the positioning of the Hectospec fields.
+ Covering the edges of (he field is relatively inefficient because (he 1 degree diameter Hectospec field overlaps the boundary., Covering the edges of the field is relatively inefficient because the 1 degree diameter Hectospec field overlaps the boundary.
+ Figure 4 shows the distribution of apparent magnitude. pepo. for all of the galaxies with measured redshifts.," Figure \ref{fig:rVSz.subaru.ps} shows the distribution of apparent magnitude, $r_{petro}$ , for all of the galaxies with measured redshifts."
+ Different symbols indicate the three subsamples., Different symbols indicate the three subsamples.
+ At every redshilt Z0.7 the redshift survey samples more than a one magnitude range in ry., At every redshift $\lesssim 0.7$ the redshift survey samples more than a one magnitude range in $r_{petro}$.
+ Figure 5 shows the redshift distribution for the entire GTO2dee? field survey and for the primary (arget sample., Figure \ref{fig:redshift.histogram.ps} shows the redshift distribution for the entire $^2$ field survey and for the primary target sample.
+ The median redshift for the primary target. sample of 3025 &alaxies is 0.336: lor the entire sample. the median redshift is 0.349.," The median redshift for the primary target sample of 3025 galaxies is 0.336; for the entire sample, the median redshift is 0.349."
+ The galaxies wilh r>21.3 populate the high redshift range ancl the SDSS galaxies fill in the survey at low redshift., The galaxies with $r > 21.3$ populate the high redshift range and the SDSS galaxies fill in the survey at low redshift.
+ The sharply defined peaks are the expected signature of the large-scale structure of the universe., The sharply defined peaks are the expected signature of the large-scale structure of the universe.
+ Figure 6 shows cone diagrams [or (he survey., Figure \ref{fig:cone.diagram.subaru.ps} shows cone diagrams for the survey.
+ Obvious voids are delineated by walls and filaments., Obvious voids are delineated by walls and filaments.
+ The contrast between the voids and the surrounding structure is enhanced relative to a redshift without color selection for red galaxies., The contrast between the voids and the surrounding structure is enhanced relative to a redshift without color selection for red galaxies.
+ The red galaxies are more clustered and preferentially populate more dense regions., The red galaxies are more clustered and preferentially populate more dense regions.
+ Fingers extending along the corresponding to massive clusters of galaxies are evident by eve al z=0.415 and al 2=0.540., Fingers extending along the line-of-sight corresponding to massive clusters of galaxies are evident by eye at $z = 0.415$ and at $z = 0.540$.
+ In Section 4.2 we show that these two prominent svstems correspond (o (lietwo most signilicant weak lensing peaks., In Section \ref{clusters} we show that these two prominent systems correspond to thetwo most significant weak lensing peaks.
+Recdshilts are crucial for iclentifving candidate svstems that correspond to weak lensing,Redshifts are crucial for identifying candidate systems that correspond to weak lensing
+2010).,.
+". In particular, the Hercules moving group, a group of stars significantly offset from the bulk of the observed velocity distribution 1)mainBodyCitationEnd46|Dehnen98b,Bovy09a, displays a wide range of ages and metallicities 2010)."," In particular, the Hercules moving group, a group of stars significantly offset from the bulk of the observed velocity distribution , displays a wide range of ages and metallicities ."
+" A compelling model has the Hercules stream originating through resonant interactions of stars in the outer disk with the bar in the central region of the Galaxy 2001), but there is no confirmation of this picture beyond the locally observed stellar kinematics that it was proposed to explain."," A compelling model has the Hercules stream originating through resonant interactions of stars in the outer disk with the bar in the central region of the Galaxy , but there is no confirmation of this picture beyond the locally observed stellar kinematics that it was proposed to explain."
+" An investigation of the Hercules members’ color-magnitude diagram found no evidence for any significant metallicity anomaly with respect to other local stars 2010),, even though such a difference should be expected as the Hercules stars originate from a few kiloparsec toward the Galactic center where the average metallicity is higher than in the Solar neighborhood."," An investigation of the Hercules members' color–magnitude diagram found no evidence for any significant metallicity anomaly with respect to other local stars , even though such a difference should be expected as the Hercules stars originate from a few kiloparsec toward the Galactic center where the average metallicity is higher than in the Solar neighborhood."
+" The absence of the predicted metallicity anomaly could potentially be explained,e.g.,, through radial mixing 2002),, but it is clear that chemodynamical modeling of the Galactic disk is not currently up to the task of separating out these effects 2008)."," The absence of the predicted metallicity anomaly could potentially be explained, through radial mixing , but it is clear that chemodynamical modeling of the Galactic disk is not currently up to the task of separating out these effects ."
+. Here I propose that a robust and unambiguous test of the bar-resonance model for the Hercules stream is a search for the distinct pattern it predicts as we trace the velocity distribution of stars near the Solar radius around the Galaxy., Here I propose that a robust and unambiguous test of the bar-resonance model for the Hercules stream is a search for the distinct pattern it predicts as we trace the velocity distribution of stars near the Solar radius around the Galaxy.
+" iis in a good position to have the final word on this, but I also show that the Hercules feature can be detected in the line-of-sight velocity distribution in selected directions on the sky."," is in a good position to have the final word on this, but I also show that the Hercules feature can be detected in the line-of-sight velocity distribution in selected directions on the sky."
+" Clear signatures are predicted to be detectable in regions a few kiloparsec from the Sun, allowing for a clean verification of the bar-resonance model for the Hercules stream."," Clear signatures are predicted to be detectable in regions a few kiloparsec from the Sun, allowing for a clean verification of the bar-resonance model for the Hercules stream."
+" To simulate the effect of the bar on stellar orbits near the Solar circle, we follow the approach of (2000)."," To simulate the effect of the bar on stellar orbits near the Solar circle, we follow the approach of ."
+". For each fixed position in the Galaxy, we evaluate the velocity distribution function at values for the radial—toward the Galactic center—and tangential—in the direction of Galactic rotation—velocity by backward-integrating these velocities in the Galactic potential to obtain the intial orbit before bar formation."," For each fixed position in the Galaxy, we evaluate the velocity distribution function at values for the radial—toward the Galactic center—and tangential---in the direction of Galactic rotation—velocity by backward-integrating these velocities in the Galactic potential to obtain the intial orbit before bar formation."
+ We posit that at this epoch the old stellar disk was in a steady-state that can be described by a simple distribution function that is a function of the energy and angular momentum of the orbit alone., We posit that at this epoch the old stellar disk was in a steady-state that can be described by a simple distribution function that is a function of the energy and angular momentum of the orbit alone.
+ It follows from the collisionless Boltzmann equation that the value of the velocity distribution at the present position and velocity is equal to that of the steady-state distribution function before, It follows from the collisionless Boltzmann equation that the value of the velocity distribution at the present position and velocity is equal to that of the steady-state distribution function before
+study this line. one has to account for partial frequency redistribution effects (e.g. Uitenbroek Bruls 1992 and Rutten Milkey 1979).,"study this line, one has to account for partial frequency redistribution effects (e.g. Uitenbroek Bruls 1992 and Rutten Milkey 1979)."
+ Here also collisions should play an important role since. besides their effects on the atomic polarization. they could change the frequency of the photons.," Here also collisions should play an important role since, besides their effects on the atomic polarization, they could change the frequency of the photons."
+ It remains a challenge to develop a general theory for partial frequency redistribution of polarized radiation in the presence of arbitrary magnetic fields and including the effects of collisions in à multilevel picture with/without hyperfine structure., It remains a challenge to develop a general theory for partial frequency redistribution of polarized radiation in the presence of arbitrary magnetic fields and including the effects of collisions in a multilevel picture with/without hyperfine structure.
+LAE may be relevant (o secondary pair creation. which requires photon energies of at least an MeV. Third. all particles in the LAEW emit LAE. ancl the associated damping of the LAEW is of potential interest in itsell. providing a simple way of relating the power in LAE io the energy in the LAEWs.,"LAE may be relevant to secondary pair creation, which requires photon energies of at least an MeV. Third, all particles in the LAEW emit LAE, and the associated damping of the LAEW is of potential interest in itself, providing a simple way of relating the power in LAE to the energy in the LAEWs."
+ Finally. there is the possibility of a maser form of LAE as a radio emission mechanism. as discussed brielly in paper 1.," Finally, there is the possibility of a maser form of LAE as a radio emission mechanism, as discussed briefly in paper 1."
+ We use the theory in paper 1 with two notable changes., We use the theory in paper 1 with two notable changes.
+" First. in paper 1 we used primes to denote «quantities in (he primed Irame in which the oscillations are purely temporal: the primed quantities are related to those in the laboratory frame. in which the LAEW have a phase speed v,, bv a Lorentz transformation with velocity c?ορ"," First, in paper 1 we used primes to denote quantities in the primed frame in which the oscillations are purely temporal; the primed quantities are related to those in the laboratory frame, in which the LAEW have a phase speed $v_\phi$ by a Lorentz transformation with velocity $c^2/v_\phi$."
+ In this paper our analvsis is restricted to the Irae in which the oscillations are purely temporal. ancl for convenience in wriling we omil (he primes on all relevant quantities.," In this paper our analysis is restricted to the frame in which the oscillations are purely temporal, and for convenience in writing we omit the primes on all relevant quantities."
+ Second. in paper 1 we concentrated on a triangular wave form. which is an excellent approximation for à LAEW in which the electrons and positrons become highly relativistic. and here we generalize to an arbitrary wave form.," Second, in paper 1 we concentrated on a triangular wave form, which is an excellent approximation for a LAEW in which the electrons and positrons become highly relativistic, and here we generalize to an arbitrary wave form."
+ This allows us to apply our results for LAE more generally., This allows us to apply our results for LAE more generally.
+ We described by the wave form by a periodic function. Z(). of phase y. and derive the orbit of a particle by expanding in inverse powers of its Lorentz factor.," We described by the wave form by a periodic function, $T(\chi)$ , of phase $\chi$, and derive the orbit of a particle by expanding in inverse powers of its Lorentz factor."
+ In Sec., In Sec.
+ 2. we derive the emissivity [or LAE. and in See.," \ref{sect:emissivity} we derive the emissivity for LAE, and in Sec."
+ 3. we evaluate it in terms of Airy integrals., \ref{sect:Airy} we evaluate it in terms of Airy integrals.
+ In Sec., In Sec.
+ 4. we compare our results with the generalized Larmor formula ancl we discuss inconsistencies (hat arise., \ref{sect:Larmor} we compare our results with the generalized Larmor formula and we discuss inconsistencies that arise.
+ We discuss our results ancl the application to pulsars in Sec., We discuss our results and the application to pulsars in Sec.
+ 5. and summarize (he conclusions in Sec. 6.., \ref{sect:discussion} and summarize the conclusions in Sec. \ref{sect:conclusions}.
+ An exact treatment of LAE is involves assuming period motion and expanding in a Fourier series (Rowe1995)., An exact treatment of LAE is involves assuming period motion and expanding in a Fourier series \citep{r95}.
+. As shown in paper 1. the emission is at harmonics. w=sQ. of the frequency of the LAEW.," As shown in paper 1, the emission is at harmonics, $\omega=s\Omega$, of the frequency of the LAEW."
+ For ahighly relativistic particle. very high harmonics dominate.," For ahighly relativistic particle, very high harmonics dominate,"
+Since the first distance bin is assumed to be complete. we can iteratively correct the number counts in the next Α ds the number of stars in bin j. Nds the number of stars in the previous bin (j;— 1). up to the limit given by the bin j. and N is the number of stars from that limit up to the limit given by bin jy—1. as indicated in Fig. Al..,"Since the first distance bin is assumed to be complete, we can iteratively correct the number counts in the next $N_j$ is the number of stars in bin $j$, $N^{'}$ is the number of stars in the previous bin $j-1$ ), up to the limit given by the bin $j$, and $N^{''}$ is the number of stars from that limit up to the limit given by bin $j-1$, as indicated in Fig. \ref{schematicCompCorr}."
+ The division into Nand N. is done to avoid correlated The first bin (;= 1) is assumed to be complete V4)., The division into $N^{'}$ and $N^{''}$ is done to avoid correlated The first bin $j=1$ ) is assumed to be complete $N_1$ ).
+" The corrected number [N57 in the second bin is then and for the third bin Thus this can generally be written as The errors of the number counts ἂν δρα and NV;, are independent of each other and thus Poissonian: oy,=VNj. [Tyr=VN,."," The corrected number $N_2^{{\rm corr}}$ in the second bin is then and for the third bin Thus this can generally be written as The errors of the number counts $N_j$, $N_{j-1}^{'}$ and $N_{j-1}^{''}$ are independent of each other and thus Poissonian: $\sigma_{N_j}=\sqrt{N_j}$, $\sigma_{N_{j-1}^{'}}=\sqrt{N_{j-1}^{'}}$, $\sigma_{N_{j-1}^{''}}=\sqrt{N_{j-1}^{''}}$."
+ Following Gaussian error propagation the error for the corrected number of stars in the second bin is Thus the error of the completeness corrected counts in the j bin is, Following Gaussian error propagation the error for the corrected number of stars in the second bin is and for the third bin Thus the error of the completeness corrected counts in the $j^{th}$ bin is
+"The inner, midplane regions of our disk model are very dense (reaching ~10'* ?), possibly dense enough for three-body (termolecular) reactions to have an effect on the chemistry.","The inner, midplane regions of our disk model are very dense (reaching $\sim$ $^{14}$ $^{-3}$ ), possibly dense enough for three-body (termolecular) reactions to have an effect on the chemistry."
+" In order to investigate the effect on benzene production in particular, we ran a third model including a representative three-body reaction: where two propargyl radicals collide with M, a third molecule."," In order to investigate the effect on benzene production in particular, we ran a third model including a representative three-body reaction: where two propargyl radicals collide with M, a third molecule."
+ This has been shown to be a very efficient formation mechanism for benzene when densities are high; see the discussion of Cherchneffetal.(1992)., This has been shown to be a very efficient formation mechanism for benzene when densities are high; see the discussion of \citet{che92}.
+". We have adopted a rate coefficient from chemical models of the Saturnian atmosphere: (Mosesetal.2000) where k,, the high pressure limit to the reaction rate, was taken to be 1.66x1015 cem ss! and ko, the low pressure limit, 1x10?"" cem? ss! 2000)."," We have adopted a rate coefficient from chemical models of the Saturnian atmosphere: \citep{mos00} where $_\infty$, the high pressure limit to the reaction rate, was taken to be $\times$ $^{-13}$ $^6$ $^{-1}$ and $_0$, the low pressure limit, $\times$ $^{-27}$ $^3$ $^{-1}$ \citep{won00}."
+". Inclusion of this reaction in the model did not result in a significant change in the abundance of gas-phase benzene, despite the presence of large fractional abundances of 04Η. in dense regions."," Inclusion of this reaction in the model did not result in a significant change in the abundance of gas-phase benzene, despite the presence of large fractional abundances of $_3$ $_3$ in dense regions."
+" Even adopting the faster rates for k,, and ko from (1.2x10719 cem? ss“! and 6.0x 1077561680/7 ccm® ss“, respectively) had little overall effect."," Even adopting the faster rates for $_\infty$ and $_0$ from \citet{leb05} $\times$ $^{-10}$ $^3$ $^{-1}$ and $\times$ $^{-28}\mathrm{e}^{1680/T}$ $^6$ $^{-1}$, respectively) had little overall effect."
+ The region in which benzene is abundant in our model (Fig. 4)), The region in which benzene is abundant in our model (Fig. \ref{fig:benzdist}) )
+" is also rich in other organic molecules: high densities mean fast collision timescales with other molecules and with dust grains, and also protection from the strong stellar UV field in this region."," is also rich in other organic molecules: high densities mean fast collision timescales with other molecules and with dust grains, and also protection from the strong stellar UV field in this region."
+" These prime conditions for the growth of organic molecules may allow molecules to grow beyond the single aromatic ring of benzene into PAHs, either by substitution of hydrogen by acetylene in the ring (Wongetal.2000;Moses2000),, or by the dimerisation of benzene (Aj) with successive phenyl rings to produce naphthalene (A3), phenanthrene (A3), pyrene (A4) and successively large PAH molecules (Cherchneff1996)."," These prime conditions for the growth of organic molecules may allow molecules to grow beyond the single aromatic ring of benzene into PAHs, either by substitution of hydrogen by acetylene in the ring \citep{won00,mos00}, or by the dimerisation of benzene $_1$ ) with successive phenyl rings to produce naphthalene $_2$ ), phenanthrene $_3$ ), pyrene $_4$ ) and successively large PAH molecules \citep{che96}."
+". Benzene itself is unlikely to persist through to the formation of planets: calculations show that benzene has a timescale for UV destruction in the diffuse solar system of only hundreds of years (Allainetal. 2005),, meaning that the benzene seen on Jupiter, Saturn and Titan is unlikely to be primordial."," Benzene itself is unlikely to persist through to the formation of planets: calculations show that benzene has a timescale for UV destruction in the diffuse solar system of only hundreds of years \citep{all96,rui05}, meaning that the benzene seen on Jupiter, Saturn and Titan is unlikely to be primordial."
+" In order for PAHs to form, the constituents must be present in a region warm enough for ring closure to occur."," In order for PAHs to form, the constituents must be present in a region warm enough for ring closure to occur."
+ PAHs only seem to form efficiently at temperatures in the region of KK (Frenklach&Feigelson1989) and not greater., PAHs only seem to form efficiently at temperatures in the region of K \citep{fre89} and not greater.
+" Production may be possible,"," Production may be possible,"
+"between these two spectra was then fit with a cut-off power-law model in PHA space (i.e., CUTOFFPL/b in XSPEC).","between these two spectra was then fit with a cut-off power-law model in PHA space (i.e., CUTOFFPL/b in XSPEC)."
+ The power-law index and cut-off energy in the CUTOFFPL model were fixed to the values found by Markevitch et OOnly the normalization was allowed to vary., The power-law index and cut-off energy in the CUTOFFPL model were fixed to the values found by Markevitch et Only the normalization was allowed to vary.
+ All spectral analysis of ACIS-S3 data presented below includes a CUTOFFPL/b model with the normalization determined from the area of the source extraction region., All spectral analysis of ACIS-S3 data presented below includes a CUTOFFPL/b model with the normalization determined from the area of the source extraction region.
+ Separate photon-weighted response and effective area files are generated for each spectrum., Separate photon-weighted response and effective area files are generated for each spectrum.
+ The effective area files are corrected for the effects of contamination on ACIS using the corrarf task., The effective area files are corrected for the effects of contamination on ACIS using the corrarf task.
+" All ACIS images presented in this paper only include photons with energies between 0.3-6.0 keV. By excluding higher energy photons, we eliminate approximately of the S3 background at the expense of only of the emission from a 9 keV thermal source."," All ACIS images presented in this paper only include photons with energies between 0.3–6.0 keV. By excluding higher energy photons, we eliminate approximately of the S3 background at the expense of only of the emission from a 9 keV thermal source."
+" In addition to the nominal background rate, the soft flare is also removed from all imaging analysis using the best fit normalization of the CUTOFFPL model derived above."," In addition to the nominal background rate, the soft flare is also removed from all imaging analysis using the best fit normalization of the CUTOFFPL model derived above."
+" Abell 1758 was observed by XMM-Newton on Nov. 11-12, 2002."," Abell 1758 was observed by XMM-Newton on Nov. 11–12, 2002."
+ All XMM-Newton analysis was done with SAS v5.4.1., All XMM-Newton analysis was done with SAS v5.4.1.
+ The data from all three EPIC cameras were first reprocessed the tasks emchain and epchain., The data from all three EPIC cameras were first reprocessed using the tasks emchain and epchain.
+" The data were then usingscreened using the standard patterns and flags specified in “An Introduction to XMM-Newton Data Analysis"".", The data were then screened using the standard patterns and flags specified in “An Introduction to XMM-Newton Data Analysis”.
+" To screen out periods of enhanced background, we generated light curves for each detector in the 10—15 keV band."," To screen out periods of enhanced background, we generated light curves for each detector in the 10–15 keV band."
+" We then performed a successive 2σ clipping of the data which yielded 49435s, 49185s, and 43471s of screened data for the MOSI, MOS2, and PN cameras, "," We then performed a successive $2 \sigma$ clipping of the data which yielded 49435s, 49185s, and 43471s of screened data for the MOS1, MOS2, and PN cameras, respectively."
+We only included events with FLAG=0 in our spectral respectively.analysis., We only included events with FLAG=0 in our spectral analysis.
+ The large field of view of the EPIC cameras allowed us to extract background spectra from the same data sets., The large field of view of the EPIC cameras allowed us to extract background spectra from the same data sets.
+ The background spectrum from all 3 cameras was extracted from within a 2’ radius circle located 11’ due north of A1758N. Separate response and effective area files were generated for each EPIC spectra using the SAS tasks rmfgen and arfgen., The background spectrum from all 3 cameras was extracted from within a $2^{\prime}$ radius circle located $11^{\prime}$ due north of A1758N. Separate response and effective area files were generated for each EPIC spectra using the SAS tasks rmfgen and arfgen.
+ A detector map was used to weight each effective area file., A detector map was used to weight each effective area file.
+" An adaptively smoothed, exposure corrected mosaic of the 3 XMM-Newton EPIC images is shown in Fig. 1.."," An adaptively smoothed, exposure corrected mosaic of the 3 XMM-Newton EPIC images is shown in Fig. \ref{fig:xmm_mosaic}."
+ A1758N and A1758S are clearly visible in this figure with a projected separation of approximately 2.0 Mpc at the redshifts of the clusters., A1758N and A1758S are clearly visible in this figure with a projected separation of approximately 2.0 Mpc at the redshifts of the clusters.
+" For comparison, the virial radii of the two clusters estimated below are 2.6 Mpc and 2.2 for A1758N and A1758S, respectively."," For comparison, the virial radii of the two clusters estimated below are 2.6 Mpc and 2.2 Mpc for A1758N and A1758S, respectively."
+ Fig., Fig.
+ | clearly showsMpc the presence of X-ray emission between the two clusters., \ref{fig:xmm_mosaic} clearly shows the presence of X-ray emission between the two clusters.
+ This emission is analyzed in detail below to search for any signs of interaction between the two systems., This emission is analyzed in detail below to search for any signs of interaction between the two systems.
+ Both A1758N and A1758S contain significant substructure., Both A1758N and A1758S contain significant substructure.
+ A1758N contains two subclusters separated along a line running from SE to NW., A1758N contains two subclusters separated along a line running from SE to NW.
+ A1758S has a large subcluster off-set towards the NE from the main component of the cluster., A1758S has a large subcluster off-set towards the NE from the main component of the cluster.
+ The source to the east of A1758S has a spatial extent consistent with that of a point source., The source to the east of A1758S has a spatial extent consistent with that of a point source.
+ Contours of the smoothed XMM-Newton mosaic are overlayed on the DSS adaptivelyimage in Fig. 2.., Contours of the adaptively smoothed XMM-Newton mosaic are overlayed on the DSS image in Fig. \ref{fig:dss_xmm_mosaic}.
+" This figure shows that the large scale contours of A1758S are roughly centered on the brightest galaxy in the cluster, while the subcluster toward the NE in A1758S is centered on a group of fainter galaxies."," This figure shows that the large scale contours of A1758S are roughly centered on the brightest galaxy in the cluster, while the subcluster toward the NE in A1758S is centered on a group of fainter galaxies."
+" The subcluster in A1758N off-set toward the NW is located near the center of a line of bright galaxies, while the SE subcluster in A1758N is off-set toward the NW from a fainter group of galaxies."," The subcluster in A1758N off-set toward the NW is located near the center of a line of bright galaxies, while the SE subcluster in A1758N is off-set toward the NW from a fainter group of galaxies."
+ The X-ray morphologies of A1758N and A1758S are discussed in more detail below., The X-ray morphologies of A1758N and A1758S are discussed in more detail below.
+" To search for signs of interaction between A1758N and A1758S, we extracted the counts from the PN image within the rectangular region shown in Fig. 2.."," To search for signs of interaction between A1758N and A1758S, we extracted the counts from the PN image within the rectangular region shown in Fig. \ref{fig:dss_xmm_mosaic}."
+ Only the PN image is used in this analysis due to its lower background compared to the MOS detectors., Only the PN image is used in this analysis due to its lower background compared to the MOS detectors.
+ A projected surface brightness profile, A projected surface brightness profile
+In most cases the distances to cataclysmic variables are highly uncertain. but there is now a growing number of systems for which parallaxes have been measured.,"In most cases the distances to cataclysmic variables are highly uncertain, but there is now a growing number of systems for which parallaxes have been measured."
+ Ten systems within our sample have measured: parallaxcs (see reftable:tum)) ancl we place particular emphasis on these systems., Ten systems within our sample have measured parallaxes (see \\ref{table:lum}) ) and we place particular emphasis on these systems.
+ Distances for other svstenis in our sample were taken from many sources. listed in 1..," Distances for other systems in our sample were taken from many sources, listed in \ref{table:general}."
+" X-ray luminosities for the systems with measured parallaxes are presented in reftablebum.. ane our full set of estimated luminosities are presented in relfig:luny, (sf."," X-ray luminosities for the systems with measured parallaxes are presented in \\ref{table:lum}, and our full set of estimated luminosities are presented in \\ref{fig:lum_hist}."
+ hegarcealsoplolledasaf unctlionoforbitalperiodinl RU, They are also plotted as a function of orbital period in \\ref{fig:lum} \\ref{fig:inc}.
+ IN (2) ?)), \\ref{fig:lum_hist} \\ref{fig:lum} \cite{wzsgeprimarymass} \ncite{whitedwarfradius})
+ IN (2) ?))., \\ref{fig:lum_hist} \\ref{fig:lum} \cite{wzsgeprimarymass} \ncite{whitedwarfradius})
+Since the discovery of its enormous IR luminosity (Lig jmi)e-1.6 x10PL.. Soifer et al.,"Since the discovery of its enormous IR luminosity $\rm
+ L_{IR}$ $\mu $ $\sim $$1.6\times 10^{12}$$\rm
+ \,L_{\odot}$, Soifer et al."
+ 1934: Emerson οἱ al., 1984; Emerson et al.
+ 1984). representing ~99% of its bolometric luminosity. 2220 has been a study of extremes when it comes to the ISM conditions in intense starbursts in the local Universe.," 1984), representing $\sim 99\%$ of its bolometric luminosity, 220 has been a study of extremes when it comes to the ISM conditions in intense starbursts in the local Universe."
+ A prominent member in a prominent class of galaxies (e.g. Sanders Ishida 2004). the large IR. luminosity and relative proximity (al Dp~T7 MMpe. it is the nearest ULIBRG). made 2220 an early target of numerous molecular line observations.," A prominent member in a prominent class of galaxies (e.g. Sanders Ishida 2004), the large IR luminosity and relative proximity (at $\rm D_L$$\sim
+ $ Mpc, it is the nearest ULIRG), made 220 an early target of numerous molecular line observations."
+ Its large CO J=10 line luminosity revealed a huge molecular gas reservoir (~LOMAL.) (Young et al.," Its large CO J=1–0 line luminosity revealed a huge molecular gas reservoir $\rm \sim
+ 10^{10}\,M_{\odot}$ ) (Young et al."
+ 1984: Sanders. Scoville. soifer 1991: Solomon οἱ al.," 1984; Sanders, Scoville, Soifer 1991; Solomon et al."
+ 1997). while luminous transitions of heavy rotor molecules such as CS J=32. and HCN. — J—10 have demonstrated that. quite unlike typical quiescent spirals. most of the molecular gas in this spectacular merger is very dense. wilh nzl0em (Solomon et al.," 1997), while luminous transitions of heavy rotor molecules such as CS J=3–2, and HCN, $^{+}$ J=1–0 have demonstrated that, quite unlike typical quiescent spirals, most of the molecular gas in this spectacular merger is very dense, with $\rm n\ga 10^5\,cm^{-3}$ (Solomon et al."
+ 1990. 1992).," 1990, 1992)."
+ Such studies have recently culminated in the most extensive molecular line survey. ever conducted for such objects (Greve οἱ al., Such studies have recently culminated in the most extensive molecular line survey ever conducted for such objects (Greve et al.
+ 2009). making 2220 a galaxy. with the best studied molecular gas reservoirs in the local Universe.," 2009), making 220 a galaxy with the best studied molecular gas reservoirs in the local Universe."
+" Ligh resolution CO J—10. 21 interferometric imaging reveals (wo compact (2r50.3"": ppc) eas concentrations c1"" ppc) apart (Scoville. Yun. Brvant 1997; Downes Solomon 1998: Sakamoto οἱ al."," High resolution CO J=1–0, 2–1 interferometric imaging reveals two compact $\la $ $''$; pc) gas concentrations $\sim $ $''$ pc) apart (Scoville, Yun, Bryant 1997; Downes Solomon 1998; Sakamoto et al."
+ 1999: Eckart. Downes 2001). whose intense starbursis dominate the IR. luminosity of the entire svstem.," 1999; Eckart, Downes 2001), whose intense starbursts dominate the IR luminosity of the entire system."
+ Finally a major advance was recently made with the imaging of ihe CO J=32 emission (and the adjacent continua at 860jn) of the warm and dense gas. that also implied substantial dust optical depths even al subi wavelengths (Sakamoto et al.," Finally a major advance was recently made with the imaging of the CO J=3–2 emission (and the adjacent continua at $\mu $ m) of the warm and dense gas, that also implied substantial dust optical depths even at submm wavelengths (Sakamoto et al."
+ 2009)., 2009).
+ We report single dish CO J=65 observations of 2220 with the James Clerk Maxwell Telescope as part ofa niilti-J CO and HCN line survey of LIRGs (Liz 10! LL. ) in the local Universe., We report single dish CO J=6–5 observations of 220 with the James Clerk Maxwell Telescope as part of a multi-J CO and HCN line survey of LIRGs $\rm L_{IR}$$\ga $ $^{11}$ $_{\odot}$ ) in the local Universe.
+ We demonstrate that an optically thick dust continuum at submim wavelengths submerges this CO line to near-blackbody emission. and strongly affects the emergent CO Spectral Line Energy Distribution (SLED) and the fine-structure line liminosity for this extreme starburst.," We demonstrate that an optically thick dust continuum at submm wavelengths submerges this CO line to near-blackbody emission, and strongly affects the emergent CO Spectral Line Energy Distribution (SLED) and the $^+$ fine-structure line luminosity for this extreme starburst."
+ Finally we discuss whether similar ISM conditions are present in dustiy starbursts in the distant Universe. lowering their observed (high-J)/(low-J) CO ratios well below those (vpical of slar-lorming molecular gas. ancl (hus limiting their diagnostic utility.," Finally we discuss whether similar ISM conditions are present in dusty starbursts in the distant Universe, lowering their observed (high-J)/(low-J) CO ratios well below those typical of star-forming molecular gas, and thus limiting their diagnostic utility."
+goes to unity and the keplerian term. Lf. dominates.,goes to unity and the keplerian term $1/x$ dominates.
+ Actually. this is an obvious consequence of the finite total mass of the model.," Actually, this is an obvious consequence of the finite total mass of the model."
+ Note. however. that. the keplerian behaviour is attained only for large values of «so that 0. is quite slowly declining and. within the observational errors. may be also compatible with a Lat rotation curve.," Note, however, that the keplerian behaviour is attained only for large values of $x$ so that $v_c$ is quite slowly declining and, within the observational errors, may be also compatible with a flat rotation curve."
+ Let us now derive the gravitational potential for PoLLS models., Let us now derive the gravitational potential for PoLLS models.
+ To this aim. the Poisson:: has to be solved.," To this aim, the Poisson: has to be solved."
+ For spherically symmetric system. the &eneral solution is (Binney&‘Tremaine1987) Using I2q.(11)) for the mass profile and I50.(10)) for the ensity law. weect: Being the total mass finite. Eq.(18)) easily shows that the eravitational potential vanishes at infinity.," For spherically symmetric system, the general solution is \cite{BT87}: Using \ref{eq: mass}) ) for the mass profile and \ref{eq: rho}) ) for the density law, we: : Being the total mass finite, \ref{eq: genphi}) ) easily shows that the gravitational potential vanishes at infinity."
+ However. a caveat is in order here.," However, a caveat is in order here."
+ H£ we consider the limit for 5ox of IEq.(19)). we find (35)=OMfr2V(2/5)/U(S/).," If we consider the limit for $x \rightarrow \infty$ of \ref{eq: phi}) ), we find $\Phi(\infty) = G M_{tot}/r_{-2} \times \Gamma(2/\gamma)/\Gamma(3/\gamma)$."
+ Since the potential is defined by (18)) up to an arbitrary additive constant. we could. add this value to Iq.(19)) in order to have Gr) vanishing at infinity as it is indeed the case.," Since the potential is defined by \ref{eq: genphi}) ) up to an arbitrary additive constant, we could add this value to \ref{eq: phi}) ) in order to have $\Phi(x)$ vanishing at infinity as it is indeed the case."
+ At the opposite end. having the model a finite force at the centre. the gravitational potential does not diverge in the origin.," At the opposite end, having the model a finite force at the centre, the gravitational potential does not diverge in the origin."
+ By applving the limit for a-0 in I2q.(19)). weτ: which is plotted in reffig: phiz..," By applying the limit for $x \rightarrow 0$ in \ref{eq: phi}) ), we: which is plotted in \\ref{fig: phiz}."
+ The central potential turns out to be an function of the slope parameter 5., The central potential turns out to be an function of the slope parameter $\gamma$.
+ As a consequence. the depth of the potential well is greater for larger values of 5 and. for these mocels. particles are more attracted toward the centre.," As a consequence, the depth of the potential well is greater for larger values of $\gamma$ and, for these models, particles are more attracted toward the centre."
+ Pherefore. for a given ur. thegravitational force and hence the number of stars within a distance wv from the centre is larger for larger values of 7 thus explaining why AMGr)/Mio is an increasing function of the slope parameter.," Therefore, for a given $x$, thegravitational force and hence the number of stars within a distance $x$ from the centre is larger for larger values of $\gamma$ thus explaining why $M(x)/M_{tot}$ is an increasing function of the slope parameter."
+ Another dynamically interesting quantity that may be straightforwardly evaluated. is the velocity dispersion., Another dynamically interesting quantity that may be straightforwardly evaluated is the velocity dispersion.
+ Assuming isotropy in the velocity space. this is given. by (Binney&Tremaine1987) Inserting Eqs.(10)) and (11)) into the above general expression. we: with This quantity may not be evaluated analytically. but it is straightforward to estimate for fixed. values of the slope parameter 5.," Assuming isotropy in the velocity space, this is given by \cite{BT87}: Inserting \ref{eq: rho}) ) and \ref{eq: mass}) ) into the above general expression, we: with This quantity may not be evaluated analytically, but it is straightforward to estimate for fixed values of the slope parameter $\gamma$ ."
+ relfig:siglig shows the velocity dispersion for three values of  using Mj;—5101M. andr»=10kpe., \\ref{fig: sigfig} shows the velocity dispersion for three values of $\gamma$ using $M_{tot} = 5 \times 10^{11} \ {\rm M_{\odot}}$ and $r_{-2} = 10 \ {\rm kpc}$.
+ The result maw be simply scaled to other values of (Alia.2) bv noting hat. because of I5q.(23)). axVEITroo.," The result may be simply scaled to other values of $(M_{tot}, r_{-2})$ by noting that, because of \ref{eq: sigma}) ), $\sigma \propto \sqrt{M_{tot}/r_{-2}}$."
+ As vet observed or the rotational velocity. for a given aw. a is higher for ueher values of 5.," As yet observed for the rotational velocity, for a given $x$, $\sigma$ is higher for higher values of $\gamma$."
+" In both cases. we have thus a degeneracy οοσα the total mass and the slope parameter since both hese quantities increase v, and e."," In both cases, we have thus a degeneracy between the total mass and the slope parameter since both these quantities increase $v_c$ and $\sigma$."
+ However. the degeneracy may be broken by determining where the rotation curve »aks since this only depends on 5.," However, the degeneracy may be broken by determining where the rotation curve peaks since this only depends on $\gamma$ ."
+ As a further remark. we note that o decreases in the outer regions of the system. husmimickine well what has been recently observed.in intermediate Luminosity elliptical galaxies 2003)..," As a further remark, we note that $\sigma$ decreases in the outer regions of the system thusmimicking well what has been recently observedin intermediate luminosity elliptical galaxies \cite{Nicola}. ."
+optical/UV to N-rav are necessary to constraim iutriusic huninosities to within a factor of a few (as we show in 3.1).,optical/UV to X-ray are necessary to constrain intrinsic luminosities to within a factor of a few (as we show in 3.1).
+ We select a sample of 945 ACNs from the Cosmic Evolution Survey field. which is based ou the 1.7 dee? HST/ACS mosaic2007).," We select a sample of 348 AGNs from the Cosmic Evolution Survey field, which is based on the 1.7 $^2$ HST/ACS mosaic."
+. These ACNs have nultiwaveleneth data in the form of Spitzer/IRAC. HST/ACS. Subaru/Suprine-Ciuu. CALEN. NAILNewton. and Chnandra observations. as described iu Table 1..," These AGNs have multiwavelength data in the form of Spitzer/IRAC, HST/ACS, Subaru/Suprime-Cam, GALEX, XMM-Newton, and Chandra observations, as described in Table \ref{tbl:obsdata}."
+" Spectroscopic identification and redshifts for these obOR comes from, archival SDSS data. MagellanTAIACS and MMT/IHectospec. 2009a).. aud VLT/VIMOS observations)."," Spectroscopic identification and redshifts for these objects comes from archival SDSS data, Magellan/IMACS and MMT/Hectospec , and VLT/VIMOS observations."
+. The sample is selected from the parent catalog of 1651 NMM-COSAMDOS point sources with opticalcounterparts 2OL0).. limited bv fosony72x101 eve s| cm7.," The sample is selected from the parent catalog of 1651 XMM-COSMOS point sources with opticalcounterparts , limited by $f_{0.5-2 {\rm keV}} > 2 \times 10^{-16}$ erg $^{-1}$ $^{-2}$."
+" Of these Nouv ""-— ποσόν, 619 objects likelihoodwith /4n have hieh-coutidence (>90% as correct) ideutifications and redshifts frou, optical spectroscopy in COSMOS."," Of these X-ray point sources, 649 objects with $i_{\rm AB}<23.5$ have high-confidence $>90\%$ likelihood as correct) identifications and redshifts from optical spectroscopy in COSMOS."
+ Most. of the N-vay point sources without spectroscopy were missed sinplv due to random slit placement constraints., Most of the X-ray point sources without spectroscopy were missed simply due to random slit placement constraints.
+ The optical spectroscopy is ~90% complete to /4p«22.5. although the completeness is redshift-depeudeut.," The optical spectroscopy is $\sim$ complete to $i_{\rm AB}<22.5$, although the completeness is redshift-dependent."
+ For broad-line AGNs. the spectroscopic completeness is lower at 0.5<κ ντ~lod. aud 2~2.1. especially at (yp>22.52009a)..," For broad-line AGNs, the spectroscopic completeness is lower at $0.5<z<1$ , $z \sim 1.4$, and $z
+\sim 2.4$, especially at $i_{\rm AB}>22.5$."
+" For uarrow- and lincless Αννα, spectroscopic coiipletenuess drops dramatically at 2>1.2. since at higher redshifts the break and the feature shift recdhward of he observed wavelength[O range."," For narrow-line and lineless AGNs, spectroscopic completeness drops dramatically at $z>1.2$, since at higher redshifts the break and the feature shift redward of the observed wavelength range."
+"ri] To ensure that X-ray objects with narrow-line aud liueless spectra are fide ACUNsS. we select only objects hunitesitywith Losqug773<10% eres ο,"," To ensure that X-ray objects with narrow-line and lineless spectra are bona-fide AGNs, we select only objects with $L_{0.5-10 {\rm keV}} > 3 \times 10^{42}$ erg $^{-1}$."
+ This N-ray Iitis generally used o separate ACUNS frou X-ray fainter starburst galaxies2001)., This X-ray luminosity limit is generally used to separate AGNs from X-ray fainter starburst galaxies.
+. We also iuclude seven woad-line ACNs without N-ray detection. six of which were selected by their Spitzer/IRAC colors and one which js a serendipitous object from the bright zCOSAIOS survey (which selected targets based only ou £ap< 22.5).," We also include seven broad-line AGNs without X-ray detection, six of which were selected by their Spitzer/IRAC colors and one which is a serendipitous object from the bright zCOSMOS survey (which selected targets based only on $i_{\rm AB}<22.5$ )."
+" While these 7 N-ray uudetected ACNs do not come from a complete sample. we include them to gain a larecr piriuueter space of ACN spectral types and accretion rates (in effect. when using them N-rayv limits. thev occupy the same Lys,ον. parameter space as a few other X-ray detected ACNs)."," While these 7 X-ray undetected AGNs do not come from a complete sample, we include them to gain a larger parameter space of AGN spectral types and accretion rates (in effect, when using their X-ray limits, they occupy the same $L_{disk}/L_X$ parameter space as a few other X-ray detected AGNs)."
+ Restricting narrow-line aud lineless ACNs be X-ray luminous and adding the 7 X-rav undetected tobroad-line AGNs makes a parent sample of 380 broad-line. 121 uarrow-line. and 19 lineless ACUNS (553 total) with bigh-confidence redshifts aud spectral identification.," Restricting narrow-line and lineless AGNs to be X-ray luminous and adding the 7 X-ray undetected broad-line AGNs makes a parent sample of 380 broad-line, 124 narrow-line, and 49 lineless AGNs (553 total) with high-confidence redshifts and spectral identification."
+ Measuring accurate black hole masses additionally constrains the sample to certain redshift ranges., Measuring accurate black hole masses additionally constrains the sample to certain redshift ranges.
+ For Type D AGNs. we require the presence of one of theCiv...Γον or broad emission lines in the observed spectral ranee. effectively limiting broad-line AGNs with IMACS or VIMOS spectra to 0.16COs. ἐκτς 2.[. and 2.7i<L9 and objects with IIectospec or SDSS spectra to z<L9.," For Type 1 AGNs, we require the presence of one of the, or broad emission lines in the observed spectral range, effectively limiting broad-line AGNs with IMACS or VIMOS spectra to $0.16<z<0.88$, $1<z<2.4$ , and $2.7<z<4.9$ and objects with Hectospec or SDSS spectra to $z<4.9$."
+ For uarrow-line aud lineless AGNs. we estimate black hole mass from the Mpg~ relation. and so we require au onate of Linty:Lügg.," For narrow-line and lineless AGNs, we estimate black hole mass from the $M_{BH} \sim L_{bulge}$ relation, and so we require an accurate estimate of $L_{bulge}$."
+ For this we use the suuple EMofobjects in COSMOS with morphological decompositions from the IST/ACS images2007).. which also effectively limits the niurow- and lineless ACGNs to :<1.2 (bevoud which he break shifts out of the ACS-/ band aud he host galaxv is much more difficult to detect).," For this we use the sample of objects in COSMOS with morphological decompositions from the HST/ACS images, which also effectively limits the narrow-line and lineless AGNs to $z<1.2$ (beyond which the break shifts out of the $i$ band and the host galaxy is much more difficult to detect)."
+" The accurate host measurements from additionally allow us to subtract the host component vefore computing the intrinsic bolometric ποσατν,", The accurate host measurements from additionally allow us to subtract the host component before computing the intrinsic bolometric luminosity.
+ Iu eeneral. the narrow-line and lineless ACINs are biased owards lower redshift aud consequently higher mass. since ACNs erow over the cosmic time.," In general, the narrow-line and lineless AGNs are biased towards lower redshift and consequently higher mass, since AGNs grow over the cosmic time."
+" The narrow-line and lineless ACNs have a mean redshift of 0.7. while the xoad-liue ACINS have a mean ""n of 1.6."," The narrow-line and lineless AGNs have a mean redshift of 0.7, while the broad-line AGNs have a mean redshift of 1.6."
+nThe final siuuple of 318 ACNs includes 256 broad-lne. ine. andnulticveleneth 27 lincless ACNs.," The final sample of 348 AGNs includes 256 broad-line, 65 narrow-line, and 27 lineless AGNs."
+ Full data exist for >95% of the AGNs in the sample in every wavelength region except the UV., Full multiwavelength data exist for $>95\%$ of the AGNs in the sample in every wavelength region except the UV.
+ A-rav data exist from both Chandra and NAIAI- we use the deeper Chandra data when available. but the Chandra observations cover oulv the ceutral 0.5 dee? of the COSMOS field.," X-ray data exist from both Chandra and XMM-Newton: we use the deeper Chandra data when available, but the Chandra observations cover only the central 0.8 $^2$ of the COSMOS field."
+ For the 7 N-rav uudetected broad-line ACINs. we use the 0.5-2 keV NMM flux limit (fosom=2<1015 cre tem 7) for their Χαν luminosity (since these ACNs have ων>10. their bolometric huninosity is dominated by their optical/UV enission and completely neglecting their X-ray eiission does not significantly chanee their bolometric Iuninositv estimate).," For the 7 X-ray undetected broad-line AGNs, we use the 0.5-2 keV XMM flux limit $f_{0.5-2 {\rm keV}} = 2 \times 10^{-16}$ erg $^{-1}$ $^{-2}$ ) for their X-ray luminosity (since these AGNs have $L_{disk}/L_X>10$, their bolometric luminosity is dominated by their optical/UV emission and completely neglecting their X-ray emission does not significantly change their bolometric luminosity estimate)."
+ We apply the zero-point offsets derived by to the IR-UV photometry., We apply the zero-point offsets derived by to the IR-UV photometry.
+ A-rav absorption and optical/UV extinction could pose a challenge to measuring the iutrindüc accretion power., X-ray absorption and optical/UV extinction could pose a challenge to measuring the intrinsic accretion power.
+ The most heavily absorbed AGNs (oe. Comptou-thick ACNs with Nyo>1075 7) me cutirely missed. by our survey because they lack detectable X- cluission2001)., The most heavily absorbed AGNs (e.g. Compton-thick AGNs with $N_H>10^{24}$ $^{-2}$ ) are entirely missed by our survey because they lack detectable X-ray emission.
+ But if au AGN is moderately absorbed aud still N-rav detected. we inight expect its disk to appear cooler because the UV lieht is prefercutially extiucted. and its N-rav slope to appear harder because the soft N-ravs are preferentially absorbed.," But if an AGN is moderately absorbed and still X-ray detected, we might expect its disk to appear cooler because the UV light is preferentially extincted, and its X-ray slope to appear harder because the soft X-rays are preferentially absorbed."
+ Some ACNs are also intrinsically reddened. decreasing their UV cussion by a factor of 2-3 aud. causing us to underestimate their accretion disk cussion.," Some AGNs are also intrinsically reddened, decreasing their UV emission by a factor of 2-3 and causing us to underestimate their accretion disk emission."
+" With absorbed soft X-ravs and extincted disk emissiou. we could sieuificantly uuderestimate L;,;/Lp."," With absorbed soft X-rays and extincted disk emission, we could significantly underestimate $L_{int}/L_{Edd}$."
+" We use N-arav column density Vy, to characterize the obscuration properties of our ACNs.", We use X-ray column density $N_H$ to characterize the obscuration properties of our AGNs.
+ Colum density aud optical extinction are roughly correlated. with Ay-/Ny~24.1078 cue2006)..," Column density and optical extinction are roughly correlated, with $A_V/N_H \sim 2 \times 10^{-23}$ $^2$."
+ Then at Ny<107 απ2. optical magnitude should be extincted ye 204 (=0.2 mag).," Then at $N_H<10^{22}$ $^{-2}$, optical magnitude should be extincted by $\lesssim 20\%$ $\lesssim 0.2$ mag)."
+ Assuming a SMC reddening aw1992).. as is most appropriate forACUNs. this otical extinction translates to a factor of ~1.2 extinction 3000Ain the UV.," Assuming a SMC reddening law, as is most appropriate forAGNs, this optical extinction translates to a factor of $\sim$ 1.2 extinction at in the UV."
+ showed hat the Ay:Nyy relation varies by up to a factor of 30 because of uuknown changes in the sgas-to-dust ratio. erain size. and/or different plysical locations of he optical and N-vav absorbing material.," showed that the $A_V-N_H$ relation varies by up to a factor of 30 because of unknown changes in the gas-to-dust ratio, grain size, and/or different physical locations of the optical and X-ray absorbing material."
+ However for all AGNs in the sample with Ly>lo’ ores LoyiNds<077 en. meaning at Ny~10°? 7 even the masimuun optical (V-band)," However for all AGNs in the sample with $L_X>10^{42}$ erg $^{-1}$ , $A_V/N_H <
+1.8 \times 10^{-22}$ $^2$ , meaning at $N_H \sim 10^{22}$ $^{-2}$ even the maximum optical $V$ -band)"
+stratification of the physical parameters as a function of the logarithm of the LOS continuum optical depth at 5000 ((log τ) can be obtained.,stratification of the physical parameters as a function of the logarithm of the LOS continuum optical depth at $5000$ $\log\tau$ ) can be obtained.
+ We tested the effects on the inversions of using different initial model atmospheres (umbra. penumbra and plage models). to asses their performance when reproducing the observed Stokes profiles.," We tested the effects on the inversions of using different initial model atmospheres (umbra, penumbra and plage models), to asses their performance when reproducing the observed Stokes profiles."
+ The penumbral model from ? seemed to yield the best results. which is consistent with what we see in the observations at the photosphere.," The penumbral model from \citet{toro94} seemed to yield the best results, which is consistent with what we see in the observations at the photosphere."
+ However. some modifications to the model (such as assuming an initial constant magnetic field strength of GG and a LOS velocity of 0.1 kmss7!) had to be implemented.," However, some modifications to the model (such as assuming an initial constant magnetic field strength of G and a LOS velocity of 0.1 $^{-1}$ ) had to be implemented."
+ Different initial. values for the inclination. and azimuth did not affect the final fit to the observed Stokes profiles., Different initial values for the inclination and azimuth did not affect the final fit to the observed Stokes profiles.
+ Macroturbulence was fixed to the same value as for the ME aversions., Macroturbulence was fixed to the same value as for the ME inversions.
+ However. for the inversion of the line. the stray-light was initialized with 30% and left as a free parameter in the fit.," However, for the inversion of the line, the stray-light was initialized with $30\%$ and left as a free parameter in the fit."
+ Stray-light profiles for each map were computed by averaging the Stokes / of non-magnetic areas. 1.8. regions where Stokes Q. U and V are negligible.," Stray-light profiles for each map were computed by averaging the Stokes $I$ of non-magnetic areas, i.e., regions where Stokes $Q$, $U$ and $V$ are negligible."
+ The necessary atomic data for the line was taken from the work of ?.., The necessary atomic data for the line was taken from the work of \citet{borrero03}.
+ In particular. the value of the logarithm of the oscillator strength times the multiplicity of the lower level that we used was log(gf)=0.363.," In particular, the value of the logarithm of the oscillator strength times the multiplicity of the lower level that we used was $\log (g f) = 0.363$."
+ To have a rough idea of the formation height of Silicon in order to associate an appropriate optical depth for the inferred vector magnetic field. response functions to magnetic field perturbations at various positions near or in the filament were calculated.," To have a rough idea of the formation height of Silicon in order to associate an appropriate optical depth for the inferred vector magnetic field, response functions to magnetic field perturbations at various positions near or in the filament were calculated."
+ Our atmospheric model covers heights from 1.2 up to -4.0 (in logr units)., Our atmospheric model covers heights from 1.2 up to -4.0 (in $\log\tau$ units).
+ The largest sensitivity for both days was found to take place at logr=—2., The largest sensitivity for both days was found to take place at $\log\tau = -2$.
+ Thus. from this point onwards. all of the figures derived from the inversions of the line are referred to this height.," Thus, from this point onwards, all of the figures derived from the inversions of the line are referred to this height."
+ Figures 7 — 9 show the results of the MELANIE and SIR inversions of the Stokes parameters for three different positions along the filament: one in a Helium dark thread. one in the spine and one at the PIL.," Figures \ref{Fig:profhethread1} – \ref{Fig:profNL3} show the results of the MELANIE and SIR inversions of the Stokes parameters for three different positions along the filament: one in a Helium dark thread, one in the spine and one at the PIL."
+ Each figure is made of 8 plots: in the row we present the AA)) observed Stokes 7. Q. U and V profiles (dots) obtained after performing the binning and the best fit achieved by the inversion code (solid line).," Each figure is made of 8 plots: in the row we present the ) observed Stokes $I$, $Q$, $U$ and $V$ profiles (dots) obtained after performing the binning and the best fit achieved by the inversion code (solid line)."
+ The exact location of the selected fits is indicated by short white arrows in Figs. 5..," The exact location of the selected fits is indicated by short white arrows in Figs. \ref{Fig:TIPmaps},"
+ 11 and 12 (see the captions of the figures for a detailed explanation)., \ref{Fig:magnetoHe} and \ref{Fig:magnetoSi} (see the captions of the figures for a detailed explanation).
+ MELANIE does not provide the uncertainties in the retrieved atmospheric parameters., MELANIE does not provide the uncertainties in the retrieved atmospheric parameters.
+ To estimate them. we used the synthetic Stokes profiles which resulted from the best fit of the model to the data.," To estimate them, we used the synthetic Stokes profiles which resulted from the best fit of the model to the data."
+ Then. several different realizations of the noise (with an amplitude of that of the noise in the observations). were added to the synthetic Stokes profiles. which were in turn inverted again.," Then, several different realizations of the noise (with an amplitude of that of the noise in the observations), were added to the synthetic Stokes profiles, which were in turn inverted again."
+ The standard deviation computed from the spread in the values of the retrieved parameters provided the errors quoted in the captions of Figs. 7..," The standard deviation computed from the spread in the values of the retrieved parameters provided the errors quoted in the captions of Figs. \ref{Fig:profhethread1},"
+ 8 and 9.., \ref{Fig:profspine2} and \ref{Fig:profNL3}.
+ The SIR inversions directly provide uncertainties which are proportional to the inverse of the response functions to changes in the physical parameters., The SIR inversions directly provide uncertainties which are proportional to the inverse of the response functions to changes in the physical parameters.
+ We found a very satisfactory performance of the ME inversion code when fitting the Stokes profiles., We found a very satisfactory performance of the ME inversion code when fitting the Stokes profiles.
+ This has already been described by ? who found in these data a ubiquitous presence of Zeeman-like signatures(see Stokes Q and U frames in Figs., This has already been described by \citet{kuckein09} who found in these data a ubiquitous presence of Zeeman-like signatures(see Stokes $Q$ and $U$ frames in Figs.
+ 7. — 9)) with no apparent contribution of atomic level polarization and its modification by the Hanle effect., \ref{Fig:profhethread1} – \ref{Fig:profNL3}) ) with no apparent contribution of atomic level polarization and its modification by the Hanle effect.
+ The presence of strong Stokes Q and U profiles is indicative of strong transverse magnetic fields in the filament., The presence of strong Stokes $Q$ and $U$ profiles is indicative of strong transverse magnetic fields in the filament.
+ In the photospheric aabsorption line we also find strong Stokes Q and U profiles. which occasionally have larger amplitudes than the corresponding Stokes V.," In the photospheric absorption line we also find strong Stokes $Q$ and $U$ profiles, which occasionally have larger amplitudes than the corresponding Stokes $V$."
+ Consequently. a predominant horizontal field is found at photospheric heights too.," Consequently, a predominant horizontal field is found at photospheric heights too."
+ Prior to the inversions of the line. we studied the possible impact of non-local thermodynamic equilibrium (NLTE) effects in the retrieved atmospheric. parameters.," Prior to the inversions of the line, we studied the possible impact of non–local thermodynamic equilibrium (NLTE) effects in the retrieved atmospheric parameters."
+ SIR. synthesizes the Stokes profiles under the assumption of local thermodynamic equilibrium., SIR synthesizes the Stokes profiles under the assumption of local thermodynamic equilibrium.
+ A recent study by ? reported that this line can be significantly affected by NLTE conditions., A recent study by \citet{bard08} reported that this line can be significantly affected by NLTE conditions.
+ The authors showed that the NLTE line core intensity is deeper than the analogous LTE result for two different model atmospheres: quiet sun (FALC model) and suspot umbra models (SPOTM model)., The authors showed that the NLTE line core intensity is deeper than the analogous LTE result for two different model atmospheres: quiet sun (FALC model) and sunspot umbra models (SPOTM model).
+ To study possible effects on our inverstor results. we introduced the departure coefficients 8 (which are defined as the ratio between population densities in NLTE over LTE) from the above cited paper (kindly provided to us by M. Carlsson). into the SIR code.," To study possible effects on our inversion results, we introduced the departure coefficients $\beta$ (which are defined as the ratio between population densities in NLTE over LTE) from the above cited paper (kindly provided to us by M. Carlsson), into the SIR code."
+ Inversions of a few selected cases where run with the B coefficients from these models (together with the LTE case) and the differences between the inferred magnetic field strength (B). inclination (y) and azimuth (@) were studied.," Inversions of a few selected cases where run with the $\beta$ coefficients from these models (together with the LTE case) and the differences between the inferred magnetic field strength $B$ ), inclination $\gamma$ ) and azimuth $\phi$ ) were studied."
+ No significant changes were found for y and ὁ between the P and non-f inversions. however. the field strength presented some variations.," No significant changes were found for $\gamma$ and $\phi$ between the $\beta$ and $\beta$ inversions, however, the field strength presented some variations."
+ Differences with an rms value of 100 GG. in the spine. and of 150 GG. in the diffuse filament region in the upper part of the maps. were found.," Differences with an rms value of $100$ G, in the spine, and of $150$ G, in the diffuse filament region in the upper part of the maps, were found."
+ We also studied the behavior of the LOS velocities taking into account the various options for the £ coefficients., We also studied the behavior of the LOS velocities taking into account the various options for the $\beta$ coefficients.
+ The rms changes found in the spine and in the orphan penumbrae were of around 0.1 and 0.2 ss?!. respectively.," The rms changes found in the spine and in the orphan penumbrae were of around $0.1$ and $0.2$ $^{-1}$, respectively."
+ These values are small. similar to other errors arising from photon noise or systematies from the velocity calibration.," These values are small, similar to other errors arising from photon noise or systematics from the velocity calibration."
+ NLTE effects of the lline are certainly non-negligible when estimating the temperature (7) stratification. which was found in our tests. but this has no impact on our purely magnetic study.," NLTE effects of the line are certainly non-negligible when estimating the temperature $T$ ) stratification, which was found in our tests, but this has no impact on our purely magnetic study."
+ In general. the temperatures abovethe range of logr=-0.5 and -1.0 ," In general, the temperatures abovethe range of $\log\tau = -0.5$ and $-1.0$ "
+"characteristic timescale fDA of diffusive shock. Production of ~GeV cosudc rays ou the tiuuescale #, estimated above is much faster than advection through the dissipation region. which happens on a characteristic time where we have imade use of the fact that in the comoving frame the dissipation region is roughly spherical (R/T~ ","characteristic timescale $t_\ditto{DSA}$ of diffusive shock Production of $\sim\rm{GeV}$ cosmic rays on the timescale $t_\ditto{heat}$ estimated above is much faster than advection through the dissipation region, which happens on a characteristic time where we have made use of the fact that in the comoving frame the dissipation region is roughly spherical $R_\ditto{diss}/\Gamma\sim R_\ditto{diss}\theta$ )."
+This is of the sanie. order as the adiabatic cooling #0).time duc to the jet. lateral expalision: where ὃςcDO is the jet expausion velocity iu the conioviue frame., This is of the same order as the adiabatic cooling time due to the jet lateral expansion: where $\beta_\ditto{exp}\simeq\Gamma\theta$ is the jet expansion velocity in the comoving frame.
+ As Table P. shows. advection and adiabatic cooling are the fastest loss processes for 1-10 GeV ious.," As Table \ref{tab1} shows, advection and adiabatic cooling are the fastest loss processes for 1-10 GeV ions."
+" The table also includes the characteristic timescales for the following processes: radiative cooling via svuchrotrou aud inverse Compton emission. clastic collisions with other protous (pp collisious. with cross section G,,2x at eucrgies of a few GeV and inclasticity kyc 1/2) and inelastic colliious. with background. photos (p collisions) — resulting iu photoaneson production and/or pair production via the Doethe-IIeitler effect7)."," The table also includes the characteristic timescales for the following processes: radiative cooling via synchrotron and inverse Compton emission, inelastic collisions with other protons $p\,p$ collisions, with cross section $\sigma_{pp}\simeq3\times10^{-26}\unit{cm^2}$ at energies of a few GeV and inelasticity $k_{pp}\simeq1/2$ ) and inelastic collisions with background photons $p\,\gamma$ collisions) – resulting in photo-meson production and/or pair production via the Bethe-Heitler effect."
+" We also discuss the limiting case iu which the proton thermal energv stored in the dissipation region (sce Cr in (L1))) is efitcicutly transferred to the cluitting electrons. which then cool via svuchrotron aux inverse Compton at the observed bolometric huninosity E,,,~NICWeeresLo the corresponding: timescale. is: a conservative lower Μπιτ which holds regardless of ""uncertainties in the efficacy of energy. exchanec between protons aud shockclectrousa"," We also discuss the limiting case in which the proton thermal energy stored in the dissipation region (see $U'_\ditto{th}$ in ) is efficiently transferred to the emitting electrons, which then cool via synchrotron and inverse Compton at the observed bolometric luminosity $L_\ditto{Bol}\sim10^{48}\unit{erg\,s^{-1}}$; the corresponding timescale is a conservative lower limit which holds regardless of uncertainties in the efficacy of energy exchange between protons and electrons."
+cceleration In the tangled and inhomogeneous fields of the dissipation region. a fraction of the shock-heated protons niav cficiently scatter with resonant maguctic fluctuations and diffuse out of the jet before being advected away with the jet flow.," In the tangled and inhomogeneous fields of the dissipation region, a fraction of the shock-heated protons may efficiently scatter with resonant magnetic fluctuations and diffuse out of the jet before being advected away with the jet flow."
+" The cosmic rax diffusion timescale across the jet is given by(38) where r, is the cosmic ray optical The resonant magnetic fluctuations providing the cosmic ray scattering may be cmbedded in the jet flow with. c.g.. a IKohlinogorov spectrin. or generated iu the dissipation region bv the accelerated cosmic ravs themselves via the so called “streamune instability”)."," The cosmic ray diffusion timescale across the jet is given by, where $\tau_\ditto{J}$ is the cosmic ray optical The resonant magnetic fluctuations providing the cosmic ray scattering may be embedded in the jet flow with, e.g., a Kolmogorov spectrum, or generated in the dissipation region by the accelerated cosmic rays themselves via the so called “streaming instability”."
+. A concise review of both scattering processes aud their ability to describe the cosmic ταν distribution fiction iu the Millkv. Way is given iu 83 of?., A concise review of both scattering processes and their ability to describe the cosmic ray distribution function in the Milky Way is given in 3 of.
+. We assume that the scattering moechanisiis which account for the interstellar cosmic rav optical depth in the Galaxy operate in the core of quasar jets as well and we emplov the same scaliues as in2., We assume that the scattering mechanisms which account for the interstellar cosmic ray optical depth in the Galaxy operate in the core of quasar jets as well and we employ the same scalings as in.
+. This may seem as a huge extrapolation. but in the absence of direct observational constraints ou the cores of radio-loud quasars. we utilize this assuuption for lack of a better choice.," This may seem as a huge extrapolation, but in the absence of direct observational constraints on the cores of radio-loud quasars, we utilize this assumption for lack of a better choice."
+" We now estimate the value of T, expected for the two scattering processes iueutioned above.", We now estimate the value of $\tau_\ditto{J}$ expected for the two scattering processes mentioned above.
+" Several theoretical models have been proposed to account for specific features of SPMI as observed in the chromosphere and possibly in the corona, but a comprehensive theory is still lacking."," Several theoretical models have been proposed to account for specific features of SPMI as observed in the chromosphere and possibly in the corona, but a comprehensive theory is still lacking."
+" To explain the presence of chromospheric hot spots rotating in phase with the planet, the first idea was that of considering a scaled version of the unipolar induction model proposed for the Jupiter-Io system."," To explain the presence of chromospheric hot spots rotating in phase with the planet, the first idea was that of considering a scaled version of the unipolar induction model proposed for the Jupiter-Io system."
+" Observations in the UV has revealed two bright spots in the Northern and the Southern hemispheres of Jupiter, respectively, located at the footpoints of the loops connecting the magnetic poles of Jupiter with Io."," Observations in the UV has revealed two bright spots in the Northern and the Southern hemispheres of Jupiter, respectively, located at the footpoints of the loops connecting the magnetic poles of Jupiter with Io."
+" A flux of Alfven waves is excited by the motion of Io across the magnetic field lines of Jupiter’s large-scale dipole field and the energy of the waves is conducted down to the loop footpoints producing the emission (see,e.g.,Zarka2007,andreferences therein).."," A flux of Alfven waves is excited by the motion of Io across the magnetic field lines of Jupiter's large-scale dipole field and the energy of the waves is conducted down to the loop footpoints producing the emission \citep[see, e.g., ][ and references therein]{Zarka07}."
+ This mechanism can work even if Io has no intrinsic magnetic field because it is sufficient its orbital motion across the Jupiter's field lines and some surface electric conductivity to excite the Alfven waves., This mechanism can work even if Io has no intrinsic magnetic field because it is sufficient its orbital motion across the Jupiter's field lines and some surface electric conductivity to excite the Alfven waves.
+assess (heir importance for the overall ABB formation.,assess their importance for the overall KBB formation.
+ Our paper is structured as follows: In 82 we outline our assumptions. explain our choice of parameters and define variables (hat will be used throughout this paper.," Our paper is structured as follows: In 2 we outline our assumptions, explain our choice of parameters and define variables that will be used throughout this paper."
+" We calculate the L and L?s formation rates for sub-ITill MBO velocities in 83 and 84 respectively,", We calculate the $L^3$ and $L^2s$ formation rates for sub-Hill KBO velocities in 3 and 4 respectively.
+ We compare the L?s and L formation rates in the sub-Iill velocity regime in 85., We compare the $L^2s$ and $L^3$ formation rates in the sub-Hill velocity regime in 5.
+ In &6 we discuss how (hese formation rates are mocdilied for super-Ilill IKRBO velocities., In 6 we discuss how these formation rates are modified for super-Hill KBO velocities.
+ The frequency of long-livecl transient binaries and their significance for the overall ABB formation is caleulated in $7., The frequency of long-lived transient binaries and their significance for the overall KBB formation is calculated in 7.
+ sSunmmnmary and conclusions follow in 38., Summary and conclusions follow in 8.
+ The Hill radius denotes the distance from a IKDO at which the tidal forees due to the Sun and the gravitational force due to the NBO. both acting on a test particle. ave in equilibrium.," The Hill radius denotes the distance from a KBO at which the tidal forces due to the Sun and the gravitational force due to the KBO, both acting on a test particle, are in equilibrium."
+ It is given by where @ is (he semi-major axis ancl 1M the mass of the KBO., It is given by where $a$ is the semi-major axis and $M$ the mass of the KBO.
+ AZ. is the mass of the sun., $M_{\sun}$ is the mass of the sun.
+ We use the 6wo-group approximation (Goldreichetal.2002.2004) which consists of the identification of two groups of objects. small ones. (hat contain most of the total mass with surface mass densitv o. and large ones. that contain onlv a small fraction of the total mass with surface mass density X«c.," We use the `two-group approximation' \citep{GLS02,GLS04}
+ which consists of the identification of two groups of objects, small ones, that contain most of the total mass with surface mass density $\sigma$, and large ones, that contain only a small fraction of the total mass with surface mass density $\Sigma \ll \sigma $."
+ We assume σ~0.3gem7 which is the extrapolation of ihe minimum-nmass solar nebular to a heliocentric distance of 40AU., We assume $\sigma \sim 0.3 \rm{g~cm^{-2}}$ which is the extrapolation of the minimum-mass solar nebular to a heliocentric distance of $40\rm{AU}$.
+ Estimates from current Ixuiper Belt survevs (Trujillo&Brown2003:Trujilloetal.2001) vield X—3x10tgem? for IKBOs with radii of 2.—100kin.," Estimates from current Kuiper Belt surveys \citep{TB03,TJL01} yield $\Sigma \sim 3 \times 10^{-4} \rm{g~cm^{-2}}$ for KBOs with radii of $R \sim 100~\rm{km}$."
+ We use this value of X. assuming that. X curing the formation of INBDs was (the same as it is now.," We use this value of $\Sigma$, assuming that $\Sigma$ during the formation of KBBs was the same as it is now."
+ Our choice for X and σ is also consistent with results from numerical coagulation simulations bv IXenvon&Lim(1999)., Our choice for $\Sigma$ and $\sigma$ is also consistent with results from numerical coagulation simulations by \citet{KL99}.
+. Large bodies grow bv the accretion of small bodies., Large bodies grow by the accretion of small bodies.
+ Large INDOs viscously stir the small bodies. increasing the small bodies! velocity dispersion uv.," Large KBOs viscously stir the small bodies, increasing the small bodies' velocity dispersion $u$."
+ As a result 4 grows on (he same timescale as 2 provided that mutual collisions among the small bodies are not vet important., As a result $u$ grows on the same timescale as $R$ provided that mutual collisions among the small bodies are not yet important.
+ In this case. u is given by where a=R/Ry10! at 40AU (Goldreichοἱal.2002).," In this case, $u$ is given by where $\alpha = R/R_{H}\sim 10^{-4}$ at $40\rm{AU}$ \citep{GLS02}."
+". vy is the LHL velocity of the laree bodies which is given by ey=OR, where € is the orbital Ireeuency. around the sun.", $v_H$ is the Hill velocity of the large bodies which is given by $v_H = \Omega R_H$ where $\Omega$ is the orbital frequency around the sun.
+ The velocity ο of large INDOs increases due to mutual viscous stirring. but is damped," The velocity $v$ of large KBOs increases due to mutual viscous stirring, but is damped"
+ | (OO right)| 7003) z Then. using //ce—T— (3) MA) right)ym )| one πια» fel,") = - + ( + ^4) ) = + + ^3) Then, using /c = ) ^4) )^2/c^2 = ) ^5) one finds ) = /c+ + ( ) ^4)"
+"but only by a small margin since μεν=1.046. Residual material in the gap leads to additional eccentricity damping unless X(a,)/X(r4)€ομως©1/3.","but only by a small margin since ${\cal C}_{CR}/{\cal C}_{eLR}=1.046$ Residual material in the gap leads to additional eccentricity damping unless $\Sigma(a_p)/\Sigma(r_1)< {\cal
+C}_{CR}/{\cal C}_{cLR} \approx 1/3$."
+ However. such a steep density. gradient near the center of the gap is incompatible with our assumption that (he disk maintains Keplerian rotation.," However, such a steep density gradient near the center of the gap is incompatible with our assumption that the disk maintains Keplerian rotation."
+ Next we exanune how departures from I[xeplerian rotation might alleet the balance between eccentricily driving and damping., Next we examine how departures from Keplerian rotation might affect the balance between eccentricity driving and damping.
+ Gas pressure perturbs the clisk’s epievclic frequency., Gas pressure perturbs the disk's epicyclic frequency.
+ ]t decreases &j at the outskirts of a gap and increases it around the Provided (he gap is sufficiently clean. only resonances in ils outer parts need be considered.," It decreases $\kappa_d$ at the outskirts of a gap and increases it around the Provided the gap is sufficiently clean, only resonances in its outer parts need be considered."
+ Numerical simulations of gap formation by panets show that this limit can be achieved. at least in two-dimensional disks.," Numerical simulations of gap formation by planets show that this limit can be achieved, at least in two-dimensional disks."
+ This requires ]«dIn€(r/hy?., This requires $1\ll |d^2\ln\Sigma/d\ln r^2|\lesssim (r/h)^2$.
+ The upper limit is sel bv the requirements (hat principal Liidblad resonances are located away [rom the planets seminiajor axis ancl by (he Ravleigh stability criterion., The upper limit is set by the requirements that principal Lindblad resonances are located away from the planet's semimajor axis and by the Rayleigh stability criterion.
+ It motivates us to investigate the eccentricity evolution of a planet that resides in a gap where the epievclic frequency. is much smaller than the orbital angular velocity., It motivates us to investigate the eccentricity evolution of a planet that resides in a gap where the epicyclic frequency is much smaller than the orbital angular velocity.
+" As αιο—0. dQy/dr4—2O04/r. and the positions of the inner and outer Lindblad resonances associated wilh a given ὦμ1, potential component. approach. from opposite sides. the location of the corresponding corotation resonance."," As $\kappa_d/\Omega_d\to 0$, $d\Omega_d/dr\to -2\Omega_d/r$, and the positions of the inner and outer Lindblad resonances associated with a given $\phi_{m-1, m}$ potential component approach, from opposite sides, the location of the corresponding corotation resonance."
+ This simplilies comparison of (he influences of the different resonances on eccentricity evolution., This simplifies comparison of the influences of the different resonances on eccentricity evolution.
+" Inserting these approximations into the expressions for the torques at first order resonances given by. (1919).. we obtain and Note that in the limit &j/,—0. each Lindblad torque dominates (he corotation torque."," Inserting these approximations into the expressions for the torques at first order resonances given by \cite{GOT79}, we obtain and Note that in the limit $\kappa_d/\Omega_d\to 0$, each Lindblad torque dominates the corotation torque."
+" llowever. the Lindblad torques have opposite signs and when summed making use of their separations crmgg/2mf3,we discover that"," However, the Lindblad torques have opposite signs and when summed making use of their separations $\mp r\kappa_d/2m\Omega_d$ we discover that"
+The paper is organized as follows.,The paper is organized as follows.
+ In Section 2? we briefly review the formalism of the equations and the numerical method used to solve them., In Section \ref{sec::formulation} we briefly review the formalism of the equations and the numerical method used to solve them.
+ Then we describe the new functional form of the rotation profile., Then we describe the new functional form of the rotation profile.
+ In Section ?? νο construct equilibrium sequences of rotating stars using the new rotation profile., In Section \ref{sec::results} we construct equilibrium sequences of rotating stars using the new rotation profile.
+ Discussions are given in the final section., Discussions are given in the final section.
+ We construct configurations for an and rotating perfect fluid in general relativity., We construct configurations for an and rotating perfect fluid in general relativity.
+ The spacetime is assumed to be asymptotically flat and the flow 1s assumed to be circular (the velocity is only in the azimuthal direction)., The spacetime is assumed to be asymptotically flat and the flow is assumed to be circular (the velocity is only in the azimuthal direction).
+ In this case we have two Killing vectors (?) and can choose a coordinate system in such a way that the metric of the spacetime is written as (e.g.. 2)). where the spherical polar coordinates (7.0.6) are used.," In this case we have two Killing vectors \citep{Bardeen1970}
+ and can choose a coordinate system in such a way that the metric of the spacetime is written as (e.g., \citet{Komatsu1989a}) ), where the spherical polar coordinates $(r,\theta,\phi)$ are used."
+ The metric potentials c.f.v and w are functions of r and 6 only.," The metric potentials $\alpha,\beta,\nu$ and $\omega$ are functions of $r$ and $\theta$ only."
+" The energy momentum tensor Τ0 of a perfect fluid is where e. p and i"" are the total energy density. the pressure and the four velocity. respectively."," The energy momentum tensor $T^{ab}$ of a perfect fluid is where $\epsilon$, $p$ and $u^a$ are the total energy density, the pressure and the four velocity, respectively."
+" The basic equations are: 1) Einstein's equation for the metric potentials Gy,=SzT,4. 2) rest mass conservatior γρ=06. which ts trivially satisfied. under the present assumptions and 3) stress-energy conservation V,7°=0."," The basic equations are: 1) Einstein's equation for the metric potentials $G_{ab}=8\pi T_{ab}$, 2) rest mass conservation $\nabla_a(\rho u^a)=0$, which is trivially satisfied under the present assumptions and 3) stress-energy conservation $\nabla_bT^{ab}=0$."
+ As described in ?.. the components of Einstein's equatior are cast into 4+ equations for the potentials. three of which are elliptic partial differential equations.," As described in \citet{Komatsu1989a}, the components of Einstein's equation are cast into 4 equations for the potentials, three of which are elliptic partial differential equations."
+ They are transformec into convenient integral equations by using appropriate Green's functions., They are transformed into convenient integral equations by using appropriate Green's functions.
+ The spatial linear velocity V of the flow with respect to ar observer with zero angular momentum is given by where Q=a/u is the angular frequency of the fluid measured in the asymptotic inertial frame., The spatial linear velocity $V$ of the flow with respect to an observer with zero angular momentum is given by where $\Omega=u^\varphi/u^t$ is the angular frequency of the fluid measured in the asymptotic inertial frame.
+ We next introduce the specific angular momentum j: Using these quantities. we can write the equilibrium equations for a stationary configuration as Assuming the fluid to be barotropic. the condition. of integrability for Eq.(5)). can be written as by introducing an arbitrary functional g. which is detailed in the next subsection.," We next introduce the specific angular momentum $j$: Using these quantities, we can write the equilibrium equations for a stationary configuration as Assuming the fluid to be barotropic, the condition of integrability for \ref{hydrostatic eq}) ), can be written as by introducing an arbitrary functional $g$, which is detailed in the next subsection."
+ The first integral of motion for a stationary solution can be written as ερThe components of Einstein's equation and the first integral of hydrostatic equation can be iteratively solved as described in 9 To exploit the first integral of the hydrostatic equation. we need to impose the integrability condition for Eq.(6)).," The first integral of motion for a stationary solution can be written as The components of Einstein's equation and the first integral of hydrostatic equation can be iteratively solved as described in \citet{Komatsu1989a}
+ To exploit the first integral of the hydrostatic equation, we need to impose the integrability condition for \ref{integrability of hydrostat}) )."
+ Different choices of functional form g(Q) would lead to various classes of differential-rotation profiles., Different choices of functional form $g(\Omega)$ would lead to various classes of differential-rotation profiles.
+" All previous studies of differentially-rotating relativistic stellar models assume the simplest linear functional form. where A 1s a constant with the dimensions of a length and Q,. is the angular velocity on the axis of rotation."," All previous studies of differentially-rotating relativistic stellar models assume the simplest linear functional form, where $A$ is a constant with the dimensions of a length and $\Omega_c$ is the angular velocity on the axis of rotation."
+ This choice is termed as 7j-const., This choice is termed as “j-const.
+" law"" since the Newtonian limit of the specific angular momentum / is that of “j-const."," law"" since the Newtonian limit of the specific angular momentum $j$ is that of “j-const."
+ law” in Newtonian rotating stars (?)..," law"" in Newtonian rotating stars \citep{EriguchiMueller1985}."
+ We here introduce a more flexible form of the rotation profile as where a. Ro and (2. are constants.," We here introduce a more flexible form of the rotation profile as where $\alpha$, $R_0$ and $\Omega_c$ are constants."
+" The corresponding specific angular momentum is The physical significance of the parameters a. Ro and Q,. in the profile is easily seen 1f we consider its Newtonian limit."," The corresponding specific angular momentum is The physical significance of the parameters $\alpha$, $R_0$ and $\Omega_c$ in the profile is easily seen if we consider its Newtonian limit."
+ Then the angular frequency is written as where R=rsind., Then the angular frequency is written as where $R=r\sin\theta$.
+ For R«Ro. we have Q~Q.. while for Rc»Ro. This means that the rotation profile has an inner plateau inside R~Ro and a power law envelope for R>>Ro.," For $R \ll R_0$, we have $\Omega\sim\Omega_c$, while for $R\gg R_0$, This means that the rotation profile has an inner plateau inside $R\sim R_0$ and a power law envelope for $R\gg R_0$."
+ The introduction of the index « is an important feature of the new profile: as shown in Table I. itis possible to reproduce different rotation laws by choosing different values for c.," The introduction of the index $\alpha$ is an important feature of the new profile: as shown in Table 1, it is possible to reproduce different rotation laws by choosing different values for $\alpha$."
+ For some interesting cases (a€σαν22. 4]. it ts possible to integrate analytically the expression for g(Q) (see Appendix).," For some interesting cases $\alpha \in \left\{-1,-2,-4\right\} $ ), it is possible to integrate analytically the expression for $g(\Omega)$ (see Appendix)."
+ It is important to note that for e=-I the j- type of law is recovered — the two functional forms must agree in the limit of weak gravity., It is important to note that for $\alpha=-1$ the $j$ -constant type of law is recovered – the two functional forms must agree in the limit of weak gravity.
+ This particular case is marginally stable under Rayleigh's local stability criterion of axisymmetric instability. which states that the specific angular momentum must not decrease outward in a stable star.," This particular case is marginally stable under Rayleigh's local stability criterion of axisymmetric instability, which states that the specific angular momentum must not decrease outward in a stable star."
+ Values of a smaller than —1 satisfy this condition (and hence are, Values of $\alpha$ smaller than $-1$ satisfy this condition (and hence are
+(RPL. WR).,"(RPL, KP)."
+ This is associatecdd with a transition from oscillatory to evanescent behaviour of eigenfunctions., This is associated with a transition from oscillatory to evanescent behaviour of eigenfunctions.
+ We do not plot cigenfunetions here. as an extensive discussion of their properties can be found in. RPL. WP. LOOS and 7..," We do not plot eigenfunctions here, as an extensive discussion of their properties can be found in RPL, KP, LO98 and \citet{Og98}."
+ We only mention that the eigenfunctions of the p- and e-mocles are trapped: near the surfaces and. decay towards the midplane. while the cigenfunetions of the r-mode are concentrated near the midplane and. decay towards: the surfaces.," We only mention that the eigenfunctions of the p- and g-modes are trapped near the surfaces and decay towards the midplane, while the eigenfunctions of the r-mode are concentrated near the midplane and decay towards the surfaces."
+ As for the cigenfunetions of the fundamental [-mode. they have no nodes anc monotonically decay. from the surfaces to the midplane.," As for the eigenfunctions of the fundamental f-mode, they have no nodes and monotonically decay from the surfaces to the midplane."
+ Below we show that these properties of eigenfunctions are altered due to self-eravity., Below we show that these properties of eigenfunctions are altered due to self-gravity.
+ Specifically. the number of nodes along cach branch of the dispersion diagrams. which is preserved in the non-self-eravitating limit. is no longer constant in the presence of sel-eravity.," Specifically, the number of nodes along each branch of the dispersion diagrams, which is preserved in the non-self-gravitating limit, is no longer constant in the presence of self-gravity."
+ Phe frequencies of the p- and r-modes increase with mode number., The frequencies of the p- and r-modes increase with mode number.
+ The frequencies ancl growth rates. of the convectively stable ancl unstable g-mocdes. respectively. decrease with the mode number.," The frequencies and growth rates of the convectively stable and unstable g-modes, respectively, decrease with the mode number."
+ These are the well-known general properties of vertical modes in polvtropic disces., These are the well-known general properties of vertical modes in polytropic discs.
+" Although cach mode tvpe has its dominant. restoring orce. one out of the above mentioned four types (compressibility. surface gravity. buovaney. inertial forces). he other three forces also contribute to a certain degree for small racial wavenumbers Afx,1."," Although each mode type has its dominant restoring force, one out of the above mentioned four types (compressibility, surface gravity, buoyancy, inertial forces), the other three forces also contribute to a certain degree for small radial wavenumbers $kh \lsim 1$."
+ For example. the p- and -modes are modified both by rotation/inertial forces and )uovancy. the [mode is also modified. by compressibility. he e-mode is mocified by rotation and compressibility and he r-mode is modified by buovaney and: compressibility.," For example, the p- and f-modes are modified both by rotation/inertial forces and buoyancy, the f-mode is also modified by compressibility, the g-mode is modified by rotation and compressibility and the r-mode is modified by buoyancy and compressibility."
+ For the sake of comparison with the relevant. result. from oevious studies. we are particularly interested in to what degree compressibility. mocifies generally non-compressive -. ge and r-modes.," For the sake of comparison with the relevant result from previous studies, we are particularly interested in to what degree compressibility modifies generally non-compressive f-, g- and r-modes."
+ So. we decided to juxtapose the dispersion. diagrams for these modes computed. separately in the compressible and. incompressible cases.," So, we decided to juxtapose the dispersion diagrams for these modes computed separately in the compressible and incompressible cases."
+ We take the incompressible limit by formally letting the acdiabatic index eo to infinity. x (?).," We take the incompressible limit by formally letting the adiabatic index go to infinity, $\gamma\rightarrow \infty$ \citep{Og98}."
+ After that equations (16-18) without self-gravity take the form We solved. these equations with the same boundary conditions and found the dispersion. diagrams of the [-. &- and r-modes that survive in this limit. (obviously. the acoustic pmocde disappears).," After that equations (16-18) without self-gravity take the form We solved these equations with the same boundary conditions and found the dispersion diagrams of the f-, g- and r-modes that survive in this limit (obviously, the acoustic p-mode disappears)."
+ The results are plotted in Fig., The results are plotted in Fig.
+ 2 with dashed lines., 2 with dashed lines.
+ Notice that by taking the incompressible limit in this manner. we have been able to retain the [-mocoe. as expected.," Notice that by taking the incompressible limit in this manner, we have been able to retain the f-mode, as expected."
+ LInstead. KP set. the density perturbations to zero in the linearized continuity equation (anelastic approximation) that resulted. in the. [mode disappearing in their incompressible limit.," Instead, KP set the density perturbations to zero in the linearized continuity equation (anelastic approximation) that resulted in the f-mode disappearing in their incompressible limit."
+ We see that compressibility most stronely alfects the convectively stable and unstable e-mode branches with small mode numbers (Figs., We see that compressibility most strongly affects the convectively stable and unstable g-mode branches with small mode numbers (Figs.
+ 2b.2g). or equivalently with vertical extent comparable το the disc height. even at large racial wavenumbers (?)..," 2b,2g), or equivalently with vertical extent comparable to the disc height, even at large radial wavenumbers \citep{Og98}."
+ The frequencies of the and r-modes do not change. much. indicating that these moces are nearly incompressible.," The frequencies of the f- and r-modes do not change much, indicating that these modes are nearly incompressible."
+ Lavine characterized all types of axisvnimetric normal modes in a stratified disc. let us now compute the dispersion diagrams taking into account sell-eravity in the perturbation equations and in the equilibrium.," Having characterized all types of axisymmetric normal modes in a stratified disc, let us now compute the dispersion diagrams taking into account self-gravity in the perturbation equations and in the equilibrium."
+" ""This will allow us to understand. how the frequencies and the structure of the cigenfunetions of the above-deseribec mode types are altered by selberavity. which mode acquires the largest) positive growth rate in the presence of scll-eravity ane. thus determines the onset criterion and nature of gravitational instability of a disc."," This will allow us to understand how the frequencies and the structure of the eigenfunctions of the above-described mode types are altered by self-gravity, which mode acquires the largest positive growth rate in the presence of self-gravity and, thus determines the onset criterion and nature of gravitational instability of a disc."
+ In other words. we return to the boundary value problem formulated. in Section. 2. which is represented by equations (16-18). supplemented with boundary concitions (21).(22) at the surface and. conditions (19).(20) of the even and. odd. symmetry. of a solution at the midplane.," In other words, we return to the boundary value problem formulated in Section 2, which is represented by equations (16-18) supplemented with boundary conditions (21),(22) at the surface and conditions (19),(20) of the even and odd symmetry of a solution at the midplane."
+ For cach Q and s. we first. determine the corresponding vertical distribution of the equilibrium quantities in equations (16-18). po.cL.du.NG. with height as described in Section 2 and then based on this compute the normal mocles.," For each $Q$ and $s$, we first determine the corresponding vertical distribution of the equilibrium quantities in equations (16-18), $\rho_0, c_s^2, g_0, N_0^2$, with height as described in Section 2 and then based on this compute the normal modes."
+ Figures 3.4.5 and 6 show the typical dispersion diagrams [or the p-. [-. g- and r-modes in a selí-gravitating disc [or subadiabatic. acliabatic and. superacliahatic vertical stratifications.," Figures 3,4,5 and 6 show the typical dispersion diagrams for the p-, f-, g- and r-modes in a self-gravitating disc for subadiabatic, adiabatic and superadiabatic vertical stratifications."
+ We separately plot the dispersion diagrams of the even ancl odd: parities for cach mode tvpe to clearly see the inlluence of self-gravity on them. which. as evident [rom these figures. depends on the mode parity.," We separately plot the dispersion diagrams of the even and odd parities for each mode type to clearly see the influence of self-gravity on them, which, as evident from these figures, depends on the mode parity."
+ Unlike in the non-self-eravitating limit. the number of nodes of the vertical velocity and: pressure perturbations in the presence of scl-eravity ave not. preserved. along cach moce branch.," Unlike in the non-self-gravitating limit, the number of nodes of the vertical velocity and pressure perturbations in the presence of self-gravity are not preserved along each mode branch."
+ However. as we will see below. at large &. the influence of self- on the mode dynamics is small and the dispersion diagrams merge with their non-self-gravitating counterparts (shown with dashed curves in Figs.," However, as we will see below, at large $k$, the influence of self-gravity on the mode dynamics is small and the dispersion diagrams merge with their non-self-gravitating counterparts (shown with dashed curves in Figs."
+ 3. 5).," 3, 5)."
+ So. the naming of the modes for smaller A. where the ellect of self-gravity. is important. is done by continuity with the large-A limit for cach mode branch.," So, the naming of the modes for smaller $k$, where the effect of self-gravity is important, is done by continuity with the $k$ limit for each mode branch."
+ Consider first the subadiabatic case (Fig., Consider first the subadiabatic case (Fig.
+ 3). where we again have the p- [-. g- and r-modes. but their dispersion diagrams are mocified/shifted. from their. non-sell-gravitating counterparts towards lower values due to self-eravity.," 3), where we again have the p-, f-, g- and r-modes, but their dispersion diagrams are modified/shifted from their non-self-gravitating counterparts towards lower values due to self-gravity."
+ As illustrated in Figs., As illustrated in Figs.
+ 3a.3d. self-gravity reduces the frequencies. of all branches of the p- and [2mocdes. for both even and odd parities. but it more allects the f-2mocdes.," 3a,3d, self-gravity reduces the frequencies of all branches of the p- and f-modes, for both even and odd parities, but it more affects the f-modes."
+ The situation for the &- and r-modes is different (Figs., The situation for the g- and r-modes is different (Figs.
+ 3b.3c.3e.3PD): only the frequencies of the even g- and rmodes are reduced by self-gravity mostly for racial wavenumbers in the range 0x&2 (dips on the corresponding dispersion curves in Figs.," 3b,3c,3e,3f): only the frequencies of the even g- and r-modes are reduced by self-gravity mostly for radial wavenumbers in the range $0\leq k \leq 2$ (dips on the corresponding dispersion curves in Figs."
+ 3b.3c indicating deviations [from the non-sell-eravitating dashed ones). whereas the frequencies of the odd. & and r-mocdes are almost unallected. by scll-eravity.," 3b,3c indicating deviations from the non-self-gravitating dashed ones), whereas the frequencies of the odd g- and r-modes are almost unaffected by self-gravity."
+ The frequencies of the fundamental f-modes. the first. few branches of the p-niocles ancl also the first. few branches of the even ο- and rmoces are reduced. noticeablv.," The frequencies of the fundamental f-modes, the first few branches of the p-modes and also the first few branches of the even g- and r-modes are reduced noticeably."
+. With increasing mode number and/or radial wavenumber. the ellect of self-gravitv on the eigenfrequencies gradually. falls ol. because the corresponding eigenfunctions become of," With increasing mode number and/or radial wavenumber, the effect of self-gravity on the eigenfrequencies gradually falls off, because the corresponding eigenfunctions become of"
+and the ecceutricitics between 0 and 0.035.,and the eccentricities between 0 and 0.035.
+ Realistic changes in the set of iuclinations iux eccentricities do not change the fnal result. except for the fraction of bodies effectively trapped ii resonances.," Realistic changes in the set of inclinations and eccentricities do not change the final result, except for the fraction of bodies effectively trapped in resonances."
+" The mass of the planet has been taken tween Jupiter and Saturn mass,", The mass of the planet has been taken between Jupiter and Saturn mass.
+ A ess lnassive planet gives less efficient trapping., A less massive planet gives less efficient trapping.
+ The eccentricity of the planet allows to test the possible influences on the asvunuetry of the dust disk produced by the bodies trapped in “evaporating orbit”., The eccentricity of the planet allows to test the possible influences on the asymmetry of the dust disk produced by the bodies trapped in “evaporating orbit”.
+ The conclusion of severa test smauulatious can be stbunarized as follows: - The 1:2 resonance is a very efficient resonance for trapping. even with a very fast migration (7~LO’ vears).," The conclusion of several test simulations can be summarized as follows: - The 1:2 resonance is a very efficient resonance for trapping, even with a very fast migration $\tau \sim 10^6$ years)."
+ But the variation of the periastron distance is too small to allow the evaporation (Aq=q;iaw10 AU). -, But the variation of the periastron distance is too small to allow the evaporation $\Delta q=q_f-q_i\approx - 10$ AU). -
+ The 1:5 ECSOMALLCE ds cticient if rzLO‘ vears and can give a larec decrease of the periastron (Aqx30 lo AU)., The 1:3 resonance is efficient if $\tau \ga 10^7$ years and can give a large decrease of the periastron $\Delta q\approx -30$ – $-40$ AU).
+ But with a planet eccentricity of 0.05. no azimuthal asvuuuetry bas been observed on the distribution of the periastron of the trapped bodies.," But with a planet eccentricity of 0.05, no azimuthal asymmetry has been observed on the distribution of the periastron of the trapped bodies."
+" Finally, there is no chauge iu the distribution of the inclination which remain low. -"," Finally, there is no change in the distribution of the inclination which remain low. -"
+ The 1: Lresonance give interesting results preseuted iu Fie., The 1:4 resonance give interesting results presented in Fig.
+ E. aud 5.., \ref{qei2tres4} and \ref{pna2t}.
+" The trapping has been fone to be efficient with rather extreme parameters: the mass of the planet uust be M,<0.5Mjj. where Mj is the uiass of Jupiter. he mueration rate mist o low (r=5+10* vears)."," The trapping has been found to be efficient with rather extreme parameters: the mass of the planet must be $M_p \ga 0.5 M_J$, where $M_J$ is the mass of Jupiter, the migration rate must be low $\tau \ga 5\cdot 10^{7}$ years)."
+ The eccentricity of the planet vas been taken to be 0.05., The eccentricity of the planet has been taken to be 0.05.
+ With hese couditions. 19/500 bodies have been found to be rapped in the 1:1 resonance. aud the periastron decreased wAgqzlO AU.," With these conditions, 19/500 bodies have been found to be trapped in the 1:4 resonance, and the periastron decreased by $\Delta q\approx -40$ AU."
+ A significant increase of the inclination as been observed after few τας well as a large asvuuuctry in the distribution of the periastrous longitude., A significant increase of the inclination has been observed after few $\tau$ as well as a large asymmetry in the distribution of the periastrons longitude.
+ Tn evaluating the evaporation rate οἳ this configuration. we find tha the periastron decrease due to the 1:l resonance presseuts two favorable particularitics. -," In evaluating the evaporation rate of this configuration, we find that the periastron decrease due to the 1:4 resonance sents two favorable particularities. -"
+" The 1:5 resonance is efficient im trapping only if the parameters of the migrating plauect are extreme with AM,£Mj.vc, ZOL aud τZ5-10* years."," The 1:5 resonance is efficient in trapping only if the parameters of the migrating planet are extreme with $M_p\ga M_J$, $e_p \ga 0.1$ and $\tau \ga 5\cdot 10^{7}$ years."
+ With these conditions. we find 10/600 bodies trapped im the resonance.," With these conditions, we find 10/600 bodies trapped in the resonance."
+ The consequence ou the CO-evaporation rate sa fiction of time is very simular to the one found in the 1:1 resonance., The consequence on the CO-evaporation rate as a function of time is very similar to the one found in the 1:4 resonance.
+ Among the 10 bodies trapped in the resonance one has evolved m a very eccentric orbit with large inclination., Among the 10 bodies trapped in the resonance one has evolved in a very eccentric orbit with large inclination.
+ The luk with the FEDs is possible but PAill to be investigated., The link with the FEBs is possible but still to be investigated.
+ As the dust life-time is very short (fy~105 vv iu the | ddisk but also f;—109 vr around PsA)) aud shorter than the stellar ages. the deusitv of dust observed around the main sequence stars is directly related to the actual production rate of dust.," As the dust life-time is very short $t_d\sim 10^4$ yr in the $\beta\:$ disk but also $t_d\sim 10^6$ yr around ) and shorter than the stellar ages, the density of dust observed around the main sequence stars is directly related to the actual production rate of dust."
+ The time variation of the dust production rate. and consequently of the disk deusity. have been evaluated in the sinulatious described in Sect. 5.2..," The time variation of the dust production rate, and consequently of the disk density, have been evaluated in the simulations described in Sect. \ref{migration}. ."
+ Fig., Fig.
+ 5.Γ shows the production as the fiction of time for the 1:1 resonauce., \ref{pna2t} shows the production as the function of time for the 1:4 resonance.
+ The production of dust occurs. between 3°105 and 7-10* vears which is very consistent with the estimated age of + ((Crifo et al., The production of dust occurs between $3\cdot 10^{7}$ and $7\cdot 10^{7}$ years which is very consistent with the estimated age of $\beta\:$ (Crifo et al.
+ 1997)., 1997).
+ We also conclude that the dust production rate by evaporation is trausieut and cau be large duriug phase during which the bodies trapped in the resonances are enteriug in the evaporation linüt., We also conclude that the dust production rate by evaporation is transient and can be large during phase during which the bodies trapped in the resonances are entering in the evaporation limit.
+ The )jpietoris ddisk eau be im such a state while other Vega--like stars are in quiescenut phase., The $\beta\:$ disk can be in such a state while other -like stars are in quiescent phase.
+ If the bodies are trapped in a resonance with a planet on ecceutrie orbit. there can be an asvuuuetrv in the distribution of the periastron as already seeu in Sect. [.1..," If the bodies are trapped in a resonance with a planet on eccentric orbit, there can be an asymmetry in the distribution of the periastron as already seen in Sect. \ref{random walk}."
+ For example. the Fig.," For example, the Fig."
+ 5. gives the dust production rate by the bodies trapped iu the 1:1 resonance with a planet on an eccentric orbit (ey= 0.05)., \ref{pna2t} gives the dust production rate by the bodies trapped in the 1:4 resonance with a planet on an eccentric orbit $e_p=0.05$ ).
+ The production rate is larger in the direction of the periastron of the planet than in the opposite direction., The production rate is larger in the direction of the periastron of the planet than in the opposite direction.
+ The disk thus produced must be asvuunuetrce with a larger deusitv in the direction of the apoastron of the σταιο planet., The disk thus produced must be asymmetric with a larger density in the direction of the apoastron of the migrating planet.
+ From the comparison of Fie., From the comparison of Fig.
+ Lo and 5.. we cau conclude lat. in this configuration. the production of dust takes ace between 3:10* and 7:10 vears when the inclination f the pareut bodies are still low. but already larecr than 1e initial inclination (< 27).," \ref{qei2tres4} and \ref{pna2t}, we can conclude that, in this configuration, the production of dust takes place between $3\cdot 10^{7}$ and $7\cdot 10^{7}$ years when the inclination of the parent bodies are still low, but already larger than the initial inclination $<2^\circ$ )."
+ Here. the opening angle of 16 produced disk 11ust be around 3 or | degrees.," Here, the opening angle of the produced disk must be around 3 or 4 degrees."
+ This angle is smaller than the 7 degrees observed iu 1ο) Pictoris ccase., This angle is smaller than the 7 degrees observed in the $\beta\:$ case.
+ However. we aust remember tha we have taken ouly one planet. the ascending node of its orbit is then constaut.," However, we must remember that we have taken only one planet, the ascending node of its orbit is then constant."
+ But iu a more realisticcase with several eiat plauets. the precession of the asceudiugo nodes can produce a significant increase of the paren," But in a more realisticcase with several giant planets, the precession of the ascending nodes can produce a significant increase of the parent"
+erowth. the CALC evows at a somewhat slower rate but over an extended period of time.,"growth, the CMC grows at a somewhat slower rate but over an extended period of time."
+ These are models SD_CG1583. SD_CC3083 and €GG5082. which are all eas-rich and have the lowest resolution.," These are models G15S3, G30S3 and G50S3, which are all gas-rich and have the lowest resolution."
+ In these models. the gas content of the bar has formed without a visibly pronunent CALC.," In these models, the gas content of the bar has formed without a visibly prominent CMC."
+ Instead of a rapid influx to the ceuter. the eas in the bar coutracts rather slowly. eraduallv increasing its density at the center.," Instead of a rapid influx to the center, the gas in the bar contracts rather slowly, gradually increasing its density at the center."
+ The process saturates when the eas is completely contained within the central πο0.1., The process saturates when the gas is completely contained within the central $R\sim 0.1$.
+ Models (ν.δ aud. (10 show a rate of growth that increases with time at the eud of the run., Models G8S2 and G15S3 show a rate of growth that increases with time at the end of the run.
+ This behavior is related to the late period of accelerated bar growth., This behavior is related to the late period of accelerated bar growth.
+ Aneulhu momentum evolution in the disk las Όσοι ollowed within a uuuber of Characteristic radii. namely. he CR (CJ444)) aud the disk radius. (which coutains of the disk mass by defiuition).," Angular momentum evolution in the disk has been followed within a number of characteristic radii, namely, the CR ) and the disk radius, (which contains of the disk mass by definition)."
+" The latter exhibits a clear evolutionary sequence as a ""unctiou of (c.g. Fie."," The latter exhibits a clear evolutionary sequence as a function of (e.g., Fig."
+ 7)., 7).
+" Both the amount of J lost bv he disk over ανασα] aud secular evolution decreases nonotouically frou, pure stellar disks to progressively nore gas-rich ones.", Both the amount of $J$ lost by the disk over dynamical and secular evolution decreases monotonically from pure stellar disks to progressively more gas-rich ones.
+ Du a way this is a reflection of weaker xus alonef., In a way this is a reflection of weaker bars along.
+.. Even more interesting is the evolution of (Fig., Even more interesting is the evolution of (Fig.
+ 8)., 8).
+ After the initial adjustineut. the disk within the CR is losing its angular momentuu.," After the initial adjustment, the disk within the CR is losing its angular momentum."
+ The loss is not dramatic. ~105420%. and somewhat icreascs withfo.," The loss is not dramatic, $\sim 10\%-20\%$, and somewhat increases with."
+.. This decline iu continues until the vertical buckling sets in., This decline in continues until the vertical buckling sets in.
+ For low disks. the bucklue is associated with a ucarly sudden increase In Which restores the pre-buckline value followed by a subsequent eradual decline:," For low disks, the buckling is associated with a nearly sudden increase in which restores the pre-buckling value followed by a subsequent gradual decline."
+ Cas-rich disk show no sudden merease m but rather a slow decline., Gas-rich disk show no sudden increase in but rather a slow decline.
+ stays quite fat until f£~150. a long time after the buckling. as can bee seen from Fig.," stays quite flat until $t\sim 150$, a long time after the buckling, as can bee seen from Fig."
+ 8 (with the exception of a σπα initial decline and increase. as mentioned above).," 8 (with the exception of a small initial decline and increase, as mentioned above)."
+ Interestingly. this time corresponds to ji linthegas— poormodels( Fig.," Interestingly, this time corresponds to $ < 1$ in the gas-poor models (Fig."
+ D., 4).
+CrossingR = lborderaf feetsthobargrouwth( Fig., Crossing $ = 1$ border affects the bar growth (Fig.
+ 3). whichsaturatesinimediatel," 3), which saturates immediately."
+" V y, which increases abruptly thereafter.", It also is reflected in the evolution of which increases abruptly thereafter.
+ Clearly. the bar erowtl ceases at some point when the exceeds aud the bar cannot capture additional orbits aud be fed bv the augular momentum from these orbits.," Clearly, the bar growth ceases at some point when the exceeds and the bar cannot capture additional orbits and be fed by the angular momentum from these orbits."
+ This situation is simular to that analyzed im Paper I (Section 3.6). where the aneular momentum of the disk inside stavs about constant as long as reais within the disk.," This situation is similar to that analyzed in Paper I (Section 3.6), where the angular momentum of the disk inside stays about constant as long as remains within the disk."
+ Our present models can be directly compared to the Standard model of Paper I. We return o this point in Section I., Our present models can be directly compared to the Standard model of Paper I. We return to this point in Section 4.
+ The loss of the augular momentum.αν with xw the outer disk appears to be umeh more severe thanJag.," The loss of the angular momentum, with by the outer disk appears to be much more severe than."
+ But oue should remember that the outer disk. ic. outside the CR. accounts for less mass.," But one should remember that the outer disk, i.e., outside the CR, accounts for less mass."
+ Before the uckliug. behaves similarly toJq44.. but ciffereutlv at the later nues.," Before the buckling, behaves similarly to, but differently at the later times."
+ The trend here is that declines steeply with time for gas poor disks after the nckliue period., The trend here is that declines steeply with time for gas poor disks after the buckling period.
+ It stavs nearly constant with time for he eas rich models., It stays nearly constant with time for the gas rich models.
+ Models with different show a behavior cousistcut with our understanding., Models with different show a behavior consistent with our understanding.
+ There is a slight increase in the J trauster duriug dvuaimical aud secular stages of the bar evolution., There is a slight increase in the $J$ transfer during dynamical and secular stages of the bar evolution.
+ Much less difference is observed in the evolution of and i this case., Much less difference is observed in the evolution of and in this case.
+ (e.g..Antonucci1993).. (e.g..Tran2001.&Huang2002).. (e.g..Trumpetal.2009).. (seeHo2008.andreferencestherein).. (e.g..Nicastroet2009).," \citep*[e.g.,][]{1993ARA&A..31..473A}. \citep*[e.g.,][]{2001ApJ...554L..19T,2003ApJ...583..632T,2002ApJ...579..205G}, \citep*[e.g.,][]{2009ApJ...700...49T}. \citep*[see][for a review and references
+therein]{2008ARA&A..46..475H}. \citep*[e.g.,][]{2003ApJ...589L..13N,2003ApJ...590...86L,2006ApJ...648L.101E,2009ApJ...701L..91E}."
+ Laor(2003) (Kaspi2000).., \citet{2003ApJ...590...86L} \citep{2000ApJ...533..631K}.
+ disk (Emmeringetal.1992).., disk \citep{1992ApJ...385..460E}.
+ Nicastro(2000) assumed that the winds from the accretion disk are triggered by the thermal instability of radiation pressure dominated region of the disk (Shakura&Sunyaev1976).., \citet{2000ApJ...530L..65N} assumed that the winds from the accretion disk are triggered by the thermal instability of radiation pressure dominated region of the disk \citep{1976MNRAS.175..613S}.
+ The transition radius between the radiation pressure dominated and gas pressure dominated regions in the disk increases with the dimensionless mass accretion rate 71 (Shakura&Sunyaev1973).., The transition radius between the radiation pressure dominated and gas pressure dominated regions in the disk increases with the dimensionless mass accretion rate $\dot{m}$ \citep{1973A&A....24..337S}.
+ In this scenario. the transition radius becomes smaller than the marginal stable orbit of the black hole for low accretion rates (low luminosities). and the winds are switched off and no BLR ca be formed in LLAGNs (Nicastroetal.2003)..," In this scenario, the transition radius becomes smaller than the marginal stable orbit of the black hole for low accretion rates (low luminosities), and the winds are switched off and no BLR can be formed in LLAGNs \citep{2003ApJ...589L..13N}."
+ A correlatio between the width of BLR and the luminosity is expected in this model. which is consistent with the observations of AGN samples (Warneretal.2004:Xu&Cao2007)..," A correlation between the width of BLR and the luminosity is expected in this model, which is consistent with the observations of AGN samples \citep{2004ApJ...608..136W,2007ChJAA...7...63X}."
+ A alternative disk-wind scenario was suggested for the BLR and dust torus. in which both the BLR and torus disappear when the bolometric luminosity is low (Elitzur&Shlosma2006:Elitzur&Ho 2009)..," An alternative disk-wind scenario was suggested for the BLR and dust torus, in which both the BLR and torus disappear when the bolometric luminosity is low \citep{2006ApJ...648L.101E,2009ApJ...701L..91E}."
+ The outflow from the accretio disk being switehed off is à key ingredient in these scenarios when accretion rates are low. though the detailed physics of the outflow dynamics has not been included in these works.," The outflow from the accretion disk being switched off is a key ingredient in these scenarios when accretion rates are low, though the detailed physics of the outflow dynamics has not been included in these works."
+ Low mass accretion rate #7 may lead to the accretion flows to be advection-dominated (Narayan&Yi1994. 1995b)..," Low mass accretion rate $\dot{m}$ may lead to the accretion flows to be advection-dominated \citep{1994ApJ...428L..13N,1995ApJ...452..710N}."
+ Advection dominated accretionflows(ADAFs)aresuggested to be present in LLAGNs (seeNarayan2002.forareview therein).. which can successfully explain most observational features of LLAGNSs (e.g..Lasotaetal.1996; 2009)..," Advection dominated accretionflows(ADAFs)aresuggested to be present in LLAGNs \citep*[see][for a review and references
+therein]{2002luml.conf..405N}, , which can successfully explain most observational features of LLAGNs \citep*[e.g.,][]{1996ApJ...462..142L,1999ApJ...516..177G,1999ApJ...525L..89Q,2009RAA.....9..401X}. ."
+of secing on the mean ellective surface brightness. defined ὃν using Eq. (,"of seeing on the mean effective surface brightness, defined as: By using Eq. ("
+15) and the conservation of the flux. it immecdiatelv follows that In Fig.,"15) and the conservation of the flux, it immediately follows that In Fig."
+ 6 we show how the mean elective surface brightness changes as a function of rz /ENΗΝ for cillerent values of m and intrinsic. ellipticities., 6 we show how the mean effective surface brightness changes as a function of $r_{\rm e}^{\rm c}/$ FWHM for different values of $n$ and intrinsic ellipticities.
+ This figure clearly shows that galaxies allected by seeing have apparent mean surfaces brightnesses lower than their true values., This figure clearly shows that galaxies affected by seeing have apparent mean surfaces brightnesses lower than their true values.
+ Lower values of; produce greater effects on this quantity., Lower values of $\beta$ produce greater effects on this quantity.
+ Also. as the intrinsic elliptieity of the object increases the ellects of the seeing on the mean ellective surface brightness also increase.," Also, as the intrinsic ellipticity of the object increases the effects of the seeing on the mean effective surface brightness also increase."
+ ]t is possible to obtain the parameters of the Sérrsic profiles (secing-[ree. quantities) from the convolved. quantities., It is possible to obtain the parameters of the Sérrsic profiles (seeing-free quantities) from the convolved quantities.
+ We present an easy. based on the use of the plots of Figures 3. 4 and 5.," We present an easy based on the use of the plots of Figures 3, 4 and 5."
+ Phe steps which observer must follows are, The steps which observer must follows are:
+of a larger sample that was selected. ou the basis of optical morphological evidence for strong tidal interactions and has been well studied at optical. infrared. aud radio wavelengths.,"of a larger sample that was selected on the basis of optical morphological evidence for strong tidal interactions and has been well studied at optical, infrared, and radio wavelengths."
+ The goal of this study is not ouly to mcasure the molecular eas content of interacting galaxies. but more müportantlv to examine the relationship between eas content and star formation characteristics. as well as interaction streneth.," The goal of this study is not only to measure the molecular gas content of interacting galaxies, but more importantly to examine the relationship between gas content and star formation characteristics, as well as interaction strength."
+ Questions that this study will address ποιο whether or uot a rich molecular gas disk is a prerequisite for a starburst iu interacting svstenis or if enhanced molecular eas production i$ occunug: whether optically “dormant” pair members have molecular material present: how the molecular eas content. far-TR huuimositv aud colors. aud optical tudicators of star formation are related in strongly interacting galaxies: aud what the relationship is between molecular and atomic gas conteut.," Questions that this study will address include whether or not a rich molecular gas disk is a prerequisite for a starburst in interacting systems or if enhanced molecular gas production is occuring; whether optically “dormant” pair members have molecular material present; how the molecular gas content, far-IR luminosity and colors, and optical indicators of star formation are related in strongly interacting galaxies; and what the relationship is between molecular and atomic gas content."
+ By addressing these questions. we hope to achieve a better idea of the circumstances under which starbursts occur in interactions aud thus of the mechanisius which trigeer the burst.," By addressing these questions, we hope to achieve a better idea of the circumstances under which starbursts occur in interactions and thus of the mechanisms which trigger the burst."
+ The galaxies for this study have been selected frou the ]arge suuple of interacting svstenis compiled by Bushouse (1986)., The galaxies for this study have been selected from the large sample of interacting systems compiled by \cite{bus86}.
+.. Members of this sample were originally chosen ou the basis of optical morphology ouly. prefercutially selecting svstenis with features characteristic of strong interactions. such as tidal tails. loops. and bridges.," Members of this sample were originally chosen on the basis of optical morphology only, preferentially selecting systems with features characteristic of strong interactions, such as tidal tails, loops, and bridges."
+ Heuce it is not intentionally biased towards mfrared- or raclio-bright systems. nor galaxies containing active ealactic nuclei (ΔένS). although many svstenis m the pareut siuple do have these characteristics.," Hence it is not intentionally biased towards infrared- or radio-bright systems, nor galaxies containing active galactic nuclei (AGN's), although many systems in the parent sample do have these characteristics."
+ Another advantage in choosing xvsteuis from this sample is that au exteusive and well-studied database of optical spectra. Mages. near and far-IR photometry and imaging. aud 22] cin data exists (Bushouse 1986. L987: Dushouse.&Werner 1988: Bushouse&Werner 1990: Dushousc.Telesco.&Werner 1998)).," Another advantage in choosing systems from this sample is that an extensive and well-studied database of optical spectra, images, near- and far-IR photometry and imaging, and 21 cm data exists (Bushouse 1986, 1987; \cite{bus88}; \cite{bus90}; \cite{bus98}) )."
+" We have obtained ""CO (10) observations for 21 systenis frou the sample of Bushouse(1986).", We have obtained $^{12}$ CO (1–0) observations for 24 systems from the sample of \cite{bus86}.
+. Some of the observed systems were selected. because of their relatively lugh far-IR fluxes. which. because of the welbkuowu correlation between far-IR aud CO fluxes for ealaxics. sugeested that they would be casily detectable iu the CO chussion line.," Some of the observed systems were selected because of their relatively high far-IR fluxes, which, because of the well-known correlation between far-IR and CO fluxes for galaxies, suggested that they would be easily detectable in the CO emission line."
+ Other systems observed iu this study were chosen in an attempt to fully sample the parameter space of kuown optical. infrared. and radio characteristics of the class of interacting galaxies.," Other systems observed in this study were chosen in an attempt to fully sample the parameter space of known optical, infrared, and radio characteristics of the class of interacting galaxies."
+ For example. while iufrared-biieht systems often have optical indicators of lugh current star formation rates (SFR3s). some svstenis were chosen because they have optical indications of low SERs. regardless of their infrared flux levels.," For example, while infrared-bright systems often have optical indicators of high current star formation rates (SFR's), some systems were chosen because they have optical indications of low SFR's, regardless of their infrared flux levels."
+ Otlers were chosen because they have optical indicators of high current SERs. vet have low far-IR emission levels.," Others were chosen because they have optical indicators of high current SFR's, yet have low far-IR emission levels."
+ Many svstems were also chosen because of the relatively laree aneular separation of the two galaxies in the pair. allowing separate CO moeasuremoeuts for cach of the galaxies and subsequent analysis of the properties of the individual galaxies.," Many systems were also chosen because of the relatively large angular separation of the two galaxies in the pair, allowing separate CO measurements for each of the galaxies and subsequent analysis of the properties of the individual galaxies."
+ We also searched the literature for CO observations of systems contained m the Bushouse(1986) parent sample., We also searched the literature for CO observations of systems contained in the \cite{bus86} parent sample.
+ This search vielded a total of 19 observations of galaxics in 13 interacting svstenis., This search yielded a total of 19 observations of galaxies in 13 interacting systems.
+ Most of these data are for iufrared-bright systems. as previous studies mainly used IB-selected samples.," Most of these data are for infrared-bright systems, as previous studies mainly used IR-selected samples."
+ The combined sample of iteracting systems analyzed im this study is not complete in either a Hux-limited or vohuue-linited scuse: rather. the galaxies were selected to span a wide range of parameter space. including far-IR fux level aud optical indications of star ormation properties.," The combined sample of interacting systems analyzed in this study is not complete in either a flux-limited or volume-limited sense; rather, the galaxies were selected to span a wide range of parameter space, including far-IR flux level and optical indications of star formation properties."
+ Table 1 contains the list of interacting svsteis frou the sample of Bushouse(1986) for which CO observations ive been obtained., Table \ref{tbl_codata} contains the list of interacting systems from the sample of \cite{bus86} for which CO observations have been obtained.
+ It also contaius a comment on the rature of each system and the CO observations that lave con obtained., It also contains a comment on the nature of each system and the CO observations that have been obtained.
+" There are a total of 19 paired «πο», refered to as “complete pairs”. for which individual CÓ observations of the two galaxies in cach pair are available."," There are a total of 19 paired systems, refered to as “complete pairs”, for which individual CO observations of the two galaxies in each pair are available."
+ There are another 5 systems for which data for oulv oue of the two galaxies (vl of 27) in the pair has heen obtained., There are another 5 systems for which data for only one of the two galaxies (“1 of 2”) in the pair has been obtained.
+ Finally. there are 13 svstenis that have a sinele CO observation that encompasses the eutire pair.," Finally, there are 13 systems that have a single CO observation that encompasses the entire pair."
+ Nine of these are comipact svstenis in which the pair of galaxies fits within a suele CO aperture., Nine of these are compact systems in which the pair of galaxies fits within a single CO aperture.
+" The remaining four svstenisNGC σι, UGC 966. TOC 1509. and. UGC 8387απο eoncrally accepted to be systems that are in an advanced stage of increimg. where the two original galaxies are in the process of coalescing into a single object."," The remaining four systems---NGC 1614, UGC 966, UGC 4509, and UGC 8387—are generally accepted to be systems that are in an advanced stage of merging, where the two original galaxies are in the process of coalescing into a single object."
+ CO emissiou lias been detected in £1 of the 56 observations., CO emission has been detected in 41 of the 56 observations.
+ We have also used a sample of isolated spiral galaxies from the Five College Radio Astronomy Observatory (FCRAO) Extragalactic CO Survey. (Young et al., We have also used a sample of isolated spiral galaxies from the Five College Radio Astronomy Observatory (FCRAO) Extragalactic CO Survey (Young et al.
+ 1995) as a conrparison sanuple in our study of the interacting systems;, 1995) as a comparison sample in our study of the interacting systems.
+ We selected 80 galaxies from this survey. raugiug in inorphological type from SO to Sd. intentionally exchiding iregulars and known iuteracting svstenis.," We selected 80 galaxies from this survey, ranging in morphological type from S0 to Sd, intentionally excluding irregulars and known interacting systems."
+ The ealaxies were chosen simply on the basis of available aud CO observations., The galaxies were chosen simply on the basis of available and CO observations.
+ The 0ο (10) observations were obtained in December 1988 using the 12 10 NRAO (IIPBW = aat 115 Cz) ou Kitt Peak., The $^{12}$ CO (1–0) observations were obtained in December 1988 using the 12 m NRAO (HPBW = at 115 GHz) on Kitt Peak.
+ The telescope was equipped with dual polarization SIS receivers (Tuan110 KR)., The telescope was equipped with dual polarization SIS receivers $T_{ssb} \sim 110$ K).
+ Two 256 x 2 MIIz channel filter banks. one for cach polarization. provided a total velocity coverage of approximately 1360 wwith a resolution of ~5.6 pper channel.," Two 256 x 2 MHz channel filter banks, one for each polarization, provided a total velocity coverage of approximately 1360 with a resolution of $\sim$ 5.6 per channel."
+ Telescope pointing was uonitored bv observations of planets aud was estimated to be accurate to (xis)., Telescope pointing was monitored by observations of planets and was estimated to be accurate to (rms).
+ All of the observations were taken usine the nutatiug subreflector. which eave very fat baselines.," All of the observations were taken using the nutating subreflector, which gave very flat baselines."
+ Thespectra were calibrated using an ambicut temperature chopper wheel aud scaled by the telescope cficieucy on a spatially extended source. yas HALoo4 to give line temperatures T5.," Thespectra were calibrated using an ambient temperature chopper wheel and scaled by the telescope efficiency on a spatially extended source, $\eta_{fss}$ = $\eta_{Moon}$ to give line temperatures $T_{R}^{*}$."
+ After co-adding both polarizations. the CO spectra were Caussian smoothed to a resolution of 152O]ans i. eiviug au average riis noise of 3. Luks per chiuncl.," After co-adding both polarizations, the CO spectra were Gaussian smoothed to a resolution of 15–20 , giving an average rms noise of 3–4 mK per channel."
+ The CO spectra were obtained with the spectrometer, The CO spectra were obtained with the spectrometer
+(arius) of the interferometer are desigued to be tilted towards each other aloug the slit clirection such that several fringes are formed aloug each v channel.,(arms) of the interferometer are designed to be tilted towards each other along the slit direction such that several fringes are formed along each $\nu$ channel.
+" The intersection of a stellar absorption line and a WLC moves (from 2, to P in Fig. 1))", The intersection of a stellar absorption line and a WLC moves (from $P_o$ to $P$ in Fig. \ref{fig:DFDI_setup}) )
+ if there is a shift of an absorption liue due to a chauge of stellar RV., if there is a shift of an absorption line due to a change of stellar RV.
+" Consequently. a small chauge of © in the dispersion direction. AXo,.. is iuduced: where LV is defined. as phase-to-velocity scale (PV scale)."," Consequently, a small change of $\phi$ in the dispersion direction, $\Delta\phi_x$, is induced: where $\Gamma$ is defined as phase-to-velocity scale (PV scale)."
+" It is determiued by the GD of au interferometer. which becomes explicit if Equation (1)) aud (3)) are combiued: At resolutions typically adopted by the DEDI inethod (5.000<[tx20.000). stellar lii (line wicth~0.! LA) are not resolved anc a measurement of A®, is extremely cillicult."," It is determined by the GD of an interferometer, which becomes explicit if Equation \ref{eq:GD_def}) ) and \ref{eq:fring_phase_shift}) ) are combined: At resolutions typically adopted by the DFDI method $5,000\le R\le20,000$ ), stellar lines (line $\sim$ $\AA$ ) are not resolved and a measurement of $\Delta\phi_x$ is extremely difficult."
+" Instead. Nó. ghase shift aloig y direction cau be measured. which is ecual to Nó, i. the combs generated by interferomeer are parallel to each other."," Instead, $\Delta\phi_y$, phase shift along $y$ direction can be measured, which is equal to $\Delta\phi_x$ if the combs generated by an interferometer are parallel to each other."
+ This is a gooc approxinmationat very higli orders of interference., This is a good approximationat very high orders of interference.
+" The advantage of measuring Ao, instead of Ao, is seen [f‘om Fig. 1.."," The advantage of measuring $\Delta\phi_y$ instead of $\Delta\phi_x$ is seen from Fig. \ref{fig:DFDI_setup},"
+ in which t johysical shift i1 the v direction is amplified iu y direction. he amplificaion rate Is determined by he relative ane[n]le between the interferometer combs aud a stellar absorption line.," in which the physical shift in the $\nu$ direction is amplified in $y$ direction, the amplification rate is determined by the relative angle between the interferometer combs and a stellar absorption line."
+" Thereloο, AO, is ‘elatively easie ""to measure compared to ὦ aud it is unieasuος by fitting a well-sampled yeriodical ]ux sienal alot& the y direction in the DEDI method."," Therefore, $\Delta\phi_y$ is relatively easier to measure compared to $\phi_x$ and it is measured by fitting a well-sampled periodical flux signal along the $y$ direction in the DFDI method."
+ Coripared to coiventioual high-'esolutiou Eclele method. the number of freedom for the DEDI metiod in the filing process ls uuch less all stnall Dopjer phase shift can be relatively easier detected with a siuple functional orm. Le.. inisoldal fuictiou.," Compared to conventional high-resolution Echelle method, the number of freedom for the DFDI method in the fitting process is much less and small Doppler phase shift can be relatively easier detected with a simple functional form, i.e., a sinusoidal function."
+ However. we want to poiut out that while the DEDI metloc| provides a boost in instruuent Doppler sensitivitv. the Doppler sensitivity is not strongly depencleHu on e slijication Lee because flux slope decreases as atplificatiou rate increases. whicl Legales the 1 of pliase sope.," However, we want to point out that while the DFDI method provides a boost in instrument Doppler sensitivity, the Doppler sensitivity is not strongly dependent on the amplification rate because flux slope decreases as amplification rate increases, which negates the gain of phase slope."
+ The paper is organized as [ollows: in 82.. we present a new method of GD measu'emienut usine a DEDI Doppler instrument. which exploits WLCs generated by a fixed delay. interferometer.," The paper is organized as follows: in \ref{sec:OpdMeasurement}, we present a new method of GD measurement using a DFDI Doppler instrument, which exploits WLCs generated by a fixed delay interferometer."
+ 83. we preseut another method of GD calibration using au RV reference star.," In \ref{sec:OPD_solution}, , we present another method of GD calibration using an RV reference star."
+ Observation results alter implementing the newly measured GD are presented in δ1.., Observation results after implementing the newly measured GD are presented in \ref{sec:Implementation}.
+ In 35.. we suniarize and clisc the new results in this paper.," In \ref{sec:Conclusion}, we summarize and discuss the new results in this paper."
+ MARVELS (Multi-objeet Apache Point Observatory Racial Velocity Exoplanet Large-area Survey) is part of the Sloan Digital Sky Survey (SDSS) ILE (??)..," MARVELS (Multi-object Apache Point Observatory Radial Velocity Exoplanet Large-area Survey) is part of the Sloan Digital Sky Survey (SDSS) III \citep{Gunn2006,Eisenstein2011}. ."
+ The iustrumeut covers awavelength, The instrument covers awavelength
+distributionn. While in agreement with standard shock theory. this scenario also agrees with the idea that the luminosity of blazars is produced through internal shocks. which naturally lead to shocks lasting for a finite time.,"n. While in agreement with standard shock theory, this scenario also agrees with the idea that the luminosity of blazars is produced through internal shocks, which naturally lead to shocks lasting for a finite time."
+sinnce October 2002.,nce October 2002.
+circunustellar disks.,circumstellar disks.
+ There are some examples of wisaligned svstems inferred frou jets enmianatiug with different position angles (Caedel Reith 1993: BReiputh ct al., There are some examples of misaligned systems inferred from jets emanating with different position angles (Gredel Reipurth 1993; Reipurth et al.
+ 1993)., 1993).
+ Wowever. this iuethod is wohlematic because two well-aligued outflows would be wisidentified asa sinele outflow.," However, this method is problematic because two well-aligned outflows would be misidentified as a single outflow."
+ The best way to test the disk aliguunent may be the direct inaeiug of both disks., The best way to test the disk alignment may be the direct imaging of both disks.
+ If III] 24 MAIS is a binary syste as discussed in he previous section. it is one of the rare examples of xotobinarv svsteni nuaged with au augular resolution Heh enough to investigate the deeree of alieumenut.," If HH 24 MMS is a binary system as discussed in the previous section, it is one of the rare examples of protobinary system imaged with an angular resolution high enough to investigate the degree of alignment."
+ Each source may consist of a compact accretion disk sirrounded by a pseudocdisk. aud the axis of the accretion disk is verv likely to be parallel to the axis of the pseudodisk around it.," Each source may consist of a compact accretion disk surrounded by a pseudodisk, and the axis of the accretion disk is very likely to be parallel to the axis of the pseudodisk around it."
+ Then the axis direction of the accretion disk can be inferred from the minor axis of the elliptical ft to the 6.9 mun source., Then the axis direction of the accretion disk can be inferred from the minor axis of the elliptical fit to the 6.9 mm source.
+ The major axes of the two components of the ΠΠ 21 AIMS system are not aligned at all., The major axes of the two components of the HH 24 MMS system are not aligned at all.
+ The ciffereuce iu the position angle between the binary compoucuts. from the Gaussian fit (Table 1). is ~ ((or ~ uf they are retrograde).," The difference in the position angle between the binary components, from the Gaussian fit (Table 1), is $\sim$ (or $\sim$ if they are retrograde)."
+ HII 21 MOMS imakes a stark contrast to the other known example of resolved binary disks. L1551 IRS 5.," HH 24 MMS makes a stark contrast to the other known example of resolved binary disks, L1551 IRS 5."
+ In the L1551 IRS 5 system. the two disks are very well aligned (Rodríguez et al.," In the L1551 IRS 5 system, the two disks are very well aligned guez et al."
+ 1998)., 1998).
+ The differcuce between the two systems is probably owing to the binary separation: 360 AU in ΠΠ 21 AIMS and 15 AU in L1551 TRS 5., The difference between the two systems is probably owing to the binary separation: 360 AU in HH 24 MMS and 45 AU in L1551 IRS 5.
+ Iu the classification scheme of multiple vouus stellar objects suggested by Looney et al. (, In the classification scheme of multiple young stellar objects suggested by Looney et al. (
+2000). III] 21 AIMS belongs to the conumaion-envelope type while L1551 IRS 5 belongs to the commmou-disk type.,"2000), HH 24 MMS belongs to the common-envelope type while L1551 IRS 5 belongs to the common-disk type."
+ We speculate that the rotational fracinentation is iÀuportaut in close binaries aud the turbulent fragmentation is important iu wide binaries. though more examples are certainly needed to draw a statistically significant conclusion.," We speculate that the rotational fragmentation is important in close binaries and the turbulent fragmentation is important in wide binaries, though more examples are certainly needed to draw a statistically significant conclusion."
+ Other examples may include TRAS 162932122 and NGC 1333 IRAS LA (Looney et al., Other examples may include IRAS 16293–2422 and NGC 1333 IRAS 4A (Looney et al.
+ 2000). but their structures are complicated and their disk orientations are less obvious.," 2000), but their structures are complicated and their disk orientations are less obvious."
+ Though studies through direct imagine is difficult. indirect evidences show that misalieued svstenis are not recessarily unusual.," Though studies through direct imaging is difficult, indirect evidences show that misaligned systems are not necessarily unusual."
+ Wieh-resolution ucar-IR polarimetry revealed plenty of examples of misaligned binaries in the Class II phase (Alouin et al., High-resolution near-IR polarimetry revealed plenty of examples of misaligned binaries in the Class II phase (Monin et al.
+ 2007)., 2007).
+ Disks in T Tauri üuaries are not perfectly coplanar (Jensen et al., Disks in T Tauri binaries are not perfectly coplanar (Jensen et al.
+ 2001). and the relative alizumoent of the disks mav decrease after ormation (Bate et al.," 2004), and the relative alignment of the disks may decrease after formation (Bate et al."
+ 2000)., 2000).
+ Therefore. IIT 21 MIMS can ο considered as aprecursor of such misaligned T Tauri systems.," Therefore, HH 24 MMS can be considered as a of such misaligned T Tauri systems."
+ As TIT 241 AIMS is a rare example of a wide nary that can be unaged well enough to study the disk alieument. this svstem eives us a chance to exanune the initial couditious of binary evolution.," As HH 24 MMS is a rare example of a wide binary that can be imaged well enough to study the disk alignment, this system gives us a chance to examine the initial conditions of binary evolution."
+ We thank WAT. Iii for helpful discussions aud encouragement., We thank K.-T. Kim for helpful discussions and encouragement.
+ M. Ix. is erateful to J. Diegiue for helpful discussions., M. K. is grateful to J. Bieging for helpful discussions.
+ This work was partially supported by the LRG program of KASL, This work was partially supported by the LRG program of KASI.
+ M. K. acknowledges support from the Iorea Research Foundation eraut KRE-2007-612-C'O0050. funded by the Korean Covermment (MOEIIRD)., M. K. acknowledges support from the Korea Research Foundation grant KRF-2007-612-C00050 funded by the Korean Government (MOEHRD).
+"indicated a very high-redshift nature, with basically zero dust absorption (Cucchiara et al 2009b, Olivares et al 2009).","indicated a very high-redshift nature, with basically zero dust absorption (Cucchiara et al 2009b, Olivares et al 2009)."
+" Our group used the Italian 3.6m Telescopio Nazionale Galileo, in the island of La Palma, to obtain a low-resolution spectrum using the Amici prism with the spectrograph NICS (Oliva 2003), and measured its redshift to be z=δ.Ι03 (Thoene et al 2009, Fernandez-Soto et al 2009, Salvaterra et al 2009)."," Our group used the Italian 3.6m Telescopio Nazionale Galileo, in the island of La Palma, to obtain a low-resolution spectrum using the Amici prism with the spectrograph NICS (Oliva 2003), and measured its redshift to be $z=8.1^{+0.1}_{-0.3}$ (Thoene et al 2009, Fernandez-Soto et al 2009, Salvaterra et al 2009)."
+" A compatible result (z= 8.23*006) was reached independently by Tanvir et al (2009b, 2009c) using two sets of higher-quality data obtained with the VLT in Chile."," A compatible result $z=8.23^{+0.06}_{-0.07}$ ) was reached independently by Tanvir et al (2009b, 2009c) using two sets of higher-quality data obtained with the VLT in Chile."
+" The Amici spectrum covers in a single exposure the wavelength range 0.8—2.5um with very low resolution 50) but very high efficiency, and thus became the ideal choice for this kind of analysis."," The Amici spectrum covers in a single exposure the wavelength range $0.8-2.5 \mu\rm{m}$ with very low resolution $R \approx 50$ ) but very high efficiency, and thus became the ideal choice for this kind of analysis."
+ We obtained a total of 128 minutes of on-target exposure time., We obtained a total of 128 minutes of on-target exposure time.
+" The exposures were dithered following the usual NIR technique, and combined into a single two-dimensional frame, which is showed in Figure 1."," The exposures were dithered following the usual NIR technique, and combined into a single two-dimensional frame, which is showed in Figure 1."
+" The slit was positioned with the help of a nearby star, whose extracted spectrum will be one of the keys in our analysis."," The slit was positioned with the help of a nearby star, whose extracted spectrum will be one of the keys in our analysis."
+ The position of the star along the slit (measured at the reference position X=600 in the CCD frame) is Y=765., The position of the star along the slit (measured at the reference position X=600 in the CCD frame) is Y=765.
+" The angular distance between the reference star and the afterglow was ~30 arcseconds, which corresponds to 120 pixels along the slit."," The angular distance between the reference star and the afterglow was $\approx 30$ arcseconds, which corresponds to 120 pixels along the slit."
+" In order to avoid possible issues caused by misalignments or the effect of distortions in the focal plane, we use this distance only as a reference, and perform a careful recentering, as described in the next Section."," In order to avoid possible issues caused by misalignments or the effect of distortions in the focal plane, we use this distance only as a reference, and perform a careful recentering, as described in the next Section."
+" Our aim will be to reproduce as perfectly as possible the spectrum of the afterglow of GRB090423, and to reconstruct the spectral equivalent of the wavelength-dependent as generated when the light passes through the atmosphere, telescope and instrument optics, and reaches the detector."," Our aim will be to reproduce as perfectly as possible the spectrum of the afterglow of GRB090423, and to reconstruct the spectral equivalent of the wavelength-dependent as generated when the light passes through the atmosphere, telescope and instrument optics, and reaches the detector."
+ We present in this Section the different steps that we perform to reach this objective., We present in this Section the different steps that we perform to reach this objective.
+" We create a library of model spectra, where the basic input parameters are three: the redshift z, the slope α in the power- spectral model f,«Α”, and the total neutral Hydrogen column density in the host interstellar medium N(HD, that produces a strong Damped Lyman Alpha (DLA) profile at the host redshift."," We create a library of model spectra, where the basic input parameters are three: the redshift $z$ , the slope $\alpha$ in the power-law spectral model $f_\nu \propto \lambda ^{\alpha}$, and the total neutral Hydrogen column density in the host interstellar medium N(HI), that produces a strong Damped Lyman Alpha (DLA) profile at the host redshift."
+" It is important to include this profile in the analysis, because adense (N(HI)= 10?'cm""?) DLA profile"," It is important to include this profile in the analysis, because adense $\rm{N(HI)} \gtrsim 10^{21} \rm{cm}^{-2}$ ) DLA profile"
+two earlier maps Clavlor et al.,two earlier maps (Taylor et al.
+ 1986) that were separated by only 75 clays., 1986) that were separated by only 75 days.
+ The upper limit of the surface magnetic field of a typical white dwarl is around LOG (Bond and Chanmugam. 1982).," The upper limit of the surface magnetic field of a typical white dwarf is around $10$ G (Bond and Chanmugam, 1982)."
+ Assuming a dipole geometry. this will fall olf as Bywpr ὃν ," Assuming a dipole geometry, this will fall off as $B_{\rm WD} \propto r^{-3}$ ."
+Por a stellar radius of O.OLR. the upper limit on the magnetic field at SOOAU is SO µας which is far too low to eive ellicient svnchrotron emission.," For a stellar radius of $_\odot$ the upper limit on the magnetic field at 500AU is $80~\mu$ G, which is far too low to give efficient synchrotron emission."
+ The speed. of the unshockec preexisting cool wind is likely to be much less than the jet speed. as a typical cool eiant wind has a velocity of the order of 40kkm s! (Vogel et al.," The speed of the unshocked pre–existing cool wind is likely to be much less than the jet speed, as a typical cool giant wind has a velocity of the order of km $^{-1}$ (Vogel et al."
+ 1994)., 1994).
+ Phe nonthermal emission is then caused bv particle acceleration. at. the resulting bow shock. as the local magnetic field is compressed. ancl enhanced: and electrons are accelerated. to velocities close to c.," The non–thermal emission is then caused by particle acceleration at the resulting bow shock, as the local magnetic field is compressed and enhanced and electrons are accelerated to velocities close to $c$."
+" 1n. such an environment. the local azimuthal field. //,., is defined by Equation 42 of Bode Ixahn (1985) the density of the shock. the temperature of the giant wind. (assumed to be constant at ~ LOW)m and the mean particle. mass which Da.is taken to be the proton mass."," In such an environment, the local azimuthal field, $H_\phi$ is defined by Equation 42 of Bode Kahn (1985) the density of the shock, the temperature of the giant wind (assumed to be constant at $\sim 10^4$ K) and the mean particle mass which is taken to be the proton mass."
+ The shock is assumed to be strong and totally ionised. so that the shock density is equal to + times the density of the unshocked. wind at this distance from the star.," The shock is assumed to be strong and totally ionised, so that the shock density is equal to 4 times the density of the unshocked wind at this distance from the star."
+ Ht is assumed that the disruptions caused to the giant wind by the hot component prior to mass ejection are negligible at this distance so that a uniform. spherical racial wind exists.," It is assumed that the disruptions caused to the giant wind by the hot component prior to mass ejection are negligible at this distance so that a uniform, spherical radial wind exists."
+ The denisty of the wind is proportional to the mass loss rate ofthe giant. AL=L9-10PM.vr+ (Skopal. 1996).," The denisty of the wind is proportional to the mass loss rate ofthe giant, $\dot{M} = 1.9\times 10^{-6} {\rm M}_\odot~{\rm yr}^{-1}$ (Skopal, 1996)."
+" Η the unshockecl wind speed. eu is taken to be the escape velocity of the red. giant. assuming Maii=0.93M. and Roan=10123420. (GSkopal 1907) then e,=53.8kins+ and IH,224 mG. Another magnetic field estimate can be derived [rom the technique used in studsing high energy. extragalactie sources."," If the unshocked wind speed, $v_w$ is taken to be the escape velocity of the red giant, assuming $M_{\rm giant} = 0.93M_\odot$ and $R_{\rm giant} = 123R_\odot$ (Skopal 1997) then $v_w = 53.8~{\rm
+km~s^{-1}}$ and $H_\phi\simeq 2.4$ mG. Another magnetic field estimate can be derived from the technique used in studying high energy extragalactic sources."
+ The total energy of a radio source is split between that of the electrons. the protons and the magnetic field.," The total energy of a radio source is split between that of the electrons, the protons and the magnetic field."
+ At equipartition. the magnetic energy. is approximately equal to the particle energy (Pachoezvk. 1970). and. the ratio of energy carried by protons to that carried by electrons is A.," At equipartition, the magnetic energy is approximately equal to the particle energy (Pachoczyk, 1970), and the ratio of energy carried by protons to that carried by electrons is $K$."
+ bor a radio source of known angular dimensions aud depth. assuming an ellipsoidal geometry. the magnetic field. under equipartition. //g. can be found using the method described by Miley. (1980).," For a radio source of known angular dimensions and depth, assuming an ellipsoidal geometry, the magnetic field under equipartition, $H_E$, can be found using the method described by Miley (1980)."
+ To make use of this model. one needs to determine several parameters that deseribe the region: filling factor. η]: the angle between the magnetic field and the line of sight. x: the Lux density of the emitting region. ορ.) at frequeney v (Gllz) and the maximum and. minimum. frequencies of radio emission (yas and Min repectivelv. in 11) from the ποιος spectrum.," To make use of this model, one needs to determine several parameters that describe the region: filling factor, $\eta$; the angle between the magnetic field and the line of sight, $\chi$; the flux density of the emitting region, $S_\nu(Jy)$, at frequency $\nu$ (GHz) and the maximum and minimum frequencies of radio emission $\nu_{\rm max}$ and $\nu_{\rm min}$ repectively, in GHz) from the source spectrum."
+ The filling factor would be more useful in extragalactic work when it would be expected that there would be unresolved. blobs of emitting material., The filling factor would be more useful in extragalactic work when it would be expected that there would be unresolved blobs of emitting material.
+ Since the measure sizes ob the knots in the extended emission are of a similar angular size to the beam. y=1.," Since the measured sizes of the knots in the extended emission are of a similar angular size to the beam, $\eta = 1$."
+ The angle x is importan due to the fact that the radiation observed. is beamed from electrons moving towards the observer., The angle $\chi$ is important due to the fact that the radiation observed is beamed from electrons moving towards the observer.
+ The fact that eclipses are seen in the system. means that the orbital plane of the stars is observed. edgeon., The fact that eclipses are seen in the system means that the orbital plane of the stars is observed edge–on.
+ Bipolar jets are believe to be aligned. with the rotation axes of their parent. star »erpendicular to the orbital plane) so the jets of CLL Cveni are believed to lic almost in the plane of the sky. (Crocker et al.," Bipolar jets are believed to be aligned with the rotation axes of their parent star (perpendicular to the orbital plane) so the jets of CH Cygni are believed to lie almost in the plane of the sky (Crocker et al.,"
+ in preparation). ó=907.," in preparation), $\phi \simeq
+90^\circ$."
+ The cutoll frequencies of the radio emission. Minin abel [uas are taken to be 01016111 and. 100€ilIz respectively (Miley. 1980).," The cutoff frequencies of the radio emission, $\nu_{\rm
+min}$ and $\nu_{\rm max}$ are taken to be 0.01GHz and 100GHz respectively (Miley, 1980)."
+ Phe value of AN. is the main unknown in this equation as the exact details of the mechanism leading to svnchrotron emission are not known., The value of $K$ is the main unknown in this equation as the exact details of the mechanism leading to synchrotron emission are not known.
+ A value of fy=LOO would be appropriate or the electrons being produced. following collisions in a circumstellar mecium (Pacholezyvk 1970). but a value of A= lis more consistent with the minimum energy. requirement (Miley 1950).," A value of $K=100$ would be appropriate for the electrons being produced following collisions in a circumstellar medium (Pacholczyk 1970), but a value of $K=1$ is more consistent with the minimum energy requirement (Miley 1980)."
+ However. since. Lf only depends upon Ae. wing ἐν LOO times larger would only increase fp by a actor of 3.," However, since $H$ only depends upon $K^\frac{2}{7}$, having $K$ 100 times larger would only increase $H_E$ by a factor of 3."
+ 1n order to determine the size of the jet bullets. two dimensional gaussians were fitted to the emitting regions in he NW and SE jets.," In order to determine the size of the jet bullets, two dimensional gaussians were fitted to the emitting regions in the NW and SE jets."
+ The depth was then found by assuming hat the regions exhibited. evlindrical svmmetry along the axis of ejection., The depth was then found by assuming that the regions exhibited cylindrical symmetry along the axis of ejection.
+" For the NW jet. 6,=0.39 aresee. #2=0.34 aresec and d=91 AU."," For the NW jet, $\theta_1 = 0.39$ arcsec, $\theta_2 = 0.34$ arcsec and $d = 91$ AU."
+" The SIS jet had dimensions of 6,=3T arcsec. 62=0.35 arcsec and d=94 AU."," The SE jet had dimensions of $\theta_1 = 0.37$ arcsec, $\theta_2 =
+0.35$ arcsec and $d = 94$ AU."
+ Using the mean value of the spectral index in the jets of à=0.23d: 0.03. for the NW jet and in the SE.," Using the mean value of the spectral index in the jets of $\alpha = -0.23 \pm 0.03$ , for the NW jet and in the SE."
+ These values are comparable to that found above from the physies of the shock. so there is sullicient magnetic ficld in the enhanced regions for svnchrotron radiation.," These values are comparable to that found above from the physics of the shock, so there is sufficient magnetic field in the enhanced regions for synchrotron radiation."
+Given realistic data. these models. will be inhomogeneous.,"Given realistic data, these models will be inhomogeneous."
+ Lndeed this is the reason that non-zero evolution functions have been introduced: (otherwise. observations are incompatible with a FLAW universe).," Indeed this is the reason that non-zero evolution functions have been introduced (otherwise, observations are incompatible with a FLRW universe)."
+ Subject to the conditions of appendix D.. for any given isotropic observations ἐς). ης). and any given LPB model. source evolution functions L(z) mls) can be found that make the LTB observational relations fit these observations.," Subject to the conditions of appendix \ref{ap2}, for any given isotropic observations $\ell(z)$ $n(z)$, and any given LTB model, source evolution functions $\hat{L}(z)$ $\hat{m}(z)$ can be found that make the LTB observational relations fit these observations."
+ We adapt the above algorithmic procedure to prove this., We adapt the above algorithmic procedure to prove this.
+ Again we assert that if the given observations and LTD model are ‘reasonable’. then the derived evolution Functions will be ‘reasonable’.," Again we assert that if the given observations and LTB model are `reasonable', then the derived evolution functions will be `reasonable'."
+ The idea is that we can vary the LTB moclel to which we fit the observations to some extent. but still keep the required: source evolution functions within a ‘realistic’ range.," The idea is that we can vary the LTB model to which we fit the observations to some extent, but still keep the required source evolution functions within a `realistic' range."
+ Source evolution functions can be found that make the dust. ELI observational relations fit any observations., Source evolution functions can be found that make the dust FLRW observational relations fit any observations.
+ An obvious consequence of (£5., An obvious consequence of (C).
+ for a parabolic case. hat /?(z) and plz) may not be single valued. and that the HOcand + plots can loop.," $\Box$ for a parabolic case, that $\hat{R}(z)$ and $\hat{\rho}(z)$ may not be single valued, and that the $\hat{R}-z$ and $\hat{\rho}$ $z$ plots can loop."
+ However. in compiling the real observational data. we merely add all the galaxies we see at a particular redshift. to get a number count.," However, in compiling the real observational data, we merely add all the galaxies we see at a particular redshift, to get a number count."
+ Similarly. we merely take an average over the Iuminosities observed at a articular redshift. ascribing the variation to natural scatter in intrinsic. properties ancl observational error. rather than o a multiply valued function.," Similarly, we merely take an average over the luminosities observed at a particular redshift, ascribing the variation to natural scatter in intrinsic properties and observational error, rather than to a multiply valued function."
+ Thus we make /(2) and p(s) single valued by construction., Thus we make $\hat{R}(z)$ and $\hat{\rho}(z)$ single valued by construction.
+ So the data functions we are rving to fit may not lead to such a good model., So the data functions we are trying to fit may not lead to such a good model.
+ In other words. assuming we succeed in constructing a well behaved LTD model from the data. it may not be the LTB model that yest represents the real universe.," In other words, assuming we succeed in constructing a well behaved LTB model from the data, it may not be the LTB model that best represents the real universe."
+ I seems unlikely though not entirely impossible that there will be a reliable way of de-convolving the superposed parts of these observational data curves. or even of discerning whether loops are present.," It seems unlikely — though not entirely impossible — that there will be a reliable way of de-convolving the superposed parts of these observational data curves, or even of discerning whether loops are present."
+ ]t is hard to predict how Likely this scenario is., It is hard to predict how likely this scenario is.
+ We have shown that a LED model (a Lemaittre-Tolman-3ondi spherically symmetric dust cosmologv) can be found to fit any given. set of observations of source counts n(z) and Luminosity/area distance Ri2). averaged over all angles. and any evolution functions for source Iunminosity L(z) and mass per source ος).," We have shown that a LTB model (a tre-Tolman-Bondi spherically symmetric dust cosmology) can be found to fit any given set of observations of source counts $n(z)$ and luminosity/area distance $\hat{R}(z)$, averaged over all angles, and any evolution functions for source luminosity $\hat{L}(z)$ and mass per source $\hat{m}(z)$."
+ In other words. even if we accept isotropy. then demonstrating homogeneity rather than assuming it must hold. because of the Copernican principle requires more than these observations.," In other words, even if we accept isotropy, then demonstrating homogeneity — rather than assuming it must hold because of the Copernican principle — requires more than these observations."
+ Conversely. our result can be used to determine the degree of inhomogencity from the observations and ogiven source evolution functions.," Conversely, our result can be used to determine the degree of inhomogeneity from the observations and given source evolution functions."
+ If the demonstration. of homogencity depends on knowing the source evolution. and validation. of source evolution theories depends on knowing the cosmological model is homogeneous. then neither is proved.," If the demonstration of homogeneity depends on knowing the source evolution, and validation of source evolution theories depends on knowing the cosmological model is homogeneous, then neither is proved."
+ Thus we need methods of validating source evolution. models that don't depend on assumptions of homogencity to establish the age at any given z., Thus we need methods of validating source evolution models that don't depend on assumptions of homogeneity to establish the age at any given $z$.
+ Similarly. deep cosmological distance measures that don't depend. on luminosity and are not influencecl by source evolution would help pin down the cosmological model better., Similarly deep cosmological distance measures that don't depend on luminosity and are not influenced by source evolution would help pin down the cosmological model better.
+ There are various promising developments. in particular: (a) distance measurement by Supernovae:," There are various promising developments, in particular: (a) distance measurement by Supernovae;"
+Dshock diffuse below 5.,$\rho_{shock}$ $\rho_{diffuse}$ below 5.
+" As expected, the compression factor is always greater than 1."," As expected, the compression factor is always greater than 1."
+" Sub-distributions, in green and red, gives the compression factor respectively for the green and red dots obtained in Figure17d."," Sub-distributions, in green and red, gives the compression factor respectively for the green and red dots obtained in Figure."
+". The green histogram, formed of points suggesting possible shock excitation, is very similar statistically to the black distribution with a mean of 7.5."," The green histogram, formed of points suggesting possible shock excitation, is very similar statistically to the black distribution with a mean of 7.5."
+" The red histogram, this time formed of points with a high probability of shock excitation, is slightly different, peaking at a greater psnock/paiffuse Value."," The red histogram, this time formed of points with a high probability of shock excitation, is slightly different, peaking at a greater $\rho_{shock}$ $\rho_{diffuse}$ value."
+" Its mean is estimated slightly above 10 i.e., those points (red in Figure 17d)) where, we believe, shocks may certainly dominate seem to be characterized by greater compression effects with respect to points (green in Figure 17d)) where the presence of shocks appears less obvious."," Its mean is estimated slightly above 10 i.e., those points (red in Figure ) where, we believe, shocks may certainly dominate seem to be characterized by greater compression effects with respect to points (green in Figure ) where the presence of shocks appears less obvious."
+" For abiadatic jump conditions, Rankine-Hugoniot relations indicate a compression factor of precisely 4 between pre- and post-shocked gas (Hennyetal.1999,Chapter12).."," For abiadatic jump conditions, Rankine-Hugoniot relations indicate a compression factor of precisely 4 between pre- and post-shocked gas \citep[Chapter 12]{Hen1999}."
+ The relatively large widths of the distributions displayed in Figure 18 indicate that isothermal shocks may dominate., The relatively large widths of the distributions displayed in Figure 18 indicate that isothermal shocks may dominate.
+" Also, the fact that forbidden lines such as [N textscii]and[S textscii] are detected favors radiative cooling and therefore non-adiabatic effects [although these lines most likely originate from ionized material trapped in the cooling, shocked outer shell and therefore may not be related to the wind-blown bubble's current state of expansion i.e., energy- or momentum-driven]."," Also, the fact that forbidden lines such as $[$ $]$ and $[$ $]$ are detected favors radiative cooling and therefore non-adiabatic effects [although these lines most likely originate from ionized material trapped in the cooling, shocked outer shell and therefore may not be related to the wind-blown bubble's current state of expansion i.e., energy- or momentum-driven]."
+" Theoretically, for an isothermal shock (Dyson&Williams1980,Chapter6), the compression factor corresponds to the square of the shock's Mach number, Mo (see the abscissa in Figure 18)."," Theoretically, for an isothermal shock \citep[Chapter 6]{Dys1980}, the compression factor corresponds to the square of the shock's Mach number, $M_{0}$ (see the abscissa in Figure 18)."
+" Therefore, for a speed of sound of 10 km ! in the H* medium (in agreement with our choice of 74400 K for the electron temperature in 11805; see § 4.2.3), shocks with typical velocities between 15 and 30 km s! (i.e., Mj ==22—10) may have led to shock excitation near Melotte 15."," Therefore, for a speed of sound of 10 km $^{-1}$ in the $^{+}$ medium (in agreement with our choice of 400 K for the electron temperature in 1805; see $\S$ 4.2.3), shocks with typical velocities between 15 and 30 km $^{-1}$ (i.e., $M_{0}^2$ $-$ 10) may have led to shock excitation near Melotte 15."
+ This is in agreement with typical expansion velocities found for wind-blown bubbles (see § 1)., This is in agreement with typical expansion velocities found for wind-blown bubbles (see $\S$ 1).
+" Shock velocities up to 50 km ! (i.e., Mg ~ 225) are also found although such dynamics is certainly peculiar."," Shock velocities up to 50 km $^{-1}$ (i.e., $M_{0}^2$ $\sim$ 25) are also found although such dynamics is certainly peculiar."
+" Note, however, that the compression measured here could be an upper limit since photoeroded gas located in the molecular cloud's envelope, rather than the diffuse foreground/background component, may act as pre-shocked material."," Note, however, that the compression measured here could be an upper limit since photoeroded gas located in the molecular cloud's envelope, rather than the diffuse foreground/background component, may act as pre-shocked material."
+" The cloud's envelope is expected to be, at least, a few tens of particles per cm? denser than typical values found for paisfuse in this work (Rathborneetal. 2009)."," The cloud's envelope is expected to be, at least, a few tens of particles per $^{3}$ denser than typical values found for $\rho_{diffuse}$ in this work \citep{Rat2009}."
+". Assuming a coplanar geometry, roughly 3—4 pc separate the most massive star of the Melotte 15 cluster and the central, ionized structure."," Assuming a coplanar geometry, roughly $-$ 4 pc separate the most massive star of the Melotte 15 cluster and the central, ionized structure."
+" Using the standard model for expanding wind-blown bubbles (Weaveretal. 1977),, the distance reached by a given bubble's post-shocked shell can be estimated from the mechanical luminosity of the winds (Ly 15Mwv%), the density of the pre-shocked medium (paif fuse) and the total time during which winds have been blown by the central star (t)."," Using the standard model for expanding wind-blown bubbles \citep{Wea1977}, , the distance reached by a given bubble's post-shocked shell can be estimated from the mechanical luminosity of the winds $L_{w}$ $\frac{1}{2}$$\dot{M}_{w}$$v_{\infty}^2$ ), the density of the pre-shocked medium $\rho_{diffuse}$ ) and the total time during which winds have been blown by the central star $t_{w}$ )."
+" Given t,, 22.5 Myr, the age of the Melotte 15 cluster (see § 2), we expect that two massive giant and supergiant O4 stars have left justvery recently the main-sequence branch."," Given $t_{w}$ 2.5 Myr, the age of the Melotte 15 cluster (see $\S$ 2), we expect that two massive giant and supergiant O4 stars have left justvery recently the main-sequence branch."
+" Therefore, for the"," Therefore, for the"
+a Group 5 star J5.,a Group 5 star –.
+6).. No star in the 2dl sample has the PCCAGLEL band stronger than 33124. but a number of stars achieve j-inclices that are just as high.," No star in the 2dF sample has the $^{13}{\rmn C}^{13}{\rmn C}~\lambda6144$ band stronger than 3124, but a number of stars achieve j-indices that are just as high."
+ This star is Star6-6 (Crabtreeetal. quoted as class)., This star is Star6-6 \cite{crabtree76} quoted as class.
+. The present work has H1814 as a Group 1 star4.8)... be. à star with weak carbon.," The present work has 1814 as a Group 1 star, i.e., a star with weak carbon."
+ The dillercnee could lie in. the statement by Crabtree et al., The difference could lie in the statement by Crabtree et al.
+ (1976). that the (0.00) and 11) CAFC bands are moderately strong. but the (0.22) HCC band and others in the Av |2 series are weak.," \shortcite{crabtree76} that the 0) and 1) $^{12}{\rmn C}^{12}{\rmn C}$ bands are moderately strong, but the 2) $^{12}{\rmn C}^{12}{\rmn C}$ band and others in the $\Delta\nu$ +2 series are weak."
+ The star is also 778 and its classification by Richer et al., The star is also 78 and its classification by Richer et al.
+ (1979) is the same as that by Crabtree et al. (1976)., \shortcite{richer79} is the same as that by Crabtree et al. \shortcite{crabtree76}.
+. It is also 1185 in the list of suspected. LMC CLIE stars (Llartwick&Cowley1988)., It is also 185 in the list of suspected LMC CH stars \cite{hartwick88}.
+. Phere is no evidence of significant photometric variation as seen in the infrared photometry [rom2-\LASS.. DENIS and Feast Whitelock (1992).. and the £ photometry from DIENIS and IKDMINOL.," There is no evidence of significant photometric variation as seen in the infrared photometry from, DENIS and Feast Whitelock \shortcite{feast92}, and the $I$ photometry from DENIS and KDMK01."
+ According to Richer et al. (1979)...," According to Richer et al. \shortcite{richer79}, ,"
+ this, this
+Z —0.02imodel,$Z$ =0.02 model.
+ The finalabundancesineachpulsecanbeidenti fiedasthosecorrespondinglolhepulsen umber markinther—avis., The final abundances in each pulse can be identified as those corresponding to the pulse number tick-mark in the x-axis.
+AMIhebeginningofathermalpulse. whiletheconvecliveinstabililyisingestingl hell burningashes Hsproducedandilsabundancesharplyinereases.," At the beginning of a thermal pulse, while the convective instability is ingesting the H-burning ashes, is produced and its abundance sharply increases."
+MIhesamelimelheabundancco[!? ddecreasesbecauseo[hediluliono{theinlershellmaterialiwith ll—burningasheswheretheabundancea {OF jissolar., At the same time the abundance of decreases because of the dilution of the intershell material with H-burning ashes where the abundance of is solar.
+"Subsequen|ly, iislransformedintro? EF."," Subsequently, is transformed into."
+" Inthermalpulses[ollowedbyT DU inourmodel. i.e. from (he lO"" thermal pulse onward. almost all iis destroved."," In thermal pulses followed by TDU in our model, i.e. from the $^{\rm th}$ thermal pulse onward, almost all is destroyed."
+" The maxinnun temperature at the base of the 10"" thermal pulse is =2.52x10% IX and iis reduced to about 1/10"" of its initial abundance in this pulse.", The maximum temperature at the base of the $^{\rm th}$ thermal pulse is $= 2.52 \times 10^8$ K and is reduced to about $^{\rm th}$ of its initial abundance in this pulse.
+ In later pulses the temperature erows reaching 3.05x105 Ix in the last thermal pulse so that iis destroved with even higher elliciency., In later pulses the temperature grows reaching $3.05 \times 10^8$ K in the last thermal pulse so that is destroyed with even higher efficiency.
+ In the very last few. pulses also about of the pproduced is destroved., In the very last few pulses also about of the produced is destroyed.
+ The effect of the partial mixing zone appears alter the 11' thermal pulse where we observe large changes to the intershellabundances., The effect of the partial mixing zone appears after the $^{\rm th}$ thermal pulse where we observe large changes to the intershellabundances.
+ For example. the amount of aand ssucddenly increase: in the 11 thermal pulse the abundance of iis about 2.5 times higher than that in the 10 thermal pulse.," For example, the amount of and suddenly increase: in the $^{\rm th}$ thermal pulse the abundance of is about 2.5 times higher than that in the $^{\rm th}$ thermal pulse."
+ The final abundance of iin the intershell is ~ higher with respect to the case with no partial mixing zone included (shown in 1))., The final abundance of in the intershell is $\simeq$ higher with respect to the case with no partial mixing zone included (shown in ).
+ The extent in mass and the proton profile of the partial mixing zone are very uncertain parameters., The extent in mass and the proton profile of the partial mixing zone are very uncertain parameters.
+ Most studies that have sell-consistentlv. produced a partially mixed zone find that the extent in mass is smaller than the 0.001. vvalue that we have used., Most studies that have self-consistently produced a partially mixed zone find that the extent in mass is smaller than the 0.001 value that we have used.
+ The computed M. is of the order of  wwith rotation. of LO? wwith overshoot (but depending on the free overshoot parameter!)," The computed $M_{pmz}$ is of the order of $^{-6}$ with rotation, of $^{-5}$ with overshoot (but depending on the free overshoot parameter!)"
+ and of ! wwith gravitational waves., and of $^{-4}$ with gravitational waves.
+ A partial mixing zone of larger extent. 5 x 'AL... was reported to result [rom semiconvection in a low-metallicity star (ILollowell&Iben1955).," A partial mixing zone of larger extent, 5 $\times$ $^{-4}$, was reported to result from semiconvection in a low-metallicity star \citep{hollowell:88}."
+. On the other hand previous nucleosvuthesis studies have artificially. considered partial mixing zones of extent up to 1/10 of the mass of the convective pulse., On the other hand previous nucleosynthesis studies have artificially considered partial mixing zones of extent up to 1/10 of the mass of the convective pulse.
+ To check the uncertainty introduced by the extent of the partial mixing zone we varied (his parameter (hus computing three cases in total: one without zone included. and the other two wilh the mass of the zone equal to {ντ 0.001 and 0.002. AL.," To check the uncertainty introduced by the extent of the partial mixing zone we varied this parameter thus computing three cases in total: one without zone included, and the other two with the mass of the zone equal to $M_{pmz}$ = 0.001 and 0.002 ."
+.. The resulis are presented in ancl show that the variation of the final abundance, The results are presented in and show that the variation of the final abundance
+extrapolating the hyclro-dyvnamical quantities instead of performing a new data exchange (see Section. 3.4).,extrapolating the hydro-dynamical quantities instead of performing a new data exchange (see Section 3.4).
+ The third svnchronization is really a problem. since it is a the end. of the computation. where all the unbalance goes accumulating.," The third synchronization is really a problem, since it is at the end of the computation, where all the unbalance goes accumulating."
+ To better analyze this point. it is worth recalling that most of the computing time is spen calculating the gravitational ancl hyedro-dyvnamical forces. and that it is done only for the subsample of particles (the so called active particles). which are occupying the smalles time bins. and are actually svnchronized. with the entire system (see Section lt seems reasonable to conclude that a good. work balance could be obtained defining a work-loac criterion only looking at the active particles.," To better analyze this point, it is worth recalling that most of the computing time is spent calculating the gravitational and hydro-dynamical forces, and that it is done only for the subsample of particles (the so called active particles), which are occupying the smallest time bins, and are actually synchronized with the entire system (see Section It seems reasonable to conclude that a good work balance could be obtained defining a work-load criterion only looking at the active particles."
+ However this can often lead: to severe memory. problems. for instance any time the number of active particles is a small fraction of the total number of particles.," However this can often lead to severe memory problems, for instance any time the number of active particles is a small fraction of the total number of particles."
+ In this case in fact most of the particles night be handled by a few To avoid such problems we decided to use in any case the criterion discussed in Section 3.2., In this case in fact most of the particles might be handled by a few To avoid such problems we decided to use in any case the criterion discussed in Section 3.2.
+ In this section we are going to present some applications of our code. aiming at testing its capability to reproduce," In this section we are going to present some applications of our code, aiming at testing its capability to reproduce"
+emission atAA.,emission at.
+. We will ¢liscuss it more in detail in the next section., We will discuss it more in detail in the next section.
+ We fitted the red wing of the DLA using the MIDAS fitlyman package., We fitted the red wing of the DLA using the MIDAS fitlyman package.
+ The resulting column. density was determined to be log Njj/em =20.73+0.05., The resulting column density was determined to be log $_H$ $^{-2}=$ $\pm$ 0.05.
+ The HI column density lies below the average neutral hydrogen column density for GRB-DLAs of log Nj/cm 221.6 (Jakobsson et al., The HI column density lies below the average neutral hydrogen column density for GRB-DLAs of log $_H$ $^{-2}=$ 21.6 (Jakobsson et al.
+ 2006: Fynbo et al., 2006; Fynbo et al.
+ 2009)., 2009).
+ This is in agreement with the simulations of Nagamine et al. (, This is in agreement with the simulations of Nagamine et al. (
+2008) which show that the mean DLA column density decreases with increasing redshift.,2008) which show that the mean DLA column density decreases with increasing redshift.
+ On the other hand. the relatively low number of GRBs at redshift >3—4 with measured HI column density and the probable observational bias against the most dusty environments (Jakobsson et al.," On the other hand, the relatively low number of GRBs at redshift $\geq 3-4$ with measured HI column density and the probable observational bias against the most dusty environments (Jakobsson et al."
+ 2006: Fynbo et al., 2006; Fynbo et al.
+ 2009) do not enable yet to firmly check the existence of this anti-correlation., 2009) do not enable yet to firmly check the existence of this anti-correlation.
+ The majority of the absorption lines detected are saturated and therefore do not allow a reliable determination of the column density., The majority of the absorption lines detected are saturated and therefore do not allow a reliable determination of the column density.
+ Mildly saturated lines. which we define here as lines with an EW of «AA.. do not lie on the linear part of the curve of growth and the derived column densities can only be considered as lower limits.," Mildly saturated lines, which we define here as lines with an EW of $<$, do not lie on the linear part of the curve of growth and the derived column densities can only be considered as lower limits."
+ For SIL. Sill* and OI (EWs from 0.30 to 0.49. Tab.," For SII, SiII* and OI (EWs from 0.30 to 0.49, Tab."
+ 2) we obtain. assuming a linear relation between the EW and the colunn density. log N/cm 215.3. log N/em 213.4 and log ΝΟΥ 214.8. respectively.," 2) we obtain, assuming a linear relation between the EW and the column density, log $^{-2}=$ 15.3, log $^{-2}=$ 13.4 and log $^{-2}=$ 14.8, respectively."
+ We take the column density limit derived from SII to determine the metallicity of the host as SII is not affected by depletion onto dust., We take the column density limit derived from SII to determine the metallicity of the host as SII is not affected by depletion onto dust.
+ Using the solar abundances reported in. Asplund et al. (, Using the solar abundances reported in Asplund et al. (
+2009). we find the metallicity in the host along the line of sight to be or Z>0.27Ze.,"2009), we find the metallicity in the host along the line of sight to be $[M/H] > -0.57$ or $Z>0.27\;\Zsun$."
+ We continued to monitor the field of GRB 090205 at late times to further study the GRB host., We continued to monitor the field of GRB 090205 at late times to further study the GRB host.
+ We obtained an image in the R-band with the FORSI camera ~20.3 days after the trigger., We obtained an image in the $R-$ band with the FORS1 camera $\sim 20.3$ days after the trigger.
+ A faint object at a magnitude of R45~26.4+0.3 has been identified very close to the position of the GRB afterglow., A faint object at a magnitude of $R_{AB} \sim 26.4 \pm 0.3$ has been identified very close to the position of the GRB afterglow.
+ This detection represents a flattening in the R—-band light curve. that we interpret as due to the host galaxy (Fig.," This detection represents a flattening in the $R-$ band light curve, that we interpret as due to the host galaxy (Fig."
+ 3)., 3).
+"4 Further I-band monitoring. carried out 46.3 days after reveals an object with Jy,=25.2+0.1."," Further $I-$ band monitoring, carried out $\sim 46.3$ days after reveals an object with $I_{AB} = 25.2 \pm 0.1$."
+ Comparing this detection with the previous one obtained in the A—-band. the resulting unabsorbed R—-7 color is consistent with that of the afterglow.," Comparing this detection with the previous one obtained in the $R-$ band, the resulting unabsorbed $R-I$ color is consistent with that of the afterglow."
+ This suggests a flattening of the light curve in the /—band too. in agreement with the hypothesis that we detected the host," This suggests a flattening of the light curve in the $I-$ band too, in agreement with the hypothesis that we detected the host"
+mass regime is negligible.,mass regime is negligible.
+" In Figure 4,, we plot the resulting distributions and the evolution of the dust distribution slope when we vary the parametersQsc, S, and s."," In Figure \ref{fig:QSs}, we plot the resulting distributions and the evolution of the dust distribution slope when we vary the parameters$Q_{\rm sc}$, $S$, and $s$."
+" These are all variables in the tensile strength curve, which is given as (Benz&Asphaug1999) The variable Qs. is a global scaling factor, S is the scaling of the strength-dominated regime, s is the power dependence on particle size of the strength-dominated regime, G is the scaling of the gravity-dominated regime, and g is the power dependence on particle size of the gravity-dominated regime."," These are all variables in the tensile strength curve, which is given as \citep{benz99}
+ The variable $Q_{\rm sc}$ is a global scaling factor, $S$ is the scaling of the strength-dominated regime, $s$ is the power dependence on particle size of the strength-dominated regime, $G$ is the scaling of the gravity-dominated regime, and $g$ is the power dependence on particle size of the gravity-dominated regime."
+" Variations in the gravity-dominated regime of the curve (G and g) do not have significant effects on the equilibrium dust-mass distribution, so we do not consider these parameters further."," Variations in the gravity-dominated regime of the curve $G$ and $g$ ) do not have significant effects on the equilibrium dust-mass distribution, so we do not consider these parameters further."
+ The tensile strength curve has been extensively studied for decades., The tensile strength curve has been extensively studied for decades.
+" However, as it is dependent on various material properties and the collisional velocity (Stewart&Leinhardt2009;Stewart2012),, its parameters do not have universally applicable values."," However, as it is dependent on various material properties and the collisional velocity \citep{stewart09,leinhardt12}, its parameters do not have universally applicable values."
+ Determining the tensile strength curve at large and small sizes is also extremely difficult experimentally., Determining the tensile strength curve at large and small sizes is also extremely difficult experimentally.
+ The slope s of the strength curve in the strength-dominated regime depends on the Weibull flaw-size distribution., The slope $s$ of the strength curve in the strength-dominated regime depends on the Weibull flaw-size distribution.
+ Its measured values range anywhere between -0.7 and -0.3 (Holsappleetal.2002)., Its measured values range anywhere between -0.7 and -0.3 \citep{holsapple02}.
+". Steeper values of s make smaller materials harder to disrupt, which results in a steeper dust distribution slope."," Steeper values of $s$ make smaller materials harder to disrupt, which results in a steeper dust distribution slope."
+" At s—1.0, the smallest particles are hard enough to resist catastrophic disruption even when the projectile mass equals the target mass."," At $s=1.0$, the smallest particles are hard enough to resist catastrophic disruption even when the projectile mass equals the target mass."
+" This results in a mass distribution with a slope equal to the redistribution slope y=1.83 at the smallest scales, while in the fitted region it is significantly steeper."," This results in a mass distribution with a slope equal to the redistribution slope $\gamma=1.83$ at the smallest scales, while in the fitted region it is significantly steeper."
+" At s=0.6, the smallest particles are still able to destroy each other and generate a dust distribution slope that is close to 1.91."," At $s=0.6$, the smallest particles are still able to destroy each other and generate a dust distribution slope that is close to 1.91."
+ The scaling constants of the tensile strength curve are the dominant parameters in the evolution toward the quasi steady-state distribution., The scaling constants of the tensile strength curve are the dominant parameters in the evolution toward the quasi steady-state distribution.
+" When reducing the complete tensile strength curve scaling (ες, wave structures form more easily as a particle becomes capable of affecting the evolution of particles much larger than itself (see Paper I)."," When reducing the complete tensile strength curve scaling $Q_{\rm sc}$, wave structures form more easily, as a particle becomes capable of affecting the evolution of particles much larger than itself (see Paper I)."
+" When upscaling the tensile strength curve, the quasi steady-state distribution slope starts to resemble the redistribution slope, as it is the particle redistributions that lead the evolution of the particle mass distribution."," When upscaling the tensile strength curve, the quasi steady-state distribution slope starts to resemble the redistribution slope, as it is the particle redistributions that lead the evolution of the particle mass distribution."
+" When varying the scaling of only the strength side of the curve S, similar effects can be seen."," When varying the scaling of only the strength side of the curve $S$, similar effects can be seen."
+ We have shown that drastic offsets in the collision parameter values result in slight changes to the quasi steady state mass distribution slope., We have shown that drastic offsets in the collision parameter values result in slight changes to the quasi steady state mass distribution slope.
+" We conclude that, for all reasonable values of the collisional parameters, the quasi steady-state dust-mass distribution slope is larger or equal to 1.88."," We conclude that, for all reasonable values of the collisional parameters, the quasi steady-state dust-mass distribution slope is larger or equal to $1.88$."
+" There are a number of parameters that can change from one collisional system to another: the material density p, the minimum and the maximum particle mass in the system m4, and Mmax, the radial distance R, height h, and width AR of the disk, and the spectral type of the central star."," There are a number of parameters that can change from one collisional system to another: the material density $\rho$, the minimum and the maximum particle mass in the system $m_{\rm min}$ and $m_{\rm max}$, the radial distance $R$, height $h$, and width $\Delta R$ of the disk, and the spectral type of the central star."
+" All these parameters affect three properties of the collisional model: the blow-out mass, the collisional velocity, and the number density of particles."," All these parameters affect three properties of the collisional model: the blow-out mass, the collisional velocity, and the number density of particles."
+ Varying these parameters will change the timescale of the evolution and affect the quasi steady-state distribution slope., Varying these parameters will change the timescale of the evolution and affect the quasi steady-state distribution slope.
+" In this subsection, we analyze the effects of varying the radial distance on the equilibrium mass distribution through the dependence of the collisional velocity on radius."," In this subsection, we analyze the effects of varying the radial distance on the equilibrium mass distribution through the dependence of the collisional velocity on radius."
+ Modifying either disk parameters AR and h or the spectral-type of the star would have similar effects., Modifying either disk parameters $\Delta R$ and $h$ or the spectral-type of the star would have similar effects.
+ We defer discussion of the variations in the timescales to a future paper., We defer discussion of the variations in the timescales to a future paper.
+" In the left panel of Figure 5,, we show the effects of varying the radial distance, R, on the mass distribution evolved to 10 Gyr."," In the left panel of Figure \ref{fig:R}, we show the effects of varying the radial distance, $R$ , on the mass distribution evolved to 10 Gyr."
+" Decreasing the radial distance will increase the collisional velocity, resulting in the appearance of waves at the small-mass end of the distribution."," Decreasing the radial distance will increase the collisional velocity, resulting in the appearance of waves at the small-mass end of the distribution."
+ It also generates a much more pronounced, It also generates a much more pronounced
+lin Jin lin Jin,.1in .1in .1in .1in
+"This ""central image” or ""odd image” corresponds (o the central maxinmum in the surface.",This “central image” or “odd image” corresponds to the central maximum in the time-delay surface.
+ However. of the ~80 known svstenis in which a galaxw produces multiple inages of a quasar or radio source. almost all consist of either 2 or 4 images.," However, of the $\sim$ 80 known systems in which a galaxy produces multiple images of a quasar or radio source, almost all consist of either 2 or 4 images."
+" One system. APM 032794-5255. definitely has three quasar images (Lewisetal.2002b).. but the third image may be due to à ""naked cusp configuration rather (han a central time-delay maximum (Lewisοἱal.2002a)."," One system, APM 08279+5255, definitely has three quasar images \citep{lewis02b}, but the third image may be due to a “naked cusp” configuration rather than a central time-delay maximum \citep{lewis02a}."
+. No central image has been identified definitivelv. even among the radio lenses. where the central image could in principle be seen throueh the light and dust of the lens galaxy.," No central image has been identified definitively, even among the radio lenses, where the central image could in principle be seen through the light and dust of the lens galaxy."
+ The absence of central images has been attributed to the demaguilication of the image bv the high central surface density of the lens galaxy (see.e.g...Narasimha.Subramanian.&Chitre19856:Wallington&Naravan1993:RusinMa2001:EvansIIunter2002:Keeton2002 ).," The absence of central images has been attributed to the demagnification of the image by the high central surface density of the lens galaxy \citep[see,
+e.g.,][]{nsc86,wn93,rm01,eh02,keeton02}."
+". An ""extra"" radio component has been detected in several radio lenses. and in three of those cases il is still unclear whether (he extra component is an additional quasar image or [anl radio emission from an active galactic nucleus (AGN)."," An “extra” radio component has been detected in several radio lenses, and in three of those cases it is still unclear whether the extra component is an additional quasar image or faint radio emission from an active galactic nucleus (AGN)."
+ Those three svstems are 0957+561 (Harvaneketal.1997).. MG. J1131--0456 (Chen&Hewitt 1993).. and the subject of this paper: PAIN J16320033 (Winnetal.2002).," Those three systems are 0957+561 \citep{harvanek97}, MG J1131+0456 \citep{chen93}, , and the subject of this paper: PMN J1632–0033 \citep{winn02}."
+. Radio components| A and B of J16320033 are two images2 of a z=3.42 quasar.[ but the third. faint component C is of unknown origin (see Eig 1)).," Radio components A and B of J1632–0033 are two images of a $z=3.42$ quasar, but the third, faint component C is of unknown origin (see Fig \ref{fig:merlin}) )."
+ In an UST image. the lens ealaxy appears (ο be a fairly. circular earlv-tvpe galaxy. wilh an optical effective radius of x(72. but it is (oo fnt to be characterizecl in detail.," In an HST image, the lens galaxy appears to be a fairly circular early-type galaxy with an optical effective radius of $\approx 0\farcs2$, but it is too faint to be characterized in detail."
+ The position of component C is near the optical lens galaxy. position. as expected for either an AGN or a central odd image.," The position of component C is near the optical lens galaxy position, as expected for either an AGN or a central odd image."
+ If A. D. and C are all images of the same quasar. then they should have similar radio spectra.," If A, B, and C are all images of the same quasar, then they should have similar radio spectra."
+ llowever. Winnetal.(2002) only presented measurements of the flux density of C al one frequency.," However, \citet{winn02} only presented measurements of the flux density of C at one frequency."
+ llere we present new VLBA data that. in combination with the previous data. allow us {ο measure (he flux densities of all three components at four widely spaced frequencies.," Here we present new VLBA data that, in combination with the previous data, allow us to measure the flux densities of all three components at four widely spaced frequencies."
+ Our analvsis shows that C has a signifieantlv different radio spectrum than A and D. We argue that there are (wo alternative explanations for this difference., Our analysis shows that C has a significantly different radio spectrum than A and B. We argue that there are two alternative explanations for this difference.
+ Either C 1s not a (bird quasar image. and instead is the active nucleus of the lens galaxy: or. Cis a third image. but its spectral index is inverted by [ree-Iree absorption in the lens galaxy.," Either C is not a third quasar image, and instead is the active nucleus of the lens galaxy; or, C is a third image, but its spectral index is inverted by free-free absorption in the lens galaxy."
+ Whichever of these possibilities is true. the additional information from the VLBA data provides more constraints on mass models of the lens galaxy (han are usually available [or a lwo-lmage lens.," Whichever of these possibilities is true, the additional information from the VLBA data provides more constraints on mass models of the lens galaxy than are usually available for a two-image lens."
+ If (he central component is a third image. (hen ils position ancl flux are valuable constraints.," If the central component is a third image, then its position and flux are valuable constraints."
+ If (he central component is the lens galaxy. then our data provide an unusually precise position lor the lens galaxy. ancl a strong upper limit on the flux of a third image.," If the central component is the lens galaxy, then our data provide an unusually precise position for the lens galaxy, and a strong upper limit on the flux of a third image."
+ The VLBA maps also reveal the relative orientations ofthe radio jets associated with, The VLBA maps also reveal the relative orientations ofthe radio jets associated with
+well delinecl relationship between radial velocity (corrected to the ealactocentric staucdard of rest) ancl distance was ciscc»vered for all members of the Sculptor group (Jejen et 11993). indicating that radial velocity iuay. be a good distance iudicator for Sculptor C'oup members.,"well defined relationship between radial velocity (corrected to the galactocentric standard of rest) and distance was discovered for all members of the Sculptor group (Jerjen et 1998), indicating that radial velocity may be a good distance indicator for Sculptor Group members."
+ Jerjen et ((1998) propose tliat he Sculptor Group and the Local GroIp are pa1 )‘the same Superealaclic structure (a part of tlje Coma - Scuptor Cloud deliueated by Tully Fisher 1987)., Jerjen et (1998) propose that the Sculptor Group and the Local Group are part of the same Supergalactic structure (a part of the Coma - Sculptor Cloud delineated by Tully Fisher 1987).
+ A «vnaniical picture of the Sculptor Group can be fourdin Whiting (1999)., A dynamical picture of the Sculptor Group can be found in Whiting (1999).
+ Miler (1991). Cotté (1995). aid Jerjeu et ((2000) indepeucdenutly searched for aud foμιά uew dwarf galaxy meubers of the Sculptor Group.," Miller (1994), Côtté (1995), and Jerjen et (2000) independently searched for and found new dwarf galaxy members of the Sculptor Group."
+ The Sculptor Ciroup ix now known to coualu 16 dwar. irregular galaxies (Cótté e 11997). some of which are amongst the lowest luminosity dwarl ir'egulars known.," The Sculptor Group is now known to contain 16 dwarf irregular galaxies (Côtté et 1997), some of which are amongst the lowest luminosity dwarf irregulars known."
+ The present Sculptor Croup membe‘ship situation can be found in Jejen et ((]1998: 2000)., The present Sculptor Group membership situation can be found in Jerjen et (1998; 2000).
+ Miler (1996) noted that several Sculptor Ciroup dis were undeected iu Ha eCLUSSLOL al 'easonably sensitive levels., Miller (1996) noted that several Sculptor Group dIs were undetected in $\alpha$ emission at reasonably sensitive levels.
+ This has several implications for the present studs., This has several implications for the present study.
+ First. as Miller (1996) poiuted out. the Sculpto .-‘oup dis inay have relatively low aveage SFRs when compared o Local Group cs.," First, as Miller (1996) pointed out, the Sculptor Group dIs may have relatively low average SFRs when compared to Local Group dIs."
+ Secoud. alhoih some of the Miller (1996) sample have now been detected iu Ha at low levels (see 2.2). there remain three dwarf galaxies iu tle Sculptor Group (SDIC:. DDO 6. and LOCA 195) with ilὁ properties of “trainition type” οιdaxies (no detectable HII 'eglous aud Mua Ευ values in he range of dl galaxies.," Second, although some of the Miller (1996) sample have now been detected in $\alpha$ at low levels (see 2.2), there remain three dwarf galaxies in the Sculptor Group (SDIG, DDO 6, and UGCA 438) with the properties of “transition type” galaxies (no detectable HII regions and $_{HI}$ $_B$ values in the range of dI galaxies)."
+" Thus. the Seuptor Ciroup sample allows is to Investigate this less common ype of galaxy (see. e(δι, Miller et 22001)."," Thus, the Sculptor Group sample allows us to investigate this less common type of galaxy (see, e.g., Miller et 2001)."
+ Third. Ha surveys iceitilvine HID regions are a iecessary first step in s«lying the ISM abuudauces of galaxies.," Third, $\alpha$ surveys identifying HII regions are a necessary first step in studying the ISM abundances of galaxies."
+ Since relatively high 5urface brightuess HII regions are isually required for an accurate abuudance allaysis. this meaus that we may uot be able to accuratey measure the chemical abunclauces iu all oftie known Sculptor Group cls.," Since relatively high surface brightness HII regions are usually required for an accurate abundance analysis, this means that we may not be able to accurately measure the chemical abundances in all of the known Sculptor Group dIs."
+" Here we present deep CCD Ha tnagine of eight SciIptor group dls obtained at the “Danish” 1οι telescope located at the European Southern Obse""atory.", Here we present deep CCD $\alpha$ imaging of eight Sculptor group dIs obtained at the “Danish” 1.5m telescope located at the European Southern Observatory.
+ We have detected. HII regious iu all eight dIs observed. in addition to the two dls detected in Ha by Miller (1996).," We have detected HII regions in all eight dIs observed, in addition to the two dIs detected in $\alpha$ by Miller (1996)."
+ We present coordinates and fluxes for the HIT regions aud estimated SFRs for these galaxies., We present coordinates and fluxes for the HII regions and estimated SFRs for these galaxies.
+ Acditionally. we have used the CTIO Lin to obtain optical spectra of ten HII regious located iu five of these dlIs aud these observatious are presented in a companion paper.," Additionally, we have used the CTIO 4-m to obtain optical spectra of ten HII regions located in five of these dIs and these observations are presented in a companion paper."
+ By comparing observatious of the dwarf e@alaxies in the next uearest group to other well defined samples. we hope to better uuclerstaucl basic questions such as: what are the average properties of the dwarf galaxies?," By comparing observations of the dwarf galaxies in the next nearest group to other well defined samples, we hope to better understand basic questions such as: what are the average properties of the dwarf galaxies?"
+ aud. are there observable signatures of euvironmental influences on galaxy evolution?," and, are there observable signatures of environmental influences on galaxy evolution?"
+ All galaxies observed were chosen from the list of Cótté et (1997)., All galaxies observed were chosen from the list of Côtté et (1997).
+ Table 1 lists some of, Table 1 lists some of
+at least one major merger and this [raction drops to around 20 for haloes with masses of around 10h.ΑΗ.,at least one major merger and this fraction drops to around 20 for haloes with masses of around $10^{11}h^{-1}M_{\odot}$.
+ We now investigate he fraction of that have had. major 1nereers., We now investigate the fraction of that have had major mergers.
+ The results are shown as a dotted [ine in the micle panel of Fig. 2.., The results are shown as a dotted line in the middle panel of Fig. \ref{fig:haloGal}.
+ Rather than rising steeply as a function of mass. t1ο galaxy major merger fraction remains close to zero up to a stelar mass of RE10Ad. anc hen rises sharply.," Rather than rising steeply as a function of mass, the galaxy major merger fraction remains close to zero up to a stellar mass of $10^{10.5} M_{\odot}$ and then rises sharply."
+ This is somewhat surprising in view of he behaviour of the same quantiv for dark matter haloes. xotted. in the right-hand. panel of Fig. 1..," This is somewhat surprising in view of the behaviour of the same quantity for dark matter haloes, plotted in the right-hand panel of Fig. \ref{fig:halofrac}."
+ For reference. we have plotted the relaion between the stellar mass of a central galaxy and the mass of its host halo in the lef panel of Fig. 2..," For reference, we have plotted the relation between the stellar mass of a central galaxy and the mass of its host halo in the left panel of Fig. \ref{fig:haloGal},"
+ as predicted. by the semi-analvtie moclels we use in this study (2)..., as predicted by the semi-analytic models we use in this study \citep{delucia2007}.
+ This mean relation can be use o transform between central eaaxv mass and halo mass in an approximate wav (this cconversion. neglects scatter )etween the two quantities and t1e fact that some galaxies are actually satellite svstenis)., This mean relation can be used to transform between central galaxy mass and halo mass in an approximate way (this conversion neglects scatter between the two quantities and the fact that some galaxies are actually satellite systems).
+ Lf11ο fraction of galaxies with major mergers followed the relation derived for their host jaloes. this would. vield the solid curve in the middle panel of Fig. 2..," If the fraction of galaxies with major mergers followed the relation derived for their host haloes, this would yield the solid curve in the middle panel of Fig. \ref{fig:haloGal}."
+ Why are the merging histories of galaxies and heir host haloes so cilferent?, Why are the merging histories of galaxies and their host haloes so different?
+ Once two dark matter haloes merge. t1ο galaxies inside hem will merge together over a timescale tjab is determined » dynamical friction.," Once two dark matter haloes merge, the galaxies inside them will merge together over a timescale that is determined by dynamical friction."
+ Upon investigation. we find that nearly all galaxies that have experienced major mergers are ocated in dar.k matter haloes that have also experienced a major merger.," Upon investigation, we find that nearly all galaxies that have experienced major mergers are located in dark matter haloes that have also experienced a major merger."
+ There are almost no galaxy major mergers hat have occiured in à halo that has only experienced. a minor merger (~0.15 percent)., There are almost no galaxy major mergers that have occurred in a halo that has only experienced a minor merger $\sim 0.15$ percent).
+ However. t1 Converse Is not rue: This is illustrated in he rieht-hanc panel of Fig. 2..," However, the converse is not true; This is illustrated in the right-hand panel of Fig. \ref{fig:haloGal}."
+ The solid line shows the number of major mergers experienced. by the progenitor of a pressent-day central galaxy as a function of its mass., The solid line shows the number of major mergers experienced by the progenitor of a present-day central galaxy as a function of its mass.
+ “Phe dasied line shows the number of major mergers experienced by their progenitorgaleries., The dashed line shows the number of major mergers experienced by their progenitor.
+ As can be seen. the number of major mergers experienced by the progenitor," As can be seen, the number of major mergers experienced by the progenitor"
+2--21-1Osawa.Mitaka.Received:accepted Asteroid is currently among the potential targets of future interplanetary exploration missions.,"-21-1\tikzmark{mainBodyEnd0} \tikzmark{mainBodyStart1}Osawa,\tikzmark{mainBodyEnd1} \tikzmark{mainBodyStart2}Mitaka,\tikzmark{mainBodyEnd2} \tikzmark{mainBodyStart3}Tokyo\tikzmark{mainBodyEnd3} \tikzmark{mainBodyStart4}181-8588,\tikzmark{mainBodyEnd4} \tikzmark{mainBodyStart5}Japan;\tikzmark{mainBodyEnd5}
+ \tikzmark{mainBodyStart6}}\tikzmark{mainBodyEnd6}
+
+ \date{Received; accepted}
+
+
+% \abstract{}{}{}{}{} 
+% 5 {} token are mandatory
+ 
+ \abstract
+ % context heading (optional)
+ % {} leave it empty if necessary 
+ {Near-Earth asteroid 162173 (1999~JU3) is a potential flyby and rendezvous target
+ for interplanetary missions because of its easy to reach orbit. The physical
+ and thermal properties of the asteroid are relevant for establishing the
+ scientific mission goals and also importa\tikzmark{mainBodyStart7}ant\tikzmark{mainBodyEnd7} \tikzmark{mainBodyStart8}in\tikzmark{mainBodyEnd8} \tikzmark{mainBodyStart9}the\tikzmark{mainBodyEnd9} \tikzmark{mainBodyStart10}context\tikzmark{mainBodyEnd10} \tikzmark{mainBodyStart11}of\tikzmark{mainBodyEnd11} \tikzmark{mainBodyStart12}near-Earth\tikzmark{mainBodyEnd12}
+ \tikzmark{mainBodyStart13}object\tikzmark{mainBodyEnd13} \tikzmark{mainBodyStart14}studies\tikzmark{mainBodyEnd14} \tikzmark{mainBodyStart15}in\tikzmark{mainBodyEnd15} \tikzmark{mainBodyStart16}general.}\tikzmark{mainBodyEnd16}
+ % aims heading (mandatory)
+ {Our goal was to derive key physical parameters such as shape, spin-vector,
+ size, geometric albedo, and surface properties of 162173 (1999~JU3).}\tikzmark{mainBodyStart17}}\tikzmark{mainBodyEnd17}
+ % methods heading (mandatory)
+ {With three sets of published thermal observations (ground-based N-band, Akari IRC,
+ Spitzer IRS), we applied a thermophysical model to derive the radiometric properties
+ of the asteroid.
+ The calculations were performed for the full range of possible shape and spin-vector
+ solutions derived from the available sample of visual lightcurve observations.}\tikzmark{mainBodyStart18}}\tikzmark{mainBodyEnd18}
+ % results heading (mandatory)
+ {The near-Earth asteroid 162173 (1999~JU3) has an effective diameter of
+ 0.87\,$\pm$\,0.03\,km and a geometric albedo of 0.070\,$\pm$\,0.006.
+ The $\chi^2$-test reveals a strong preference for a retrograde sense of rotation
+ with a spin-axis orientation of $\lambda_{\mathrm{ecl}}\,=\,73^{\circ}$, $\beta_{\mathrm{ecl}}\,=\,-62^{\circ}$ and P$_{\mathrm{sid}}$\,=\,7.63$\pm$0.01\,h.
+ The most likely thermal inertia ranges between 200 and
+ 600\,J\,m$^{-2}$\,s$^{-0.5}$\,K$^{-1}$, about a factor of 2 lower than the value
+ for 25143~Itokawa. This indicates that the surface lies somewhere between a thick-dust regolith
+ and a rock/boulder/cm-sized, gravel-dominated surface like that of 25143~Itokawa.
+ Our analysis represents the first time that
+ shape and spin-vector information has been derived from a combined data set of 
+ visual lightcurves (reflected light) and mid-infrared photometry and spectroscopy
+ (thermal emission).}\tikzmark{mainBodyStart19}}\tikzmark{mainBodyEnd19}
+ % conclusions heading (optional), leave it empty if necessary 
+ {}\tikzmark{mainBodyStart20}} Asteroid is currently among the potential targets of future interplanetary exploration missions."
+ The target is relatively easy to reach with state-of-the-art mission capabilities. and it offers high scientific potential (Binzel et citebinzelO4)).," The target is relatively easy to reach with state-of-the-art mission capabilities, and it offers high scientific potential (Binzel et \\cite{binzel04}) )."
+ This near-Earth asteroid belongs to the C-class objects. which are believed to represent primitive. volatile-rich remnants of the early solar system.," This near-Earth asteroid belongs to the C-class objects, which are believed to represent primitive, volatile-rich remnants of the early solar system."
+ Various aspects of this small body have been covered in some detail in the recent works by Hasegawa et ((2008)) and by Campins et ((2009))., Various aspects of this small body have been covered in some detail in the recent works by Hasegawa et \cite{hasegawa08}) ) and by Campins et \cite{campins09}) ).
+" Hasegawa et ((2008)) use a spherical shape nodel for their radiometric analysis. and alternatively use an ellipsoidal shape model. but without knowing the truespin-vector orientation,"," Hasegawa et \cite{hasegawa08}) ) use a spherical shape model for their radiometric analysis, and alternatively use an ellipsoidal shape model, but without knowing the truespin-vector orientation."
+ The results (radiometric diameter of 200.12kkm. visual geometric albedo of 5070) which indicate a thermal inertia. larger. thar mm ? KK! ). were based on a set of photometric Subaru and observations and connected to simplifiec shape and spin-axis assumptions.," The results (radiometric diameter of $\pm$ km, visual geometric albedo of $^{+0.020}_{-0.015}$ ) which indicate a thermal inertia larger than $^{-2}$ $^{-0.5}$ $^{-1}$ ), were based on a set of photometric Subaru and observations and connected to simplified shape and spin-axis assumptions."
+ Campins et ((2009)) have obtained a single-epochSpitzer infrared spectrograph (IRS) spectrum., Campins et \cite{campins09}) ) have obtained a single-epoch infrared spectrograph (IRS) spectrum.
+ They used a spherical shape model. anc for the spin-pole orientation they used the extreme case of an equatorial retrograde geometry and a prograde solutior published by Abe et ((2008)).," They used a spherical shape model, and for the spin-pole orientation they used the extreme case of an equatorial retrograde geometry and a prograde solution published by Abe et \cite{abe08}) )."
+ Their analysis. based οἱ the single IRS-spectrum and ignoring the data sets publishec by Hasegawa et ((2008)). yielded a value for the thermal inertia of + mm ss? KK'. a diameter estimate of x kkm. and geometric albedo of + 00.01.," Their analysis, based on the single IRS-spectrum and ignoring the data sets published by Hasegawa et \cite{hasegawa08}) ), yielded a value for the thermal inertia of $\pm$ $^{-2}$ $^{-0.5}$ $^{-1}$, a diameter estimate of $\pm$ km, and geometric albedo of $\pm$ 0.01."
+ Despite the simplifications in. shape and. spin-vector properties. both sets of published radiometric diameter and albedo values agree within the given uncertainties.," Despite the simplifications in shape and spin-vector properties, both sets of published radiometric diameter and albedo values agree within the given uncertainties."
+ Both teams also favoured a relatively high thermal inertia (close to that of Itokawa))., Both teams also favoured a relatively high thermal inertia (close to that of ).
+ The unknowns of the spin-vector orientatio cause a large uncertainty in the thermal properties., The unknowns of the spin-vector orientation cause a large uncertainty in the thermal properties.
+ Here we re-analyse all available lightcurve observationnr to derive (on the basis of standard y lightcurve inversio techniques) all matching spin-vector and shape solutions (Sect. 2)), Here we re-analyse all available lightcurve observations to derive (on the basis of standard $\chi^2$ lightcurve inversion techniques) all matching spin-vector and shape solutions (Sect. \ref{sec:sv}) ).
+ The full possible range for shape. spin-axis orientation. and rotation period was then used as input for à thermophysical y analysis of all available thermal observations (Sect. 3))," The full possible range for shape, spin-axis orientation, and rotation period was then used as input for a thermophysical $\chi^2$ analysis of all available thermal observations (Sect. \ref{sec:tpm}) )"
+ with the goal of deriving radiometric sizes. albedos and thermal inertias.," with the goal of deriving radiometric sizes, albedos and thermal inertias."
+ At the same time. we determined the most likely shape-solution. rotation period and spin-axis orientation (Sect. 4)).," At the same time, we determined the most likely shape-solution, rotation period and spin-axis orientation (Sect. \ref{sec:results}) )."
+ A detailed list of the available photometric observations 15 presented in Table 1.., A detailed list of the available photometric observations is presented in Table \ref{tbl:lc_obs}. .
+ There are about 40 dedicated visual lightcurve data sets spread over more than 270 days., There are about 40 dedicated visual lightcurve data sets spread over more than 270 days.
+ About half of the lighteurveswere calibrated: some of the data, About half of the lightcurveswere calibrated; some of the data
+It has often been assumed that binary open cluster formation in the Milkv Way is uncommon.,It has often been assumed that binary open cluster formation in the Milky Way is uncommon.
+ In contrast. our results indicate that the lives of primordial binary clusters are violent and hazardous.," In contrast, our results indicate that the lives of primordial binary clusters are violent and hazardous."
+ Close pairs (if formed) merge in a short ümescale. being the shortest for verv eccentric pairs. secondaries in low mass ratio pairs are rapidly destroved. and wicle pairs quickly ionize in the background tidal field due to mass loss.," Close pairs (if formed) merge in a short timescale, being the shortest for very eccentric pairs, secondaries in low mass ratio pairs are rapidly destroyed, and wide pairs quickly ionize in the background tidal field due to mass loss."
+ As open clusters are actually born in close proximity (complexes). these appear to be the genuine reasons behind the apparent lack of binary open clusters in our Galaxy.," As open clusters are actually born in close proximity (complexes), these appear to be the genuine reasons behind the apparent lack of binary open clusters in our Galaxy."
+" Star cluster binarity is. therefore. a transient phenomenon,"," Star cluster binarity is, therefore, a transient phenomenon."
+ The effects and trends. observed in the present set. of simulations are robust [or the range of open cluster parameters studied., The effects and trends observed in the present set of simulations are robust for the range of open cluster parameters studied.
+ As usual. it is potentially dangerous to make unwarranted extrapolations to lareger/smaller or denser clusters.," As usual, it is potentially dangerous to make unwarranted extrapolations to larger/smaller or denser clusters."
+ It could be possible that for much larger anc clenser star clusters the merging timescale is longer., It could be possible that for much larger and denser star clusters the merging timescale is longer.
+ Nevertheless. the absence of binary globular clusters in the Milkv Way appears to indicate that. long-term binary cluster stability is. in fact. unlikely.," Nevertheless, the absence of binary globular clusters in the Milky Way appears to indicate that long-term binary cluster stability is, in fact, unlikely."
+ The role of the gravogvro ellect of the evolution ol merger remnants appears to be well documented in our present work but larger simulations are needed to better understand Che statistical strength of (his process., The role of the gravogyro effect of the evolution of merger remnants appears to be well documented in our present work but larger simulations are needed to better understand the statistical strength of this process.
+ The authors would like to thank S. Aarseth ancl Ix. Nitadori for providing the code used in (his research., The authors would like to thank S. Aarseth and K. Nitadori for providing the code used in this research.
+ The authors also would like to acknowledge the help of J. L. Mazo Friaas. N. Pedone. and P. Chan of Dell Computer with the T5500+S1070 svstem.," The authors also would like to acknowledge the help of J. L. Mazo as, N. Pedone, and P. Chan of Dell Computer with the T5500+S1070 system."
+ In preparation of this paper. we made use of the NASA Astrophysics Data System and the astro-ph e-print server.," In preparation of this paper, we made use of the NASA Astrophysics Data System and the astro-ph e-print server."
+ This research has made use of the WEBDA database operated at the Institute of Astronomy of the University of Vienna. Austria.," This research has made use of the WEBDA database operated at the Institute of Astronomy of the University of Vienna, Austria."
+ This work also made use of the ALADIN. SIMDAD and VIZIER. databases. operated at the CDS. Strasbourg. France.," This work also made use of the ALADIN, SIMBAD and VIZIER databases, operated at the CDS, Strasbourg, France."
+The situation is more complex for the other three sources. which have cooler Lt colours suggestive of a strong starburst contribution.,"The situation is more complex for the other three sources, which have cooler IR colours suggestive of a strong starburst contribution."
+ As discussed. by Vignati et al. (, As discussed by Vignati et al. (
+1999). [or GC 6240 the bolometric ACN luminosity derived assuming he Quasar SED of Elvis ct al. (,"1999), for NGC 6240 the bolometric AGN luminosity derived assuming the Quasar SED of Elvis et al. ("
+L994). be. Ling~οdus iso 1e same order of the LR. luminosity.,"1994), i.e. $L_{\rm bol}\sim 30 L_{\rm 2-10 kev}$, is of the same order of the IR luminosity."
+ Pherclore. a large raction of the It flux. should arise from the reprocessing of th nuclear UV and X.ray photons. even if a significant contribution from starburst it is also possible. especially at he longer wavelengths.," Therefore, a large fraction of the IR flux should arise from the reprocessing of the nuclear UV and X–ray photons, even if a significant contribution from starburst it is also possible, especially at the longer wavelengths."
+ For Circinus and NGC 4945 the situation is cilferent., For Circinus and NGC 4945 the situation is different.
+ For low-luminosity AXGNs like these two Sevferts Gvhich are he sources with the lowest intrinsic Lx in our sample by uw). the bolometric correction should be lower (about. 10 instead of 30). and it is possible that the LE Lux is actually dominated by starburst emission (it is worth noting that the required starburst Luminosity for Cireinus and NGC 4945 is much lower than that would be required to account for the IR. luminosity of NGC 6240) Finally. a few general considerations on the vexecl question of the existence. of type 2 quasars may be appropriate here.," For low-luminosity AGNs like these two Seyferts (which are the sources with the lowest intrinsic $L_X$ in our sample by far), the bolometric correction should be lower (about 10 instead of 30), and it is possible that the IR flux is actually dominated by starburst emission (it is worth noting that the required starburst luminosity for Circinus and NGC 4945 is much lower than that would be required to account for the IR luminosity of NGC 6240) Finally, a few general considerations on the vexed question of the existence of type 2 quasars may be appropriate here."
+ At least two sources in our sample have a 210 keV luminosity of order o£ 107 erg s.|. and therefore. using the SED of Elvis et al. (," At least two sources in our sample have a 2–10 keV luminosity of order of $^{44}$ erg $^{-1}$, and therefore, using the SED of Elvis et al. ("
+1994). a bolometric luminosity equal or exceeding 107 erg well within the Quasar regime (there are also other examples. as listed by Vignati et al.,"1994), a bolometric luminosity equal or exceeding $^{45}$ erg $^{-1}$, well within the Quasar regime (there are also other examples, as listed by Vignati et al."
+ 1999)., 1999).
+ But even more important is to remark that the QSO 2 debate is largely based. on a misunderstanding. as the case of NGC 6240 makes clear.," But even more important is to remark that the QSO 2 debate is largely based on a misunderstanding, as the case of NGC 6240 makes clear."
+ H£ one assumes the classical optical spectroscopic classification. NGC 6240 is certainly not a tvpe 2 QSO. since it lacks the emission line spectrum typical of the Narrow Line Region.," If one assumes the classical optical spectroscopic classification, NGC 6240 is certainly not a type 2 QSO, since it lacks the emission line spectrum typical of the Narrow Line Region."
+ However. it must not be forgotten that the original classification. even if it has been very useful in the past. is based on AGN properties that are secondary and not needed to define an AGN in the modern sense.," However, it must not be forgotten that the original classification, even if it has been very useful in the past, is based on AGN properties that are secondary and not needed to define an AGN in the modern sense."
+ In X.ray terminology. a tvpe 2 XGN is simply à source in which the X.rav emitting region (i.c. the black hole and its immediate surroundings) is obscured. no matter if the BLR. and νι are visible or not.," In X–ray terminology, a type 2 AGN is simply a source in which the X–ray emitting region (i.e. the black hole and its immediate surroundings) is obscured, no matter if the BLR and NLR are visible or not."
+ Since Xray emission is àfundemmentad property of an AGN. a classification based on Xαν properties is clearly to be preferred. being much less ambiguous.," Since X–ray emission is a property of an AGN, a classification based on X–ray properties is clearly to be preferred, being much less ambiguous."
+ Moreover. and. contrary to the infrared. band. in Xrays aceretion is certainly much more important than emission associated with stellar processes.," Moreover, and contrary to the infrared band, in X–rays accretion is certainly much more important than emission associated with stellar processes."
+ The issue of the fraction and luminosity dependence of “twpe 27 (or. better. obscured) AGN will therefore be addressed. (ancl hopefully settled) by Xray surveys like those that will be performed by and NMM.," The issue of the fraction and luminosity dependence of “type 2"" (or, better, obscured) AGN will therefore be addressed (and hopefully settled) by X–ray surveys like those that will be performed by and XMM."
+ Comptonthick Sevfert 2 galaxies are very likely the most common subclass of AGN in the local Universe (Alaioling et al., Compton–thick Seyfert 2 galaxies are very likely the most common subclass of AGN in the local Universe (Maiolino et al.
+ 1998). and possibly also at high recshifts (Fabian 1999).," 1998), and possibly also at high redshifts (Fabian 1999)."
+ The hard Xαν band is certainly the best with which to study these sources. because part of the nuclear radiation can penetrate the obscuring matter. if the column density does not exceed a few times 1073 cn7.," The hard X–ray band is certainly the best with which to study these sources, because part of the nuclear radiation can penetrate the obscuring matter, if the column density does not exceed a few times $^{24}$ $^{-2}$."
+ This is the reason why BeppoSAX has permitted a great. advance in his limited but important field., This is the reason why BeppoSAX has permitted a great advance in this limited but important field.
+ Unfortunately. even this instrument does not allow the exploration of hard. X.rays xvond. the local Universe. and the cosmological evolution of the column density and covering factor of the absorber. which are important in order to understand the growth of he black holes and its relation with the star formation rate (c.g. Fabian 1999). is still unknown.," Unfortunately, even this instrument does not allow the exploration of hard X–rays beyond the local Universe, and the cosmological evolution of the column density and covering factor of the absorber, which are important in order to understand the growth of the black holes and its relation with the star formation rate (e.g. Fabian 1999), is still unknown."
+ Moreover. only a small raction of the extragalactic sky has been covered so far at hese energies with sullicient sensitivity. which implies that many sources like NGC 4945 and Cireinus are still awaiting discovery.," Moreover, only a small fraction of the extragalactic sky has been covered so far at these energies with sufficient sensitivity, which implies that many sources like NGC 4945 and Circinus are still awaiting discovery."
+ To make significant progresses in this field. a large improvement in sensitivity (like that will be provided. by N*)). and laree area. deep surveys (like that provided by and. even better. that proposed with the project) are needed.," To make significant progresses in this field, a large improvement in sensitivity (like that will be provided by ), and large area, deep surveys (like that provided by and, even better, that proposed with the project) are needed."
+ We thank Enzo Branchini. Roberto Maiolino and Alessandro Marconi for useful discussions.," We thank Enzo Branchini, Roberto Maiolino and Alessandro Marconi for useful discussions."
+ GAL acknowledges ASL and, GM acknowledges ASI and
+this component are much lower than in the disc.,this component are much lower than in the disc.
+ Were such a strict separation between disc and bulge in the models present also in actual galaxies. the bulge metallicity would not be casily observed.," Were such a strict separation between disc and bulge in the models present also in actual galaxies, the bulge metallicity would not be easily observed."
+ This seems to be at variance with actual observations (see.forexample.ο) that do not detect à drop in the rest-frame UV flux in the central regions of the galaxies.," This seems to be at variance with actual observations \citep[see, for example,][]{law2} that do not detect a drop in the rest-frame UV flux in the central regions of the galaxies."
+ This [act may hint that the discburst/bulge decomposition of model galaxies that well characterises local galaxies is not à good approximation of the complex structure of high redshift star forming svstenis., This fact may hint that the disc/burst/bulge decomposition of model galaxies that well characterises local galaxies is not a good approximation of the complex structure of high redshift star forming systems.
+ Modifications such that. for instance. the metal rich eas from the bulge component could enrich the disces. may help in reconciling such models with At the same time. as it can clearly be perceived. from he fact that there are many more galaxies in the z=2 igures as opposed to the z—3 ones. we predict a large number of new galaxies that appear during the elapsed ime interval.," Modifications such that, for instance, the metal rich gas from the bulge component could enrich the discs, may help in reconciling such models with At the same time, as it can clearly be perceived from the fact that there are many more galaxies in the z=2 figures as opposed to the z=3 ones, we predict a large number of new galaxies that appear during the elapsed time interval."
+ Phe voungest ones have stellar masses below 102A1. and low metallicities. thus biasing the predicted mass-metallicity relation to lower values ancl allecting the oedieted: evolution.," The youngest ones have stellar masses below $10^{10} M_{\odot}$ and low metallicities, thus biasing the predicted mass-metallicity relation to lower values and affecting the predicted evolution."
+ Therefore. in order to improve the agreement with observations. one can apply a dkRage cut o the predicted sample of galaxies.," Therefore, in order to improve the agreement with observations, one can apply a SFR/age cut to the predicted sample of galaxies."
+ Given the difficulties in oedicting the correct SElts - see above - this has not been attempted., Given the difficulties in predicting the correct SFRs - see above - this has not been attempted.
+ A closer inspection of our galaxies show that the metallicity increases with decreasing gas mass [fraction almost as expected in a simple closed. box model of chemical evolution. namely ZxdntMarfMges).," A closer inspection of our galaxies show that the metallicity increases with decreasing gas mass fraction almost as expected in a simple closed box model of chemical evolution, namely $Z\propto ln(M_{tot}/M_{gas})$."
+ The gas fraction is almost constant. (around 0.8) with total barvonic mass (at a fixed. redshift} in the high mass regime. whereas it exhibits a large scatter at the low-mass end.," The gas fraction is almost constant (around 0.8) with total baryonic mass (at a fixed redshift) in the high mass regime, whereas it exhibits a large scatter at the low-mass end."
+ On the other and it strongly decreases with stellar mass. although with increasing scatter.," On the other hand it strongly decreases with stellar mass, although with increasing scatter."
+ “This is because the SER is constant or galaxies with stellar mass below 101A1... becoming oportional to the stellar mass above this limit (ef," This is because the SFR is constant for galaxies with stellar mass below $10^{10} M_{\odot}$, becoming proportional to the stellar mass above this limit (c.f."
+ Fig., Fig.
+ 9 in ?))., 9 in \cite{Mann}) ).
+ Only for the most massive galaxies (in terms of their stellar mass). does the gas fraction rise again to (approximately) 0.8.," Only for the most massive galaxies (in terms of their stellar mass), does the gas fraction rise again to (approximately) $0.8$."
+ This is basically what we see in the niass-metallicity plots: the metallicity increases with stellar mass up to 1027A7. where it then Hlattens out in the most massive galaxies (where the gas fraction is higher again), This is basically what we see in the mass-metallicity plots: the metallicity increases with stellar mass up to $\sim 10^{10} M_{\odot}$ where it then flattens out in the most massive galaxies (where the gas fraction is higher again).
+ As explained above. in these systems the star formation is not ellicient. enough. compared to the inflow of primordial eas.," As explained above, in these systems the star formation is not efficient enough compared to the inflow of primordial gas."
+ Indeed. as shown by ?.. the high mass end. of the Z—3 mass-metallicity relation cannot be reproduced by the ποσα disc. but. proto-spheroids with very high SER. such hat their O abundance quickly jumps to solar values in a ew Myr. are needed.," Indeed, as shown by \citet{calura}, the high mass end of the z=3 mass-metallicity relation cannot be reproduced by the “local” disc, but proto-spheroids with very high SFR, such that their O abundance quickly jumps to solar values in a few Myr, are needed."
+ At variance with such a model. where he galaxy morphology is assumed. galaxies may iàve three components (disc. bulge and burst. see Section 2 and ?.. for more details) that co-exist at the same time and whose presence is linked to the evolutionary path of he galaxies.," At variance with such a model, where the galaxy morphology is assumed, galaxies may have three components (disc, bulge and burst, see Section 2 and \citet{galics1}, for more details) that co-exist at the same time and whose presence is linked to the evolutionary path of the galaxies."
+ Here we remind the reader that we show he abunclances predicted. for the disc component.1 of each galaxy only in order to have a fair comparison with the observations that sample quite a large region compared. to he assumed sizes of both the bulge and. components in the model (whenever they are present)., Here we remind the reader that we show the abundances predicted for the disc component of each galaxy only in order to have a fair comparison with the observations that sample quite a large region compared to the assumed sizes of both the bulge and components in the model (whenever they are present).
+ The results do not change if we show the abundances averaged over the hree components because the mass in the cold phase of the disc is much larger than that in the bulge (where we preclict eas mass fraction well below 0.1) or in the burst., The results do not change if we show the abundances averaged over the three components because the mass in the cold phase of the disc is much larger than that in the bulge (where we predict gas mass fraction well below 0.1) or in the burst.
+ 1n these alter components we predict 12|log(O/1)>»8.5 in the majority of the galaxies. however. such a high O abundance is diluted in the mass averaging.," In these latter components we predict $12+\log\left(O/H\right)>8.5$ in the majority of the galaxies, however, such a high O abundance is diluted in the mass averaging."
+ We have already noted wt observations sample the inner regions of the ealaxies., We have already noted that observations sample the inner regions of the galaxies.
+ Jn the other hand. we present the results pertaining to je entire cise because the nominal spatial resolution of is ~30 kpc.," On the other hand, we present the results pertaining to the entire disc because the nominal spatial resolution of is $\sim$ 30 kpc."
+ HE we had enough spatial resolution to make predictions about the inner 4 kpe only. it is likely jit: d) the density would make the SER: higher. rence the metal enrichment quicker and ii) the of 1e metals in the bulge and in the burst by means of the inner disc gas would be lower (meaning a higher metallicity in that region of the disc. lower fractional contribution in 1ο mass averaging).," If we had enough spatial resolution to make predictions about the inner $\sim$ 4 kpc only, it is likely that: i) the density would make the SFR higher, hence the metal enrichment quicker and ii) the of the metals in the bulge and in the burst by means of the inner disc gas would be lower (meaning a higher metallicity in that region of the disc, lower fractional contribution in the mass averaging)."
+ Finally. we note that on the abscissa of all the Figs.," Finally, we note that on the abscissa of all the Figs."
+ discussed so far. we plot the total stellar mass. summed over the three component.," discussed so far, we plot the total stellar mass, summed over the three component."
+ Since bulges aud oursts tend. to be more frequent at the high mass end. where galaxies are older and had more time to evolve and merge. this implies that we the mass axis without an increase in the O abundance. hence obtaining an artificially latter (by à small amount) and more extended. to. high masses mass-metallicity relation than the one expected. for a pure One possible way out would involve making predictions that take into account the three component. without computing an average O abundance.," Since bulges and bursts tend to be more frequent at the high mass end, where galaxies are older and had more time to evolve and merge, this implies that we the mass axis without an increase in the O abundance, hence obtaining an artificially flatter (by a small amount) and more extended to high masses mass-metallicity relation than the one expected for a pure One possible way out would involve making predictions that take into account the three component, without computing an average O abundance."
+ For instance. one could. use the discs to explain the low mass end of the mass metallicity relation. whereas bursts and bulges could explain the high mass end (seealso2)..," For instance, one could use the discs to explain the low mass end of the mass metallicity relation, whereas bursts and bulges could explain the high mass end \citep[see also][]{calura}."
+ As noted above. the bulges cannot be directly. compared: with the observations because itis assumed that the SER in this component can involve only eas returned. from stars. hence the predicted SER. are lower than those in the disc.," As noted above, the bulges cannot be directly compared with the observations because it is assumed that the SFR in this component can involve only gas returned from stars, hence the predicted SFR are lower than those in the disc."
+ Using the bulges instead of the discs at the high mass end. where they are more abundant. will lead to an increase in the slope of the predicted mass-metallicity relation. at the expense of an even poorer agreement between the observed and. predicted At z—2. nearly of the galaxies. feature a burst component.," Using the bulges instead of the discs at the high mass end, where they are more abundant, will lead to an increase in the slope of the predicted mass-metallicity relation, at the expense of an even poorer agreement between the observed and predicted At z=2, nearly of the galaxies feature a burst component."
+ In such a component. the galaxies exhibit SFRs comparable to. or even a factor of ~2 higher than. the disc and the predicted O abundances are S«12|log(OfH)9. so one may be tempted to use only the predicted properties of the. only to compare with observations.," In such a component, the galaxies exhibit SFRs comparable to, or even a factor of $\sim$ 2 higher than, the disc and the predicted O abundances are $8<12+\log\left(O/H\right)<9$, so one may be tempted to use only the predicted properties of the only to compare with observations."
+ The problems with this [ast scenario are numerous., The problems with this last scenario are numerous.
+ In the first. place. the burst. component is assumed. to have a radius that is typically below 1. kpc. hence smaller than both that of the disc and the aperture used. by observers.," In the first place, the burst component is assumed to have a radius that is typically below 1 kpc, hence smaller than both that of the disc and the aperture used by observers."
+ We find it. dillicult to explain the properties of observed galaxies only by means of such a centrally concentrated component., We find it difficult to explain the properties of observed galaxies only by means of such a centrally concentrated component.
+ Moreover. one has to devise a reason to bias the observations in favour of a “burst” a given mass (101A7. ) in order to steepen the mass-metallicity relation at the high mass end.," Moreover, one has to devise a reason to bias the observations in favour of a “burst” a given mass $\sim 10^{10}M_{\odot}$ ) in order to steepen the mass-metallicity relation at the high mass end."
+ In part. such a bias can be granted by the fact. that observed. samples are selected," In part, such a bias can be granted by the fact that observed samples are selected"
+Discovered by Ruprecht on photographie plates as he was searching for open clusters (Alteretal.1961).. Ruprecht 106 existed in relative obscurity until the first CCD photometry by Buonannoetal.(1990). showed that it is metal-poor with a red IB ancl. thus. was likely to be vounger (han other GCs with similar metallicities.,"Discovered by Ruprecht on photographic plates as he was searching for open clusters \citep{al61}, Ruprecht 106 existed in relative obscurity until the first CCD photometry by \citet{bu90} showed that it is metal-poor with a red HB and, thus, was likely to be younger than other GCs with similar metallicities."
+ Subsequentwork by (1993) corroborated this result., Subsequentwork by \citet{bu93} corroborated this result.
+ The metal abundance of Ruprecht 106 has been a source of controversy since the work of Buonannoetal.(1990)., The metal abundance of Ruprecht 106 has been a source of controversy since the work of \citet{bu90}.
+. The abundanuce studies performed (hus far have fallen into two broad categories. those that find a metal abundance of |Fe/II] ~—1.9 (Duonannoetal.1990.Sarajedini&Lavden1997) based on photometric indicators and those that conclude Fe/H] ~—1.6 based on spectroscopy (DaCostaοἱal.1992:Francoiset1997).," The abundance studies performed thus far have fallen into two broad categories, those that find a metal abundance of [Fe/H] $\sim -1.9$ \citep{bu90,bu93,sl97} based on photometric indicators and those that conclude [Fe/H] $\sim -1.6$ based on spectroscopy \citep{da92,fr97}."
+. etal.(1995). found that Ruprecht 106 and M33 have similar metallicities based on time series photometry of 12 RI Lyrae variables and. separately. that [Fe/H] >—1.6 based on the relative locations of the RGB bump and the HL.," \citet{ka95} found that Ruprecht 106 and 3 have similar metallicities based on time series photometry of 12 RR Lyrae variables and, separately, that [Fe/H] $\geq -1.6$ based on the relative locations of the RGB bump and the HB."
+ The discrepancy in metallicity estimates may be due to the unusually low [o /Fe| ratio of Ruprecht 106 compared to other metal-poor Galactic GCs (Pritzl.Venn.&Irwin2005)., The discrepancy in metallicity estimates may be due to the unusually low $\alpha$ /Fe] ratio of Ruprecht 106 compared to other metal-poor Galactic GCs \citep{pvi05}.
+. Brownetal.(1997). claimed [Fe/IJ=~—1.45and [O/Fe] ~0 based on high-resolution spectra of 2 red giants., \citet{bw97} claimed $\simeq -1.45$and [O/Fe] $\sim 0$ based on high-resolution spectra of 2 red giants.
+ VLT/UVES spectroscopy of 6 red. giants indicates —1.5« ΕΠΗ) «—1.45 and —0.1< la/Fe] <0.1 FFrancois. 2010. private comnmnication).," VLT/UVES spectroscopy of 6 red giants indicates $-1.5 < $ [Fe/H] $ < -1.45$ and $-0.1 < $ $\alpha$ /Fe] $< +0.1$ Francois, 2010, private communication)."
+ The earliest CCD photometry of Palomar 15 dates back to the work of aud Harris(1991)., The earliest CCD photometry of Palomar 15 dates back to the work of \citet{sc90} and \citet{ha91}.
+. CMDs presented in both of these studies extend [vom the Gp of the RGB to about 3 mags below the HB and reveal an HD. morphology that is of the RR Lyrae instability strip., CMDs presented in both of these studies extend from the tip of the RGB to about 3 mags below the HB and reveal an HB morphology that is blue-ward of the RR Lyrae instability strip.
+ In addition. both studies conclude that Palomar 15 sulfers Irom a higher-than-expected amount of line-ol-sight reddening.," In addition, both studies conclude that Palomar 15 suffers from a higher-than-expected amount of line-of-sight reddening."
+ Whereas the &Ileiles(1982) reddening maps predict E(D—V) ~ 0.1. the CMDs of and ILarris(1991) suggest a reddening closer to E(B—V) ~ 0.4. consistent with the reddening maps of Schlegeletal. (1993).," Whereas the \citet{bh82} reddening maps predict $-$ V) $\sim$ 0.1, the CMDs of \citet{sc90} and \citet{ha91} suggest a reddening closer to $-$ V) $\sim$ 0.4, consistent with the reddening maps of \citet{sc98}. ."
+. Given this value. photometric indicators predict |Fe/H] ~ -1.9.," Given this value, photometric indicators predict [Fe/H] $\sim$ -1.9."
+ The only spectroscopic value of the metallicity was published by (1995)... who found [Fe/II]2 —2.00=0.08 based on the strength of its Calcium triplet lines.," The only spectroscopic value of the metallicity was published by \citet{da95}, , who found [Fe/H]= $-2.00\pm0.08$ based on the strength of its Calcium triplet lines."
+"x, implies a lower density for that core and hence less central condensation and a larger Love number, as Fig.","$x_\mathrm{c}$ implies a lower density for that core and hence less central condensation and a larger Love number, as Fig."
+ 1. shows., \ref{fig:2L_massratio_k2} shows.
+ The data points for the smallest mass ratios have been generated with the maximum value of q=π., The data points for the smallest mass ratios have been generated with the maximum value of $q=\pi$.
+ This means that for a model with a fixed core radius there is a minimum core mass., This means that for a two-layer model with a fixed core radius there is a minimum core mass.
+" These calculations within a simple two-layer model give an intuitive interpretation of the characteristic of ky of being a measure for the level of central condensation: the bigger the core mass, the more centrally condensed and hence the smaller the Love number ky."," These calculations within a simple two-layer model give an intuitive interpretation of the characteristic of $k_2$ of being a measure for the level of central condensation: the bigger the core mass, the more centrally condensed and hence the smaller the Love number $k_2$."
+" However, as we will show in the next section this simple deduction is no longer possible when there is another density discontinuity in the envelope of the planet."," However, as we will show in the next section this simple deduction is no longer possible when there is another density discontinuity in the envelope of the planet."
+ In this section we will investigate the behavior of kz in a more complicated model., In this section we will investigate the behavior of $k_2$ in a more complicated model.
+ In addition to the density discontinuity at the core-mantle boundary we introduce another discontinuity in the envelope of the planet., In addition to the density discontinuity at the core-mantle boundary we introduce another discontinuity in the envelope of the planet.
+" A three layer structure is a common assumption in planet modeling and has been used for modeling the solar system giants Jupiter and Saturn (see e.g. ?,, ?))."," A three layer structure is a common assumption in planet modeling and has been used for modeling the solar system giants Jupiter and Saturn (see e.g. \citet{Guillot99}, \citet{SaumonGuillot04}) )."
+ Such a separation of layers in the planetary envelope can occur as a result of demixing of hydrogen and helium (?).., Such a separation of layers in the planetary envelope can occur as a result of demixing of hydrogen and helium \citep{StevensonSalpeter77}.
+ It could also arise from double diffusive convection in the presence of a density gradient that is introduced during accretion or because of subsequent core erosion (?).., It could also arise from double diffusive convection in the presence of a density gradient that is introduced during accretion or because of subsequent core erosion \citep{Stevenson82}. .
+ Layer boundaries are also compatible with standard models of planet formation (?).., Layer boundaries are also compatible with standard models of planet formation \citep{Hubbardetal95}.
+ We define our theoretical three-layer model as follows: It consists of a core with constant density A and two polytropic envelopes described by ρι(α) (outer envelope) and p2(x) (inner envelope)., We define our theoretical three-layer model as follows: It consists of a core with constant density $A$ and two polytropic envelopes described by $\rho_1(x)$ (outer envelope) and $\rho_2(x)$ (inner envelope).
+ The same characteristics as for the two-layer model apply., The same characteristics as for the two-layer model apply.
+ The layer boundary in the envelope is placed at x., The layer boundary in the envelope is placed at $x_\mathrm{m}$.
+" We choose the parameters q, and 42 (near but smaller than z and not much different from each other).", We choose the parameters $q_1$ and $q_2$ (near but smaller than $\pi$ and not much different from each other).
+" The location of the density discontinuities x, and xm are also free parameters.", The location of the density discontinuities $x_\mathrm{c}$ and $x_\mathrm{m}$ are also free parameters.
+ The core density A is determined as for the two-layer model., The core density $A$ is determined as for the two-layer model.
+" The continuity of pressure and gravity at Xm is used to calculate x,.", The continuity of pressure and gravity at $x_\mathrm{m}$ is used to calculate $x_\mathrm{a}$.
+" We can then choose B to get a specified non-dimensionalized size of the density jump Ap=(02— p1)/Pily=x,,- ", We can then choose $B$ to get a specified non-dimensionalized size of the density jump $\Delta\rho=\left.(\rho_2-\rho_1)/\rho_1\right|_{x=x_\mathrm{m}}$ .
+"Summarizing, in this three-layer model we can vary the parameters xc, Xm, G1, 42 and Ap while A, B and xa are determined by the parameters chosen."," Summarizing, in this three-layer model we can vary the parameters $x_\mathrm{c}$, $x_\mathrm{m}$, $q_1$, $q_2$ and $\Delta\rho$ while $A$, $B$ and $x_\mathrm{a}$ are determined by the parameters chosen."
+" The parameters characterizing the density discontinuity in the envelope are xm and Ap, position and size of the discontinuity."," The parameters characterizing the density discontinuity in the envelope are $x_\mathrm{m}$ and $\Delta\rho$, position and size of the discontinuity."
+" Thus, in order to investigate the influence of the outer discontinuity on the Love number kz we varied the parameters Xm and Ap from 0.5 to 0.9 planet radii (increment 0.01) and from 0.01 to 0.5 (increment 0.01), respectively."," Thus, in order to investigate the influence of the outer discontinuity on the Love number $k_2$ we varied the parameters $x_\mathrm{m}$ and $\Delta\rho$ from 0.5 to 0.9 planet radii (increment 0.01) and from 0.01 to 0.5 (increment 0.01), respectively."
+" For the example we give here, the other free parameters are fixed at αι= 0.982, qo=0.90π and x,=0.1."," For the example we give here, the other free parameters are fixed at $q_1=0.98\pi$ , $q_2=0.99\pi$ and $x_\mathrm{c}=0.1$."
+ This choice is in order to keep the envelope structure close to an n=1 polytrope., This choice is in order to keep the envelope structure close to an $n=1$ polytrope.
+ Together with a moderate value for the core size this example for the density distribution mimics a Jupiter-like planet., Together with a moderate value for the core size this example for the density distribution mimics a Jupiter-like planet.
+ The change of the Love number kz in dependence on the parameters of the outer density discontinuity is shown by Fig. 2.., The change of the Love number $k_2$ in dependence on the parameters of the outer density discontinuity is shown by Fig. \ref{fig:3L-degeneracy}.
+ Lines of equal kz and their values are given., Lines of equal $k_2$ and their values are given.
+ This demonstrates adiscontinuity.. One can always find many different x-Ap-pairs (that means different three-layer planetary models) that give the same kz., This demonstrates a. One can always find many different $x_\mathrm{m}$ $\Delta\rho$ -pairs (that means different three-layer planetary models) that give the same $k_2$.
+ These models lie on one of the equi-k2-lines and hence one cannot distinguish between these models by a measurement of ky., These models lie on one of the $k_2$ -lines and hence one cannot distinguish between these models by a measurement of $k_2$.
+ In addition to the Love number {2 we calculated the ratio of core mass to total mass Meore/Mitota in the same parameter space., In addition to the Love number $k_2$ we calculated the ratio of core mass to total mass $M_{\mathrm{core}}/M_{\mathrm{total}}$ in the same parameter space.
+ Lines of equal mass ratio are also shown in Fig. 2.., Lines of equal mass ratio are also shown in Fig. \ref{fig:3L-degeneracy}.
+" For a specific value of ky planetary models with different mass ratios are possible, see intersections of equi-k»- and equi-Meore/Miotai- ("," For a specific value of $k_2$ planetary models with different mass ratios are possible, see intersections of $k_2$ - and $M_{\mathrm{core}}/M_{\mathrm{total}}$ -lines. ("
+"Note that the resulting Λήκοιε/Μιοιι have a very limited range because only x, and Ap change and Meore/Mtotalis more sensitive to the q parameters which are fixed in our example here.",Note that the resulting $M_{\mathrm{core}}/M_{\mathrm{total}}$ have a very limited range because only $x_\mathrm{m}$ and $\Delta\rho$ change and $M_{\mathrm{core}}/M_{\mathrm{total}}$is more sensitive to the $q$ parameters which are fixed in our example here.
+" Anyhow, other choices of the q parameters and/or x; would yield qualitativelysimilar results.)"," Anyhow, other choices of the $q$ parameters and/or $x_\mathrm{c}$ would yield qualitativelysimilar results.)"
+" In the same manner,"," In the same manner,"
+significant detection. but the upper limits on the presence of such QPOs are not very. stringent.,"significant detection, but the upper limits on the presence of such QPOs are not very stringent."
+ Although in. some sources stronger high-frequency QPOs have been found (e.g.. NTE J1859|226: Cui et al.," Although in some sources stronger high-frequency QPOs have been found (e.g., XTE J1859+226; Cui et al."
+ 2000). in other sources thehigh-[requeney. QPOs were considerable weaker than our upper limits (e.g. NPE J1550564: Homan et al.," 2000), in other sources thehigh-frequency QPOs were considerable weaker than our upper limits (e.g., XTE J1550–564; Homan et al."
+ 2001)., 2001).
+ Therefore. we cannot exclude that high frequcney QPOs were present in GRs 17392178.," Therefore, we cannot exclude that high frequency QPOs were present in GRS 1739–278."
+ This work was supported by NASA through Chandra Postdoctoral Fellowship grant number PE9-10010 awarded by οΝο which is operated. by SAO) for NASA. under contract NASS-39073.," This work was supported by NASA through Chandra Postdoctoral Fellowship grant number PF9-10010 awarded by CXC, which is operated by SAO for NASA under contract NAS8-39073."
+ “Phis research. has made use of data obtained through the LUEASARC Online Service. provided by the NASA/GSEC and quick-look results provided by the RATE team.," This research has made use of data obtained through the HEASARC Online Service, provided by the NASA/GSFC and quick-look results provided by the team."
+4.,4.
+ THEORETICAL CONSIDERATIONS The discovery rate of black hole transients per vear is now about 5-6 from he nearly continuous coverage of all-sky monitors such as BATSE and he ASM., THEORETICAL CONSIDERATIONS The discovery rate of black hole transients per year is now about 5-6 from the nearly continuous coverage of all-sky monitors such as BATSE and the ASM.
+ Two of the more recently discovered. transients. Cl Cam and NTE J1550-564. have bright (~100 mJ) radio Very.," Two of the more recently discovered transients, CI Cam and XTE J1550-564, have bright $\sim$ 100 mJy) radio counterparts."
+ long xweline interferometry indicates that the radiationcounterparts. comes from clouds of racdio-emitting plasma associated: with jets., Very long baseline interferometry indicates that the radiation comes from clouds of radio-emitting plasma associated with jets.
+ It is therefore possible that jets are common black hole transients. the persistent outbursting and jetamong production in GRS 1915|105although may be rare.," It is therefore possible that jets are common among black hole transients, although the persistent outbursting and jet production in GRS 1915+105 may be rare."
+ Alost black hole transients. such as CRO J1655-40. tend to be active in he GllIz band over the first few weeks or months of the outburst. if at all.," Most black hole transients, such as GRO J1655-40, tend to be active in the GHz band over the first few weeks or months of the outburst, if at all."
+ This associates relatively large accretion rates with jet formation. typical of the high to very high states.," This associates relatively large accretion rates with jet formation, typical of the high to very high states."
+ Phere is also evidence that substates produce jets. as in ο 1915]105. when the sourceparticular is highly variable. and undergoing drastic changes in the broadband:," There is also evidence that particular substates produce jets, as in GRS 1915+105, when the source is highly variable, and undergoing drastic changes in the broadband spectrum."
+ changes in the inner disk structure may producespectrum. a highly RapidIuminous clisk in soft N-ravs. or a fast outflowalternately of high energy. particles.," Rapid changes in the inner disk structure may alternately produce a highly luminous disk in soft X-rays, or a fast outflow of high energy particles."
+ For now. we do not know the origin of such substates.," For now, we do not know the origin of such substates."
+ Likely explanations with those of quasi-periodie oscillations in the accretion. disk. which overlapare bevond the scope of this paper (see contributions by Stella and van der Ixlis. this meeting).," Likely explanations overlap with those of quasi-periodic oscillations in the accretion disk, which are beyond the scope of this paper (see contributions by Stella and van der Klis, this meeting)."
+ work on the ancl of jets has come Considerablefrom theoreticalobservations of ACN overformation many vears., Considerable theoretical work on the formation and collimation of jets has come from observations of AGN over many years.
+collimation Lt is thought that a magnetic field is generated within the nucleus by ordered: motion of accreting and dust., It is thought that a magnetic field is generated within the nucleus by ordered motion of accreting gas and dust.
+ Vhe disk magnetic field is likely to play a role in gas material out of the disk as well as the strongoutflow Phis lau, The disk magnetic field is likely to play a strong role in launching material out of the disk as well as collimating the outflow \cite{Blan82}.
+nchingmodel. in its various forms. is attractivecollimating because magnetically-driven23].. scale to any size system. from. binarics to AGN.," This model, in its various forms, is attractive because magnetically-driven jets scale easily to any size system from binaries to AGN."
+ Reeent work jetswith three easilydimensional simulations an explanation of the transientmagnet natureohyvdrodvnamical of the jets via instabilities in providesmagnetic pressure dominated (low 37) disks 25] or “switches” from the variations in coronal particle density 26]..," Recent work with three dimensional magnetohydrodynamical simulations provides an explanation of the transient nature of the jets via instabilities in magnetic pressure dominated (low $\beta$ ) disks \cite{Mat97} or magnetic “switches"" from the variations in coronal particle density \cite{Me97}."
+ magnetic there is no way to directly measure the strength. of such Currentlyfields near black holes., Currently there is no way to directly measure the strength of such magnetic fields near galactic black holes.
+ Such. fields are inferred. from. observationsmagnetic of voung stellar galacticobjects and the central regions of quasars., Such fields are inferred from observations of young stellar objects and the central regions of quasars.
+ A weakness of magnetic field models is that provide no explanation for the of in some systems., A weakness of magnetic field models is that they provide no explanation for the presence of jets in only some systems.
+ theyVarious other observations in the presenceand infraredjets onlysuggest the presence of massive companions with mass opticalloss rates (such as in a brief of evolution) that results in highconditions that are conducive to jet. production.period similar to (νο N-3 and SS lt 433.," Various other observations in the optical and infrared suggest the presence of massive companions with high mass loss rates (such as in a brief period of evolution) that results in conditions that are conducive to jet production, similar to Cyg X-3 and SS 433."
+has also been proposed. as in AGN. that the of the black hole be a critical factor in ," It has also been proposed, as in AGN, that the spin of the black hole may be a critical factor in jet production."
+The angular spinmomentum of the mavslack hole causes the event jethorizon to production.shrink. and thus the inner edge of he can to the ," The angular momentum of the black hole causes the event horizon to shrink, and thus the inner edge of the accretion disk can extend closer to the black hole."
+increase soft accretiondisk extend clos, This can increase soft X-ray production significantly.
+erZhang et al., Zhang et al.
+ black27]. hole.have Thiscan that he superluminalN-rav. productionsources signilicantly.contain black holes with proposedmomentunir near the maximum allowed mayvalue., \cite{Zh97b} have proposed that the superluminal sources may contain black holes with angular momentum near the maximum allowed value.
+ This may account for the angular soft. emission seen at times in the broadband of the GRO largeJ1655-40 rayand GRS 1915|105., This may account for the large soft X-ray emission seen at times in the broadband spectra of the GRO J1655-40 and GRS 1915+105.
+ lt is not clear how the spectraangular momentum of the slack hole is extracted by the a possible mechanism for this was given w Blandford and Znajek jet, It is not clear how the angular momentum of the black hole is extracted by the jet; a possible mechanism for this was given by Blandford and Znajek \cite{Blan77}.
+:This scenario provides an for other black holes (such as €vg 24]..X-1 in the hard/low state and GS explanation2023|25) of not consistent with the lack of ultrasolt components havingand the prominentsteep jetspower law tail extending into the stronggamma ray regime., This scenario provides an explanation for other black holes (such as Cyg X-1 in the hard/low state and GS 2023+25) of not having prominent jets consistent with the lack of strong ultrasoft components and the steep power law tail extending into the gamma ray regime.
+ 5., 5.
+ THE FUTURE We hope to continue investigations of the known superluminal sources.," THE FUTURE We hope to continue investigations of the known superluminal sources,"
+ CJuneauetal. (Forrarese&Merritt2000:al.20y.," \citep{juneau05,bundy06,schawinski06}. \citep{ferrarese00,gebhardt00,tremaine02}."
+ 50% Tozzictal.2006:Pollettaet2008).," $\%$ \citep{comastri01,tozzi06,polletta08}."
+. recessary to sustain the massive star fornation iu their ost ealaxies (IHopkius&Ieruquist2006)., necessary to sustain the massive star formation in their host galaxies \citep{hopkins06}.
+. On the other rand. dust obscuration can arise in a dMfereut way in he unified model of ACN when accrejon disks aud xoad line regions (BLRs) around SMDIs ave viewed hrough dust torus.," On the other hand, dust obscuration can arise in a different way in the unified model of AGN when accretion disks and broad line regions (BLRs) around SMBHs are viewed through dust torus."
+ Tn any case. red. disty Ανν can shed light on the properties aud the evolutjon of the AGN yopulation im general.," In any case, red, dusty AGNs can shed light on the properties and the evolution of the AGN population in general."
+" The phenomenological definition of ""recdusty ACN& is rather broad."," The phenomenological definition of “red, dusty AGNs” is rather broad."
+ It can include ACGNs selected iun a variety of ways such as those selected by very red colors in optical through NIR aud the radio detection (6.8. RoONοὉ πας and JoAN>1.5 mae of the sample in Urrutia et al.," It can include AGNs selected in a variety of ways such as those selected by very red colors in optical through NIR and the radio detection (e.g., $R-K > 5$ mag and $J-K > 1.3$ mag of the sample in Urrutia et al."
+ 2009: also see Cutri et al., 2009; also see Cutri et al.
+ 2001). rec AUR colors (Laev et al.," 2001), red MIR colors (Lacy et al."
+ 20014: Lee et al., 2004; Lee et al.
+ 2008). anc iard A-rvav detectiouns (e... Polleta et al.," 2008), and hard X-ray detections (e.g., Polleta et al."
+ 2007)., 2007).
+ These differcut kiuds of ACNs have a conunon characteristic where their SEDs are red. due to the obscuration of rei light bw the foreground eas aud dust.," These different kinds of AGNs have a common characteristic where their SEDs are red, due to the obscuration of their light by the foreground gas and dust."
+" Ποιος, rey are cousidered to be the intermediate population )etwoeenu the dust-eushirouded star forming galaxies aru 1 unobseured ACNs."," Hence, they are considered to be the intermediate population between the dust-enshrouded star forming galaxies and the unobscured AGNs."
+ The dust-obscuration does. not OEecessarily exclude Type-l ACNs (ACNs with broac >nission lines: e.g... Alouso-Terrero et al.," The dust-obscuration does not necessarily exclude Type-1 AGNs (AGNs with broad emission lines; e.g., Alonso-Herrero et al."
+" 2006). as red. dusty ACNs with broad emission lines are found quite Τομ,"," 2006), as red, dusty AGNs with broad emission lines are found quite often."
+ OF 23 X-ray QSOs with IR power-law SED anc spectroscopic redshift identification. Ll are classified. as broad-line AGNs (Szololy et al.," Of 23 X-ray QSOs with IR power-law SED and spectroscopic redshift identification, 14 are classified as broad-line AGNs (Szokoly et al."
+ 2001: Alouso-Terrero et al., 2004; Alonso-Herrero et al.
+ 2006)., 2006).
+ More than of the red ACGNs with the radio and red optical/NIR color sclection are found to be tvpe-l ACN which can have the reddening parameter of 2 or more (Uruttia et al., More than of the red AGNs with the radio and red optical/NIR color selection are found to be type-1 AGN which can have the reddening parameter of 2 or more (Uruttia et al.
+ 2009: Clilaman et al., 2009; Glikman et al.
+ 2007)., 2007).
+ The existence of broad line ACONS among red. dusty ACUONSopens a possibility of studyiug the properties of the," The existence of broad line AGNs among red, dusty AGNsopens a possibility of studying the properties of the"
+more extreme forms for the coronal dissipation fraction f(r).,more extreme forms for the coronal dissipation fraction $f(r)$.
+ In detail. we refit the tPEDISIN model allowing f(r) to have a power-law form. f(r)xrS ," In detail, we refit the tPTDISK model allowing $f(r)$ to have a power-law form, $f(r)\propto r^{-\lambda}$."
+phe &oodness of [it is close to that for the best-fit ΕΟΟΟΙ) model., The goodness of fit is close to that for the best-fit tTORQUED model.
+ However. these fits require A>2 at the confidence level.," However, these fits require $\lambda>2$ at the confidence level."
+ Notingthe trivial fact that f(r) cannot exceed. unity. integration of the coronal dissipation across the disk implies that at most of the dissipated energy can be released in the X-ray corona.," Notingthe trivial fact that $f(r)$ cannot exceed unity, integration of the coronal dissipation across the disk implies that at most of the dissipated energy can be released in the X-ray corona."
+ Again. this violates the constraints on the total energeties of this source by a [actor of 5. even once we include the fact that the instantaneous X-ray [lux drops by a factor of 2 when the source enters the Deep Minimum state.," Again, this violates the constraints on the total energetics of this source by a factor of 5, even once we include the fact that the instantaneous X-ray flux drops by a factor of 2 when the source enters the Deep Minimum state."
+ Finally. we examine the cloubly restricted model PPDISW in which Ay=0 and raa x.," Finally, we examine the doubly restricted model PTDISK in which $\Delta\eta=0$ and $r_{\rm out}\rightarrow \infty$ ."
+ From Table, From Table
+saturation.,saturation.
+ However. as with the case of GI191D2D. we measure a polarization that is iuconsisteut with zero.," However, as with the case of G191B2B, we measure a polarization that is inconsistent with zero."
+" The polavized properties of IID282812 have been Q0=33.8"" and p(0.55j=6.29%(xU.: w Turnsheketal.(1990) and Whittetetal.(1992). respectively."," The polarized properties of HD283812 have been reported to be $p(\sim~0.55)=4\%\pm1\%, \theta=33.8\degr$ and $p(0.55)=6.29\%\pm0.05\%, 
+\theta=32\degr\pm1\degr$ by \citet{turn90} and \citet{whit92} respectively."
+" Whittetetal.(1992) also measures p(2.01)=1.31%+0.07% at ()235°42"",", \citet{whit92} also measures $p(2.04)=1.31\%\pm0.07\%$ at $\theta=35\degr\pm2\degr$.
+ As HD283812 is an extended source. we use the brightest northern component to beein the ceutroid for photometry.," As HD283812 is an extended source, we use the brightest northern component to begin the centroid for photometry."
+ Both epochs of the data for ITD283812 had. exposure times of 5.98&., Both epochs of the data for HD283812 had exposure times of 5.98s.
+ Ditheriug was only applied to the September 1997 epoch data however. where the poiut spacing was laree (5700).," Dithering was only applied to the September 1997 epoch data however, where the point spacing was large 0)."
+ The dithered photometric curves of erowth show a spread comparable to the dithered data in Figure 3.. there is no evidence of persistence.," The dithered photometric curves of growth show a spread comparable to the dithered data in Figure \ref{fig:dither}, there is no evidence of persistence."
+ We find that the dithered data shows a larger degree and orieutation of polarization than expected but the nou-dithered data is eutirelv consistent., We find that the dithered data shows a larger degree and orientation of polarization than expected but the non-dithered data is entirely consistent.
+" Turushekoetal.(1990) and οπια,Elston&Lupie(1992). report IID331591 (BD32|3739) to be au unpolarized standard with p(0.13)=0.01440.03%,", \citet{turn90} and \citet{seandl92} report HD331891 (BD32+3739) to be an unpolarized standard with $p(\sim~0.43)=0.04\%\pm0.03\%$.
+ This standard was also used in deteruinimeg the pre- aud post-NC'S coefficients., This standard was also used in determining the pre- and post-NCS coefficients.
+ Whereas alb exposures were 5.988. all but the 1997 data have been dithered with a 2733 four potut spacing.," Whereas all exposures were 5.98s, all but the 1997 data have been dithered with a 3 four point spacing."
+ There is no evidence of persistence aud none of the data shows au exaggerated spread iu the photometric curves of erowth., There is no evidence of persistence and none of the data shows an exaggerated spread in the photometric curves of growth.
+ Nevertheless. all epochs show non-zero degrees and orieutations of polarization.," Nevertheless, all epochs show non-zero degrees and orientations of polarization."
+ These results are consistent with the fudines of Ueta.Miwakawa& who use the same data with the same co-efücieuts to find p(2.05)=Lil’., These results are consistent with the findings of \citet{umandm05} who use the same data with the same co-efficients to find $p(2.05)=1.4\%$.
+ They attribute the LM polarization. ice. greater than the =1% iustrunenutal2 polarization reported by Mines.Scluinidt&Schucider (2000).. to “systematics in the data reduction procedure.”," They attribute the $\gtrsim1\%$ polarization, i.e., greater than the $\lesssim1\%$ instrumental polarization reported by \citet{hsands00}, to “systematics in the data reduction procedure.”"
+ However. the estimate of pinsX (Was based upon ground based-thermal vacuum tests aud not from on-orbit data.," However, the estimate of $p_{ins}\lesssim1\%$ was based upon ground based-thermal vacuum tests and not from on-orbit data."
+ As demonstrated. by Figure 2. we can see that the simulated (Figure 1)) aud real data compare well., As demonstrated by Figure \ref{fig:rescomp} we can see that the simulated (Figure \ref{fig:tests}) ) and real data compare well.
+ In all cases the large-scale values of p and 0 are clearly stable., In all cases the large-scale values of $p$ and $\theta$ are clearly stable.
+ The iuuer variations due to the residual PSF and sub-pixel mis-aligunieut are also seen. but are iusignificant outside of a radius which approximately correspouds to the first airv ring of the PSF (~073). i6. as long as the photometric apertures contain all of the flux spread by the PSF. consistent values of p aud 0 will be measured.," The inner variations due to the residual PSF and sub-pixel mis-alignment are also seen, but are insignificant outside of a radius which approximately corresponds to the first airy ring of the PSF $\sim0\farcs3$ ), i.e., as long as the photometric apertures contain all of the flux spread by the PSF, consistent values of $p$ and $\theta$ will be measured."
+ It is clear that there are siguificaut differences in polarizations measured at separate epochs for CILTÀ-DC-F7 and IID2835812., It is clear that there are significant differences in polarizations measured at separate epochs for CHA-DC-F7 and HD283812.
+ It is uulikelv that these variations are intrinsic to the objects., It is unlikely that these variations are intrinsic to the objects.
+ Iu both cases it is noted that the data are pre-NCS aud dithered., In both cases it is noted that the data are pre-NCS and dithered.
+ It is also unclear why such an anomaly is not seen in the pre-NC'S dithered data for IID3231891., It is also unclear why such an anomaly is not seen in the pre-NCS dithered data for HD331891.
+ However. as sueecsted in refehadefrz.. this may be an artifact from the pedestal effect.," However, as suggested in \\ref{chadcf7}, this may be an artifact from the pedestal effect."
+ There do exist several IRAF software packages o remove the pedestal effect(pedshky. podsub). but as he dither allows a check of the NICMOS polarization xoperties im cach quadraut. we cau check to see if this is the reason for the discrepant results ly verformineg our routine only ou the data from a single »outius.," There do exist several IRAF software packages to remove the pedestal effect, ), but as the dither allows a check of the NICMOS polarization properties in each quadrant, we can check to see if this is the reason for the discrepant results by performing our routine only on the data from a single pointing."
+ Accordinely we have re-creduced the raw data files through the pipeline aud determined the volavization results from cach quadrant., Accordingly we have re-reduced the raw data files through the pipeline and determined the polarization results from each quadrant.
+ The results are xeseuted in Table L., The results are presented in Table \ref{tab:pedres}.
+ It can be scen that the individual »omtiues eive polarization nieasures consistent with unu-dithered data. aud that the pre-NCS pedestal effect can induce variations of Ap~0.2% iud. AOz15? between different quadrauts.," It can be seen that the individual pointings give polarization measures consistent with un-dithered data, and that the pre-NCS pedestal effect can induce variations of $\Delta{p}\approx0.2\%$ and $\Delta\theta\approx15\degr$ between different quadrants."
+ Post-NCS the pedestal effect. does not appear to have affected the results. however. it is clear that the pedestal effect does produce slisbtlv different polarizations iu cach quadraut of the array.," Post-NCS the pedestal effect does not appear to have affected the results, however, it is clear that the pedestal effect does produce slightly different polarizations in each quadrant of the array."
+ We have measured polarization in all of the targets which have previously been listed as munpolarized., We have measured polarization in all of the targets which have previously been listed as unpolarized.
+ Sucht findings sugeest that there may be a low level of iustrmuental polarization., Such findings suggest that there may be a low level of instrumental polarization.
+ From Table 3 it cau be secu that. on average. this instrumental polarization has a," From Table \ref{tab:results} it can be seen that, on average, this instrumental polarization has a"
+lower boundary (?7).,lower boundary \citep{voegleretal2005ii}.
+ The value of the cutropy density has been chosen such that the resulting mean radiative enerev flix in the nonanagnetic run matches the observed solar value to withini, The value of the entropy density has been chosen such that the resulting mean radiative energy flux in the non-magnetic run matches the observed solar value to within.
+"τοι, In all three simmlations. the total radiative out] at the top of the simulation box fluctuates by about 4] around its mean value in the course of the 5-nüuute oscillations iu the box."," In all three simulations, the total radiative output at the top of the simulation box fluctuates by about $\pm 1$ around its mean value in the course of the 5-minute oscillations in the box."
+ Each sauulation was runi or about 150 nünu solar time after a statistically stationary state of the magueto-convective system hack evolved., Each simulation was run for about 150 min solar time after a statistically stationary state of the magneto-convective system had evolved.
+ For the detailed analysis we then considered 10 equidistant snapshots covering about 20 iin in order to suppress the effects of the 5baninute oscillations and to reduce the statistical fluctuatious., For the detailed analysis we then considered 10 equidistant snapshots covering about 20 min in order to suppress the effects of the 5-minute oscillations and to reduce the statistical fluctuations.
+ 2?) used these siuulatious to study the ceutre-to-liuh variation of the bolometric intensity iu dependence of the average vertical maenetic field., \cite{voegleretal2005ii} used these simulations to study the centre-to-limb variation of the bolometric intensity in dependence of the average vertical magnetic field.
+ We calculate the emcergeut iuteusities along parallel ravs. cach passing through a different poiut of the surface erid of the MIID suuulatious.," We calculate the emergent intensities along parallel rays, each passing through a different point of the surface grid of the MHD simulations."
+ Since we are interested in a broad spectral coverage. we use Isurucz spectral svuthesis code ATLAS (7). (which iucludes diatomic molecules) with opacity distribution functions (ODFs) on cach of those simulated model atinosphleres at disk ceutre and towards the limb.," Since we are interested in a broad spectral coverage, we use Kurucz' spectral synthesis code ATLAS9 \citep{kurucz1993} (which includes diatomic molecules) with opacity distribution functions (ODFs) on each of those simulated model atmospheres at disk centre and towards the limb."
+ ATLASS (as rewritten by J. B. Lester) aud the ODFs were obtained through CCP? (Collaborative Computing Project No. 7+))., ATLAS9 (as rewritten by J. B. Lester) and the ODFs were obtained through CCP7 (Collaborative Computing Project No. ).
+ The program solves the radiative trauster equatious under the asstuuption of local theriodyuanuie equilibria (LTE)., The program solves the radiative transfer equations under the assumption of local thermodynamic equilibrium (LTE).
+ The simplification of LTE aud the use of ODFs allows us to calculate the whole spectrum from 160 mu to 160 000. um with a spectral resolution of better than 200 in the visible., The simplification of LTE and the use of ODFs allows us to calculate the whole spectrum from 160 nm to 160 000 nm with a spectral resolution of better than 200 in the visible.
+ Towever. the TINODE/SOT/BFI filters are warower than the ATLASS resolution bius. and. in addition. the spectra are not calculated at the wavelength ceutre for the filters which leads to a comparison at slightly different waveleugths.," However, the HINODE/SOT/BFI filters are narrower than the ATLAS9 resolution bins, and, in addition, the spectra are not calculated at the wavelength centre for the filters which leads to a comparison at slightly different wavelengths."
+ Furthermore. the CN and CD regions comprise niuiv spectral lines that are ouly approximated by coarsely resolved ODFs in the ATLAS9 calculation.," Furthermore, the CN and GB regions comprise many spectral lines that are only approximated by coarsely resolved ODFs in the ATLAS9 calculation."
+ This may influence the comparison of the simulated intensities with the observations for the CN aud GB wavelengths., This may influence the comparison of the simulated intensities with the observations for the CN and GB wavelengths.
+ For an inclined view away frou disk ceutre. we allow the line of sight to cut through the simulated atinospliere at an angle 0 to the surface normal (0 is the heliocentric auele).," For an inclined view away from disk centre, we allow the line of sight to cut through the simulated atmosphere at an angle $\theta$ to the surface normal $\theta$ is the heliocentric angle)."
+ For our sinulatiou box. this means we lave to interpolate the atiiosphiere parameters along ravs with the respective heloccutric angle.," For our simulation box, this means we have to interpolate the atmosphere parameters along rays with the respective heliocentric angle."
+ For the comparison with observations. we degraded the simulated images to take iuto account optical effects: The poiut-spread fictions for the SOT as calculated by 2) were used for both dataset I and dataset IT. as discussed in Sect.," For the comparison with observations, we degraded the simulated images to take into account optical effects: The point-spread functions for the SOT as calculated by \cite{mathewetal2009} were used for both dataset I and dataset II, as discussed in Sect."
+ 2., 2.
+ Tere. we compare the resulting intensities from slinulations with observations to validate the sinulated results.," Here, we compare the resulting intensities from simulations with observations to validate the simulated results."
+ For dataset I at j=0.9 with its finer binnine. we first inspect the observed aud simulated intensity nuages visually for a qualitative comparison.," For dataset I at $\mu=0.9$ with its finer binning, we first inspect the observed and simulated intensity images visually for a qualitative comparison."
+ Inu a next step we compare dataset II with the simulations (at all linib angles) iu terius of ruis contrasts and intensity cistributions., In a next step we compare dataset II with the simulations (at all limb angles) in terms of rms contrasts and intensity distributions.
+ Iu this section. we compare observations with original aud convolved simulations at a limb angle of j(/20.9.," In this section, we compare observations with original and convolved simulations at a limb angle of $\mu=0.9$."
+ Iu Fig., In Fig.
+ 3 original aud convolvec simulation snapshots for OC are shown together with the observed image iu the respective wavelength filter., \ref{fig:muram50G_hinode_unconv_conv_all} original and convolved simulation snapshots for 50G are shown together with the observed image in the respective wavelength filter.
+ The qualitative agreement is good. the granules are of comparable size aud the visually oereeived. contrasts appear to be simular.," The qualitative agreement is good, the granules are of comparable size and the visually perceived contrasts appear to be similar."
+ The occurrence of brighter aud darker features is similar aud also changes correspoudingly for the differcut wavelength filters., The occurrence of brighter and darker features is similar and also changes correspondingly for the different wavelength filters.
+ The xiehtest features can be seen in the CN wavelength., The brightest features can be seen in the CN wavelength.
+ The xieht maenetic features that appear for the Ὀθέα (as seen in Fie. 3)), The bright magnetic features that appear for the 50G (as seen in Fig. \ref{fig:muram50G_hinode_unconv_conv_all}) )
+ and 2006 siuulations are uot seen for 1ionnmaenuetie simulations., and 200G simulations are not seen for nonmagnetic simulations.
+ Iu Fie., In Fig.
+ |. we demoustrate the effect of convolving with the PSF in a more quantitative wav., \ref{fig:hist_gc_unconv} we demonstrate the effect of convolving with the PSF in a more quantitative way.
+ We plot. as an example. the listograms for the ereen contimmuni intensity distribution in the original simulation at the liuib angle jp.=0.9 for OG. 50G. and 200€. We overplot (n erev) the corresponding histogram for the convolved station as shown in Fieure 5.," We plot, as an example, the histograms for the green continuum intensity distribution in the original simulation at the limb angle $\mu=0.9$ for 0G, 50G, and 200G. We overplot (in grey) the corresponding histogram for the convolved simulation as shown in Figure \ref{fig:muram_hinode_conv_histo_all}."
+ As expected. the convolved histograms are significantly narrower.," As expected, the convolved histograms are significantly narrower."
+ For most wavelengths we find a distribution with a double peal iu the uncouvolved simulations. which is most prominent iu the 50C simulation. aud represeuts the dark intereranular conmonents (the more prominucut peak at 2/2) <1) and the bright eranular compoucuts (the weaker peak. at {ο >1) (κου e.g. ?) and references therein).," For most wavelengths we find a distribution with a double peak in the unconvolved simulations, which is most prominent in the 50G simulation, and represents the dark intergranular components (the more prominent peak at $I/\langle I \rangle $ $<$ 1) and the bright granular components (the weaker peak at $I/\langle I \rangle $ $>$ 1) (see e.g. \cite{wedemeyervoort2009} and references therein)."
+ After he convolution. the double-pea is smeared out. but inportant differences )etyeen the nommaenetic. 506. and 200€; siuulations remain. iu particular for the visible xoacbau istributious (BC. GC. aud RC).," After the convolution, the double-peak is smeared out, but important differences between the nonmagnetic, 50G, and 200G simulations remain, in particular for the visible broadband distributions (BC, GC, and RC)."
+ Fiewre 5- shows the histograms of the iutensitv for (all) 10 averaged simulation snapshots and for average naenetic fields of OC. 50G. and 200€. From left to right we plot listograims for the convolved simulations at the hree limb angles jp=0.9. ," Figure \ref{fig:muram_hinode_conv_histo_all} shows the histograms of the intensity for (all) 10 averaged simulation snapshots and for average magnetic fields of 0G, 50G, and 200G. From left to right we plot histograms for the convolved simulations at the three limb angles $\mu=0.9$ "
+iav be lessened or completely suppressed. (Slilosiuau Noguchi. 1993).,"may be lessened or completely suppressed (Shlosman Noguchi, 1993)."
+ We have eucounutered this problem when evolving models with which prevented us from completing these rus., We have encountered this problem when evolving models with which prevented us from completing these runs.
+ Thestellar bar strength has been quautified using the Fourier amplitude of the m=2 mode normalized by the a=0 mode (Paper D., The bar strength has been quantified using the Fourier amplitude of the $m = 2$ mode normalized by the $m = 0$ mode (Paper I).
+ It is obtaimed by integration over restricted. evliudrical volumes where the bar is a dominant morphological feature. namely. over R=01 ranee. where is the bar size defined oei Section 23.3. (see also Paper Dj.," It is obtained by integration over restricted cylindrical volumes where the bar is a dominant morphological feature, namely, over $R = 0.1 - $ range, where is the bar size defined in Section 3.3 (see also Paper I)."
+ In Figure 1.. we plot as a funetion of time f for various gas fractious. and eravitational softenius in the gas.ea.," In Figure \ref{A2b}, we plot as a function of time $t$ for various gas fractions, and gravitational softening in the gas,."
+ The initial stages in the evolution of are very similar in all inodels: an initial stage of an accelerated erowth. and a peak followed by. a sudden drop.," The initial stages in the evolution of are very similar in all models: an initial stage of an accelerated growth, and a peak followed by a sudden drop."
+ The duration of this dyaiizuuical stage varies from model to model. but is conipleted by £f100.," The duration of this dynamical stage varies from model to model, but is completed by $t\sim 100$."
+ The first peak in is followed by the vertical buckling instability in the bar leading to an abrupt weakening in the bar but not a complete dissolution (c.e.. Magztinez-Valpuesta Shlosiman 2001).," The first peak in is followed by the vertical buckling instability in the bar leading to an abrupt weakening in the bar but not a complete dissolution (e.g., Martinez-Valpuesta Shlosman 2004)."
+" We observe that the bar weakening is quite iudependoeut of audενος, ax expected."," We observe that the bar weakening is quite independent of and, as expected."
+ For gas-poor models. the first peak of is iudependoenut of andfy.," For gas-poor models, the first peak of is independent of and."
+".. For eas-vich models. with fee απ, the peak is lowered gradually for essays0.05. up to a factor of 0.1."," For gas-rich models, with $\gtorder 15\%$ , the peak is lowered gradually for $\ltorder 0.05$, up to a factor of 0.4."
+ Models with έρως0.1 appear not to be affected at all by this treud., Models with $=0.1$ appear not to be affected at all by this trend.
+ The post-buckliug evolution of the bars is much inore affected by audeui. and shows a bimodal behavior.," The post-buckling evolution of the bars is much more affected by and, and shows a bimodal behavior."
+ Naiuely. dm sas-poor models. the bar resumes its erowth but at amore eradual pace as compared to the dynamical erowth.," Namely, in gas-poor models, the bar resumes its growth but at a more gradual pace as compared to the dynamical growth."
+ In eas-vich models. the bar strength declines over the simulation (i.c.. Hubble) tine.," In gas-rich models, the bar strength declines over the simulation (i.e., Hubble) time."
+ This οςαἱ behavior was noted by Bereutzen et al. (, This bimodal behavior was noted by Berentzen et al. (
+2007). but the current models show it more explicitly.,"2007), but the current models show it more explicitly."
+ We shall refer to the first bar evolution phase. iucludiug the buckling. as thedynamical plase. aud to the subsequent evolution as thesecular phase.," We shall refer to the first bar evolution phase, including the buckling, as the phase, and to the subsequent evolution as the phase."
+ The secular stage allows the separation of the secululv erowing models from the secularly declining oues., The secular stage allows the separation of the secularly growing models from the secularly declining ones.
+ Ta no models have the bars disappeared completely at least a substantial oval distortion remained., In no models have the bars disappeared completely — at least a substantial oval distortion remained.
+ The borderline between these two trends. erowthn and dechne. depends ou eaa.," The borderline between these two trends, growth and decline, depends on ."
+ For 0.016. Hiesintherangoofte~ — n—," For $=0.016$, it lies in the range of $\sim 5\%-7\%$."
+ Larger Inoves the borderline toward iiore gas-rich models., Larger moves the borderline toward more gas-rich models.
+" Models SD.GG5S82. and SD.GCGI5S3— ο close το this borderline. for104 for —0.4050ndf,— —1012% for mU.dl. andshowaccrysingularecolution afterthesuddendropAs,~coust."," Models G8S2 and G15S3 lie close to this borderline, $\sim 8\%-10\%$ for $=0.05$ and $\sim 10\%-12\%$ for $=0.1$, and show a very singular evolution --- after the sudden drop $\sim $ const."
+ for about. Af~100. and a second period of bar growth beeims.," for about $\Delta t\sim 100$, and a second period of bar growth begins."
+ This is a surprisingly iixed behavior showing the prolonged constancy in of the eas-ricli models and the secular bar growth of the gas-poor oucs., This is a surprisingly mixed behavior showing the prolonged constancy in of the gas-rich models and the secular bar growth of the gas-poor ones.
+ Model SD_GG50S2 has a peculiar evolution iu the dynamical stage that deserves special attention., Model G50S2 has a peculiar evolution in the dynamical stage that deserves special attention.
+ Eveu though this is a gas-rich model. the bar resumes is secular erowthn after the buckling.," Even though this is a gas-rich model, the bar resumes is secular growth after the buckling."
+ The rate of growth is much more nioderate than that iu the other models but is clearly noticeable., The rate of growth is much more moderate than that in the other models but is clearly noticeable.
+ The evolution of the pattern speed is shown iu Fig. 2.., The evolution of the pattern speed is shown in Fig. \ref{figps}.
+ It remains about constant during the dynamical phase of the bar evolution. slightly rising in the gas rich models.," It remains about constant during the dynamical phase of the bar evolution, slightly rising in the gas rich models."
+ The bimodality of gas-poor aud rich models shows up in the secular evolution where drops in the former aud stavs constant or rises in the latter disks., The bimodality of gas-poor and rich models shows up in the secular evolution where drops in the former and stays constant or rises in the latter disks.
+ Iu the secular phase. nearly always auticorrelates witli A sustained drop iu corresponds to a rise ind," In the secular phase, nearly always anticorrelates with A sustained drop in corresponds to a rise in."
+ons Thus. in the gas-poor models the bar tumbling slows down. while in the σαςΙσ models it stavs nearly constant.," Thus, in the gas-poor models, the bar tumbling slows down, while in the gas-rich models it stays nearly constant."
+ Fie., Fig.
+ 2 provides a lint to what oue can expect for the evolution of the augular momentum im the cisk-halo system. although in no wav does this figure account forall the angular momentum in the disk. aud even iu the bar region it represeuts oulv the tumbling of the bar.," 2 provides a hint to what one can expect for the evolution of the angular momentum in the disk-halo system, although in no way does this figure account for the angular momentum in the disk, and even in the bar region it represents only the tumbling of the bar."
+ The eas-poor disks lose their tumbling angular mioimenutuni cficicuthy. while eas-rich ones nearly conserve it or even speed up their tumbling divine the secular phase.," The gas-poor disks lose their tumbling angular momentum efficiently, while gas-rich ones nearly conserve it or even speed up their tumbling during the secular phase."
+ The leugth of the bar semiauajor axis.Πρι las been taken as the radius where the bar equatorial ellipticity drops bv off its peak (Paper D.," The length of the bar semi-major axis, has been taken as the radius where the bar equatorial ellipticity drops by off its peak (Paper I)."
+ This method has been tested iu comparison with alternative method basec ou the last stable orbit supporting the bar (Martinez-Valpuesta ot al., This method has been tested in comparison with alternative method based on the last stable orbit supporting the bar (Martinez-Valpuesta et al.
+ 2006)., 2006).
+ It is the most reliable when applied after the first imiaxiuuni of the bar strength., It is the most reliable when applied after the first maximum of the bar strength.
+ The cllipticity of the bar at different radii is obtaine by fitting ellipses to the isodensity contours in the face-on disk., The ellipticity of the bar at different radii is obtained by fitting ellipses to the isodensity contours in the face-on disk.
+ As iu the collisionless models. the bar leneth exhibits the initial period. of an accelerated erowthiux reaches adnasim which always coincides (iu finie) with the maxinuun in (Fig. 3)).," As in the collisionless models, the bar length exhibits the initial period of an accelerated growthand reaches amaximum which always coincides (in time) with the maximum in (Fig. \ref{figrbar}) )."
+ This svinholizes the, This symbolizes the
+It seems reasonable to begin with the assumption that the OI emission οι IRC+10216 is intvinsically svnuuetric.,It seems reasonable to begin with the assumption that the OH emission from IRC+10216 is intrinsically symmetric.
+ If that were the case. (hen in order to explain the OIL line shapes. ere must exist some absorbing medium within the circumstellar envelope of IRC+10216 which prevents emission from (he lar (recdshifted) side of the envelope from reaching us. while enussion from the near (blueshifted) side escapes and is detected.," If that were the case, then in order to explain the OH line shapes, there must exist some absorbing medium within the circumstellar envelope of IRC+10216 which prevents emission from the far (redshifted) side of the envelope from reaching us, while emission from the near (blueshifted) side escapes and is detected."
+ Since these observations were done al wavelengths which are large. relative to previous observations of this star. (here is one source of opacity Chat could have previously gone undetected - electron Iree-[ree absorption.," Since these observations were done at wavelengths which are large, relative to previous observations of this star, there is one source of opacity that could have previously gone undetected - electron free-free absorption."
+ However. (his absorption would be accompanied by thermal continuum enission.," However, this absorption would be accompanied by thermal continuum emission."
+ In order to be an effective absorber of the redshiltecl half of the OI line. the absorber would need to be optically thick at a radius of ~1011i em. where the OI distribution peaks (hough 1e peak radius of the OLI distribution varies somewhat depending on the parameters of the photodissociation model. LOM em is a reasonable value [for most models).," In order to be an effective absorber of the redshifted half of the OH line, the absorber would need to be optically thick at a radius of $\sim 10^{17}\,$ cm, where the OH distribution peaks (though the peak radius of the OH distribution varies somewhat depending on the parameters of the photodissociation model, $10^{17}\,$ cm is a reasonable value for most models)."
+ At such a large radius. (he envelope eas temperature is quite cold. about 111Ix. and the thermal emission from the absorber would correspond to a blackbody at (he same temperature.," At such a large radius, the envelope gas temperature is quite cold, about $11\,$ K, and the thermal emission from the absorber would correspond to a blackbody at the same temperature."
+ Unfortunately. our observations are not well suited to determining ihe continuum brightness at 1665 and 1667 MlIz. but Sahai. Claussen. Masson (1989) have observed the continuum brightness ab slightly. higher frequencies. specifically al 15 and 20 Gllz.," Unfortunately, our observations are not well suited to determining the continuum brightness at $1665$ and $1667\,$ MHz, but Sahai, Claussen, Masson (1989) have observed the continuum brightness at slightly higher frequencies, specifically at $15$ and $20\,$ GHz."
+ The predicted. επικ densities due to [ree-[ree emission at 15 and 20 GlIz are 250 and 140 mJ. respectively. for an object which is optically thick at 1665 and. 1667 MIIz.," The predicted flux densities due to free-free emission at $15$ and $20\,$ GHz are $250$ and $140\,$ mJy, respectively, for an object which is optically thick at $1665$ and $1667\,$ MHz."
+ Sahai et al. (, Sahai et al. (
+1989) find the [αν clensities to be only 6 and 1.43mJy at 15 and 20 GlIz. respectively.,"1989) find the flux densities to be only $6$ and $1.4\,$ mJy at $15$ and $20\,$ GHz, respectively."
+ The absence of strong continuum enission indicates that at these Irequencies. (he envelope of IRC+10216 cannot be optically thick at radii large compared to that of the AO beam. effectively ruling out the absorption scenario.," The absence of strong continuum emission indicates that at these frequencies, the envelope of IRC+10216 cannot be optically thick at radii large compared to that of the AO beam, effectively ruling out the absorption scenario."
+ In addition. the electron density required. (o make the envelope of IRC+10216 oplically thick would be substantially in excess of the densities presently expected based on theoretical models of (he envelope chemistry.," In addition, the electron density required to make the envelope of IRC+10216 optically thick would be substantially in excess of the densities presently expected based on theoretical models of the envelope chemistry."
+ A second possible explanation for the observed line profiles is maser amplification of the 1667 and 1665 MIIz lines.," A second possible explanation for the observed line profiles is maser amplification of the $1667$ and $1665\,$ MHz lines."
+ This possibility is attractive. because it is the only one which could potentially be responsible for a velocity dillerence between (he 1667 and 1665 MIIz lines.," This possibility is attractive, because it is the only one which could potentially be responsible for a velocity difference between the $1667$ and $1665\,$ MHz lines."
+ While the modeling of circumstellar maser emission is difficult and a detailed model of possible OIL masers around IC4-10216 is bevond the scope of this paper. we can compare our observations to other asvanptotie giant branch (AGB) stars with known OIL maser emission.," While the modeling of circumstellar maser emission is difficult and a detailed model of possible OH masers around IRC+10216 is beyond the scope of this paper, we can compare our observations to other asymptotic giant branch (AGB) stars with known OH maser emission."
+ llowever. we should be cautious about drawing strong conclusions based on (hiis comparison.," However, we should be cautious about drawing strong conclusions based on this comparison,"
+and standard I-band magnitudes.,and standard I-band magnitudes.
+ We conservatively estimate the error of calibration to be 0.1 mag., We conservatively estimate the error of calibration to be 0.1 mag.
+ We looked for the host galaxy of by fitting two-dimensional surface brightness models to the observed 1-band image., We looked for the host galaxy of by fitting two-dimensional surface brightness models to the observed i-band image.
+ We used three different models: an unresolved nucleus only. nucleus + a de Vaucouleurs profile host galaxy (Sérrsic ϱ = 0.25) and nucleus + a disk host galaxy (6 = 1.0).," We used three different models: an unresolved nucleus only, nucleus + a de Vaucouleurs profile host galaxy (Sérrsic $\beta$ = 0.25) and nucleus + a disk host galaxy $\beta$ = 1.0)."
+ Details of this fitting process can be found in ?.., Details of this fitting process can be found in \cite{1999PASP..111.1223N}.
+" In short. the model has three adjustable parameters. the magnitude of the nucleus Jeon. the magnitude of the host galaxy ος, and the effective radius of the host galaxy ray."," In short, the model has three adjustable parameters, the magnitude of the nucleus $I_{\rm core}$, the magnitude of the host galaxy $I_{\rm
+ host}$ and the effective radius of the host galaxy $r_{\rm eff}$."
+ The unresolved nucleus is assumed to be centered on the host galaxy., The unresolved nucleus is assumed to be centered on the host galaxy.
+ The three parameters are adjusted via an iterative Levenberg-Marquardt loop until the minimum chi squared between the model and the observed surface brightness distribution 1s found., The three parameters are adjusted via an iterative Levenberg-Marquardt loop until the minimum chi squared between the model and the observed surface brightness distribution is found.
+ The fit was extended to an outer radius of 9755. except that the pixels affected by a galaxy 5799 SE of were masked out from the fit.," The fit was extended to an outer radius of 5, except that the pixels affected by a galaxy 9 SE of were masked out from the fit."
+ Prior to computing the chi squared the model was convolved with the PSF., Prior to computing the chi squared the model was convolved with the PSF.
+ We used star 1 in Fig., We used star 1 in Fig.
+ 1. for the PSF due to its high signal to noise and proximity to 3C 270., \ref{kentta} for the PSF due to its high signal to noise and proximity to 3C 279.
+ As in ?.. we studied the variability of the PSF across the FOV by extracting the surface brightness profiles of stars 1-4 in Fig.," As in \cite{2008A&A...487L..29N}, we studied the variability of the PSF across the FOV by extracting the surface brightness profiles of stars 1-4 in Fig."
+ | and deriving the rms scatter between the profiles., \ref{kentta} and deriving the rms scatter between the profiles.
+ The PSF variability was then parameterized by a parabolic expression where csi) 1s the uncertainty in the PSF relative to the intensity at radius r (pixels) from the PSF center., The PSF variability was then parameterized by a parabolic expression where $\sigma_{\rm PSF}(r)$ is the uncertainty in the PSF relative to the intensity at radius $r$ (pixels) from the PSF center.
+ Wethen computed the expected variance c in a pixel with intensity 7 (ADU) and distance r from the center of from the expression where G is the effective gain and R ts the effective readout noise. and compute the y- from where M is the model intensity and the summation is over all unmasked pixels within the fitting radius.," Wethen computed the expected variance $\sigma^2$ in a pixel with intensity $I$ (ADU) and distance $r$ from the center of from the expression where $G$ is the effective gain and $R$ is the effective readout noise, and compute the $\chi^2$ from where $M$ is the model intensity and the summation is over all unmasked pixels within the fitting radius."
+ By including the PSF variability into the computation of y we can ensure that the results are not dominated by PSF errors. which are most pronounced close to the center of the object.," By including the PSF variability into the computation of $\chi^2$ we can ensure that the results are not dominated by PSF errors, which are most pronounced close to the center of the object."
+ The results of this model fitting are summarized in Table and Fig. 2.., The results of this model fitting are summarized in Table \ref{tulokset} and Fig. \ref{profiilit}.
+ We see a clear excess in 3C 279 over the PSF and this excess is clearly above any PSF variability across the FOV (see the lower panel of Fig. 2))., We see a clear excess in 3C 279 over the PSF and this excess is clearly above any PSF variability across the FOV (see the lower panel of Fig. \ref{profiilit}) ).
+ The models with a host galaxy give a much better fit than a model with pure nucleus with y indicating a near-perfect fit., The models with a host galaxy give a much better fit than a model with pure nucleus with $\chi^2$ indicating a near-perfect fit.
+ The model with ade Vaucouleurs host galaxy gives a slightly better fit that the model with a disk galaxy but the difference is not significant according to our error simulations (see below)., The model with a de Vaucouleurs host galaxy gives a slightly better fit that the model with a disk galaxy but the difference is not significant according to our error simulations (see below).
+ However. since the de Vaucouleurs profile gave formally a better fit and no blazar host galaxy has ever been associated with a disk galaxy. we concentrate in the following to the results with the de Vaucouleurs profile.," However, since the de Vaucouleurs profile gave formally a better fit and no blazar host galaxy has ever been associated with a disk galaxy, we concentrate in the following to the results with the de Vaucouleurs profile."
+ We have studied the sensitivity of the results to random noise. PSF variability and errors m the assumed profile of the host galaxy by performing 50 model fits to simulated images of279.," We have studied the sensitivity of the results to random noise, PSF variability and errors in the assumed profile of the host galaxy by performing 50 model fits to simulated images of."
+. In these simulations Jose. μου and ray were held constant at the values shown in Table .. but the contribution from the different noise sources changed from one simulation to another.," In these simulations $I_{\rm core}$, $I_{\rm host}$ and $r_{\rm eff}$ were held constant at the values shown in Table \ref{tulokset}, but the contribution from the different noise sources changed from one simulation to another."
+ The random noise in these simulations 1s assumed to raise from readout and photon noise according to the first term on the right hand side of Eq. 2.., The random noise in these simulations is assumed to raise from readout and photon noise according to the first term on the right hand side of Eq. \ref{nkaava}.
+ To model the PSF variability we created two PSF models in each simulation. slightly differing from each other.," To model the PSF variability we created two PSF models in each simulation, slightly differing from each other."
+ The simulated model was convolved with the first PSF and the model fits were made with the second PSF., The simulated model was convolved with the first PSF and the model fits were made with the second PSF.
+ Both PSFs consisted of an elliptical Moffat profile with 5 = 2.5., Both PSFs consisted of an elliptical Moffat profile with $\beta$ = 2.5.
+ The ellipticity of the first PSF was randomly drawn from an uniform distribution with a minimum of 0.0 and à maximum of 0.12 and the position angle from an uniform distribution between 0 and 180 degrees., The ellipticity of the first PSF was randomly drawn from an uniform distribution with a minimum of 0.0 and a maximum of 0.12 and the position angle from an uniform distribution between 0 and 180 degrees.
+ For the second PSF the ellipticity was drawn from a Gaussian distribution with a mean equal to the ellipticity of the first PSF and σ= 0.05., For the second PSF the ellipticity was drawn from a Gaussian distribution with a mean equal to the ellipticity of the first PSF and $\sigma = 0.05$ .
+ The position angle of the second PSF was similarly drawn from a Gaussian distribution with a mean equal to the first PSF and c=12.5 , The position angle of the second PSF was similarly drawn from a Gaussian distribution with a mean equal to the first PSF and $\sigma = 12.5$ 
+»»Utiue centers are shown in Fie. l..,pointing centers are shown in Fig. \ref{f5:poi}.
+ In four poiutiugs Πω...↽∖l.57.. which. means that offsets could amount to 15 of the signal.," In four pointings $\la
+1.8$, which means that offsets could amount to 15 - of the signal."
+" These poiutings are iubers 9. 10. 12 and 2L and the lin those poiutiugs are ~ 1.6. — 1.7. — 1.7 aud ~ 1.3 respectively,"," These pointings are numbers 9, 10, 12 and 24, and the in those pointings are $\sim$ 1.6, $\sim$ 1.7, $\sim$ 1.7 and $\sim$ 1.3 respectively."
+" All aare computed only atf positions where yx πανeam and reduced u of the linear of Αη Vag<2 (soc Sect, 3.5)).", All are computed only at positions where $\langle P \rangle > 20$ mJy/beam and reduced $\chi^2$ of the linear $\phi(\lambda^2)$ -relation $\chi^2_{red} < 2$ (see Sect. \ref{ss5:rm}) ).
+ So. care must be exercised in interpreting the polarization data from these pointing centers.," So, care must be exercised in interpreting the polarization data from these pointing centers."
+ Ou the other haud. polarization angles show a linear variation over frequency in a large part of the data. vielding good RAY determinations. which would not be possible if large-scale offsets dominate.," On the other hand, polarization angles show a linear variation over frequency in a large part of the data, yielding good $RM$ determinations, which would not be possible if large-scale offsets dominate."
+ See Iaverkorn ct ((2003a) for an extended discussion of this point., See Haverkorn et (2003a) for an extended discussion of this point.
+ The observed Stokes Q aud C intensities at the rresolutiou are shown in Fig. 2.., The observed Stokes $Q$ and $U$ intensities at the resolution are shown in Fig. \ref{f5:qu}.
+" The distributions of Q aud U values iu the field are approximately Catssian. centered around zero. aud have a width of 20 to 25 wJv/heam (equivalent to polarized brightness temperatures Tn, of 2.9 - 3.7 EK) for the five frequencies;"," The distributions of $Q$ and $U$ values in the field are approximately Gaussian, centered around zero, and have a width of 20 to 25 mJy/beam (equivalent to polarized brightness temperatures $T_{b,pol}$ of 2.9 - 3.7 K) for the five frequencies."
+ The ring structure is clearly visible in Q and (., The ring structure is clearly visible in $Q$ and $U$.
+" To emphasize the perfect circularity of the ring. a circle of radius aand centered on (0.8)=(18.05°.65.73""). which was fitted bv eve to the ring-like structure in Q and C. is superimposed in both maps."," To emphasize the perfect circularity of the ring, a circle of radius and centered on $(\alpha,\delta) = (48.05\dg,
+65.73\dg)$, which was fitted by eye to the ring-like structure in $Q$ and $U$, is superimposed in both maps."
+ The maps of polarized intensity P derived from the Stokes Q aud © inaps are shown in Fie., The maps of polarized intensity $P$ derived from the Stokes $Q$ and $U$ maps are shown in Fig.
+ 3 for all 5 frequencies., \ref{f5:pi} for all 5 frequencies.
+" White denotes the highest intensity. aud intensities above 110 utybeam (equal to T5,16 I&) ave saturated."," White denotes the highest intensity, and intensities above 110 mJy/beam (equal to $T_{b,pol} = 16$ K) are saturated."
+ The maxiuun inteusitics iu the 5 frequency bands are Lil. 122. 107. 120. and 113 mJv/beanu. respectively.," The maximum intensities in the 5 frequency bands are 141, 122, 107, 120, and 143 mJy/beam, respectively."
+ Superimposed are lues of constaut Calactic latitude basTe . and 117. and the superimposed circle is the same as the one in Fig. 2..," Superimposed are lines of constant Galactic latitude $b
+= $, and , and the superimposed circle is the same as the one in Fig. \ref{f5:qu}."
+ Several regions ofdifferent topology of polarized intensity are present iu the field: Tn the lower melt plot of Fig., Several regions of different topology of polarized intensity are present in the field: In the lower right plot of Fig.
+ 3 a map of total intensity Za 319 MITZz is shown. from which point sources > Foaindv/beam are removed.," \ref{f5:pi} a map of total intensity $I$ at 349 MHz is shown, from which point sources $> 5$ mJy/beam are removed."
+ The map has the same resolution of «5.5/ aand the same brightuess scaling as the P maps iu the other paucls of the same figure., The map has the same resolution of $\times$ and the same brightness scaling as the $P$ maps in the other panels of the same figure.
+ No structure in total intensity 7 is visible. although there is abundant structure iu polarization.," No structure in total intensity $I$ is visible, although there is abundant structure in polarization."
+ The circular structure in- the upper left corner of the £ map is artificial and caused by a very bright unpolarized extragalactic source at that position., The circular structure in the upper left corner of the $I$ map is artificial and caused by a very bright unpolarized extragalactic source at that position.
+ As an iuterferometer acts as a ligh-pass filter. the map imteeral 7 over the field is set to zero bv lack of information about its true level.," As an interferometer acts as a high-pass filter, the map integral $I$ over the field is set to zero by lack of information about its true level."
+ Frou the single dish survev of Hashun et ((1981.. 1982) at los MIIz. the total brightuess temperature at [08 MITz at this position is approximately [1.5 I& with a temperature uncertainty of ~ ," From the single dish survey of Haslam et (1981, 1982) at 408 MHz, the total brightness temperature at 408 MHz at this position is approximately 44.5 K with a temperature uncertainty of $\sim$ "
+The photosphere typically occurs at several scale heights above the midplane (as is also seen iu shearing box simulations with [ar ruore detailed thermocyuaimics: see ?)).,The photosphere typically occurs at several scale heights above the midplane (as is also seen in shearing box simulations with far more detailed thermodynamics; see \cite{HKS06}) ).
+ Since we clefine a different photospliere at each pxint (r.6) in the disk. for each frame of simulation data. the emission and rav-tracing are truly th'ee-cdimensional aud dyvuamic.," Since we define a different photosphere at each point $(r,\phi)$ in the disk, for each frame of simulation data, the emission and ray-tracing are truly three-dimensional and dynamic."
+ Only afterwards do we iutegrate over avinli aud time to produce t1e observed spectra., Only afterwards do we integrate over azimuth and time to produce the observed spectra.
+ Just like a real detector. this dyuzunic. 3-d ray-raciug allows is to accurately uodel the effects of isolated lot-spots aud velocity perturbations iu tje disk that teid to lead to alarder spectrum.," Just like a real detector, this dynamic, 3-d ray-tracing allows us to accurately model the effects of isolated hot-spots and velocity perturbations in the disk that tend to lead to a harder spectrum."
+ At intervals of LOOAL from /10000347 to !1500047. we used the 3-d simulation data as boudary couditious for the general relativistic radiation trausler code «escribed in ? ancl ?..," At intervals of $100M$ from $t=10000M$ to $t=15000M$, we used the 3-d simulation data as boundary conditions for the general relativistic radiation transfer code described in \cite{Schnittman:2009} and \cite{Schnittman:2011}."
+ With that code. for each point ou the plotosphere (rx©=36061) i each. suapshot. we followed ~104 photon packets along outware-directed rays raudomly selectedLin direction with a probability cistributiou uniorum over solid aigle iu the fIuid. frame.," With that code, for each point on the photosphere $r \times \phi
+= 360\times 64$ ) in each snapshot, we followed $\sim 10^4$ photon packets along outward-directed rays randomly selected in direction with a probability distribution uniform over solid angle in the fluid frame."
+ Most plotons reach infinity (here. r= OAL).," Most photons reach infinity (here, $r=10000M$ )."
+ A muority are captured by the black hole., A minority are captured by the black hole.
+ Another ημον strike the accretion disk somewhere ese. where they a'e scattered: with a redistributio1 function following the expression for a scatteriue-douinated atmosphere derived by ?..," Another minority strike the accretion disk somewhere else, where they are scattered with a redistribution function following the expression for a scattering-dominated atmosphere derived by \cite{Chandra:1960}."
+2 To determije the {his directed iu a given solid angle. we clOse 11 bins eveily spaced in cos0 aud grouped a| photous arriving within a sinele bin.," To determine the flux directed in a given solid angle, we chose 41 bins evenly spaced in $\cos\theta$ and grouped all photons arriving within a single bin."
+" When we cie a bolometric Iuminosity d£/dQ in a particular direction. it is computed by integratiug in requency over dL,/dt."," When we cite a bolometric luminosity $dL/d\Omega$ in a particular direction, it is computed by integrating in frequency over $dL_\nu/d\Omega$."
+ Figure 2 displays the tiue- aud aziimiuthially-averagec surface brightjess of ThinHBR. as ineasured in the local fluid. frame (i.e.. the local orbital frame).," Figure \ref{fig:ffflux} displays the time- and azimuthally-averaged surface brightness of ThinHR, as measured in the local fluid frame (i.e., the local orbital frame)."
+ TIie surface brigituness follows the model prediction a r1041 because at sucl radii. the stress al the ISCO has little elffect(—the specilic accretec angul:uw iuonjentunm (the parameter fixec in the NT mocel by the ISCO sress) is stnall compared to t1e local specific anguar monaentuni," The surface brightness follows the Novikov-Thorne model prediction at $r \gtrsim 10M$ because at such radii, the stress at the ISCO has little effect—the specific accreted angular momentum (the parameter fixed in the NT model by the ISCO stress) is small compared to the local specific angular momentum."
+ Near aud iuside the ISCO. i0wever. the surface brightiess contrasts sharply with tle NT model.," Near and inside the ISCO, however, the surface brightness contrasts sharply with the NT model."
+" I remains at ahigh level all he way to the event horizo""", It remains at ahigh level all the way to the event horizon.
+ Our results may also je compared to those of ?.., Our results may also be compared to those of \cite{BHK08}.
+ Scaling dissigation rate to stress. they oedieted fIuid-f[raime cdissip:οι rates that rose steeply through the plungiug region.," Scaling dissipation rate to stress, they predicted fluid-frame dissipation rates that rose steeply through the plunging region."
+ Que reasou our plunging region eimissiviy ds SOLjewliat less than thelr estimate is that our cooling function does ιοί leac to complete radiatii1 of all dissipated heat., One reason our plunging region emissivity is somewhat less than their estimate is that our cooling function does not lead to complete radiation of all dissipated heat.
+ Adiabatic expansion can lower the temperature low the target temperatuο despite continued dissipation., Adiabatic expansion can lower the temperature below the target temperature despite continued dissipation.
+ As a result. some beat is kept in the Iuid all the way to the horizon.," As a result, some heat is kept in the fluid all the way to the horizon."
+ The amount of heat retained is illustrated in Figure 3.. which shows," The amount of heat retained is illustrated in Figure \ref{fig:enthalpy}, , which shows"
+For more thàn a decade. radial velocity observations. with aceuraeles of order mss”! have been within reach (see for instance ? and 2)).,"For more than a decade, radial velocity observations with accuracies of order $^{-1}$ have been within reach (see for instance \citet{marbut2000} and \citet{queloz2001}) )."
+ Even accuracies of less than 1 mss! (2) are possible now.," Even accuracies of less than 1 $^{-1}$ \citep{pepe2003}
+ are possible now."
+ With these observations. more than 200 sub-stellar companions have been discovered by measuring the reflex motions of their parent stars.," With these observations, more than 200 sub-stellar companions have been discovered by measuring the reflex motions of their parent stars."
+ Most of these sub-stellar companions have been detected around F. G and K main sequence stars. but detections around an A star (2) and several subgiants (?.. ?)) have also been reported recently.," Most of these sub-stellar companions have been detected around F, G and K main sequence stars, but detections around an A star \citep{galland2006} and several subgiants \citet{johnson2006}, \citet{johnson2007}) ) have also been reported recently."
+ Moreover. 10 giant stars were reported to have sub-stellar companions (t Draconis (K21ID) ?.. HDI04985 (G9IID ?.. HD47526 (KILL) ?.. HD13189 (K2II-IID ?.. HD11977 (G5IID ?.. Pollux (ΚΟΠΗ 2.. 2.. 4UMa (ΚΤΠ) ?.. NGC2423 No3 and NGC4349 NoI27 ?.. and recently HD17092 (ΚΟΠΟ ?))!..," Moreover, 10 giant stars were reported to have sub-stellar companions $\iota$ Draconis (K2III) \citet{frink2002}, HD104985 (G9III) \citet{sato2003}, HD47526 (K1III) \citet{setiawan2003}, HD13189 (K2II-III) \citet{hatzes2005}, , HD11977 (G5III) \citet{setiawan2005}, Pollux (K0III) \citet{hatzes2006}, \citet{reffert2006}, 4UMa (K1III) \citet{dollinger2007}, NGC2423 No3 and NGC4349 No127 \citet{lovis2007}, and recently HD17092 (K0III) \citet{niedzielski2007})."
+ In addition to searches for extra-solar companions. radial velocity observations prove to be very useful for detecting solar-like oscillations 1n. stars with turbulent atmospheres. such as the dwarf a Cen A (e.g. ?).. the subgiant Procyon (e.g.??) and the giant e Ophiuchi (e.g.2).," In addition to searches for extra-solar companions, radial velocity observations prove to be very useful for detecting solar-like oscillations in stars with turbulent atmospheres, such as the dwarf $\alpha$ Cen A \citep[e.g.][]{bedding2006}, the subgiant Procyon \citep[e.g.][]{eggenberger2004,martic2004} and the giant $\epsilon$ Ophiuchi \citep[e.g.][]{deridder2006}."
+ With techniques for accurate radial velocity observations at hand. à survey was started in 1999 to verify whether K giants are stable enough to be used as astrometric reference stars for SIM/PlanetQuest (Space Interferometry Mission) (?)..," With techniques for accurate radial velocity observations at hand, a survey was started in 1999 to verify whether K giants are stable enough to be used as astrometric reference stars for SIM/PlanetQuest (Space Interferometry Mission) \citep{frink2001}."
+ This survey contains 179 stars and uses the Coudé Auxiliary Telescope (CAT) at University of California Observatories / Lick Observatory. in conjunetion with the Hamilton Echelle Spectrograph.," This survey contains 179 stars and uses the Coudé Auxiliary Telescope (CAT) at University of California Observatories / Lick Observatory, in conjunction with the Hamilton Echelle Spectrograph."
+ The survey has recently been expanded to about 380 giants and ts still ongoing., The survey has recently been expanded to about 380 giants and is still ongoing.
+ For the analysis described in the present paper onlydata from the initial 179 stars are used., For the analysis described in the present paper onlydata from the initial 179 stars are used.
+ From this survey. companions have been announced for t Draconis (2). and Pollux (?)..," From this survey, companions have been announced for $\iota$ Draconis \citep{frink2002} and Pollux \citep{reffert2006}."
+ Stars with radial. velocity variations of less than 20 mss?! have been presented as stable stars by ?.., Stars with radial velocity variations of less than 20 $^{-1}$ have been presented as stable stars by \citet{hekker2006a}.
+ In addition. some binaries discovered with this survey. as well as an extensive overview of the sample. will be presented in forthcoming papers.," In addition, some binaries discovered with this survey, as well as an extensive overview of the sample, will be presented in forthcoming papers."
+ As almost all of the stars show significant radial velocity variations. we investigate here which mechanism causes these variations.," As almost all of the stars show significant radial velocity variations, we investigate here which mechanism causes these variations."
+ Non-periodic radial velocity variations. of the order of the investigated timescales. are most likely caused by some intrinsic mechanism. while the periodic variability can also be caused by companions.," Non-periodic radial velocity variations, of the order of the investigated timescales, are most likely caused by some intrinsic mechanism, while the periodic variability can also be caused by companions."
+ We also investigate the characteristics of these companions., We also investigate the characteristics of these companions.
+ In Sect., In Sect.
+ 2. the radial velocity observations are described in detail.," 2, the radial velocity observations are described in detail."
+ In Sect., In Sect.
+ 3. the relation between the observed radial velocity amplitude and surface gravity is investigated.," 3, the relation between the observed radial velocity amplitude and surface gravity is investigated."
+ In Sect., In Sect.
+ 4. we explore the hypothesis that all periodic radial velocity variations are caused by sub-stellar companions. and we compare the inferred orbital parameters with those obtained for sub-stellar companions orbiting main sequence stars.," 4, we explore the hypothesis that all periodic radial velocity variations are caused by sub-stellar companions, and we compare the inferred orbital parameters with those obtained for sub-stellar companions orbiting main sequence stars."
+ Our conclusions are presented in Sect., Our conclusions are presented in Sect.
+ 5., 5.
+ The initial 179 stars selected for the radial velocity survey are used in the present work., The initial 179 stars selected for the radial velocity survey are used in the present work.
+ These stars have been selected from the Hipparcos catalogue (?).. based on the criteria described by ?.," These stars have been selected from the Hipparcos catalogue \citep{esa1997}, based on the criteria described by \citet{frink2001}."
+ The selected stars are all brighter than 6 mag. are presumably single and have photometric variations <0.06 mag in V.The survey started in 1999 at Lick Observatory using the Coudé Auxiliary Telescope (CAT) in conjunction with the Hamilton Echelle Spectrograph 0000).," The selected stars are all brighter than 6 mag, are presumably single and have photometric variations $< 0.06$ mag in V.The survey started in 1999 at Lick Observatory using the Coudé Auxiliary Telescope (CAT) in conjunction with the Hamilton Echelle Spectrograph 000)."
+ The system with an todine cell in the light path has been described by ? and ?.., The system with an iodine cell in the light path has been described by \citet{marcy1992} and \citet{valenti1995}. .
+" With integration times of up to thirty minutes for the faintest stars (7, = 6 mag) we reacha signal to noise ratio of about 80—100 att=5500 A. yielding à radial velocity precision of"," With integration times of up to thirty minutes for the faintest stars $m_{v}$ = 6 mag) we reacha signal to noise ratio of about $80-100$ at $\lambda = 5500$ , yielding a radial velocity precision of"
+a fold. pair.,a fold pair.
+ To derive this expression. we must first obtain the image positions at which the time delay is evaluated.," To derive this expression, we must first obtain the image positions at which the time delay is evaluated."
+ These results were derived by Keetonοἱal.(2005)..., These results were derived by \citet{Keeton-fold}.
+ We oller a summary of their analysis in Section 3.1. ancl present our new results for the time delay in Section 3.2..," We offer a summary of their analysis in Section \ref{sub:fold_image}
+ and present our new results for the time delay in Section \ref{sub:fold_time}."
+ Since we are considering à source near a fold point. we write its position in terms of a scalar parameter c. which we take to be small but finite.," Since we are considering a source near a fold point, we write its position in terms of a scalar parameter $\epsilon$, which we take to be small but finite."
+ In particular. let uo» eu.," In particular, let $\bu \to \epsilon \bu$ ."
+ Combining equations (7)) and (11)) we can write the lens equationas, Combining equations \ref{eq:lens-ortho}) ) and \ref{eq:psiexp}) ) we can write the lens equationas
+Observations of fifteen quasar candidates were obtained with the GAIOS spectrograph (Alurowinski (2003) in preparation. Hook (2003) submitted) cluring regular queuecscheduled observations on the Gemini North telescope| in. November 2001 and June-July 2002.,"Observations of fifteen quasar candidates were obtained with the GMOS spectrograph (Murowinski (2003) in preparation, Hook (2003) submitted) during regular queue-scheduled observations on the Gemini North telescope in November 2001 and June-July 2002."
+ The November observations were among the first. regular science. observations with GALOS North and. as a result. could. not be taken with he proper blocking filter to remove the second order spectra.," The November observations were among the first regular science observations with GMOS North and, as a result, could not be taken with the proper blocking filter to remove the second order spectra."
+ Conditions were photometric or near-photometric or the majority of the observations and the image quality was better than | aresec., Conditions were photometric or near-photometric for the majority of the observations and the image quality was better than 1 arcsec.
+ AX slit was used with the 1400 grating and the detector was binned 2. vielding a dispersion of 1.4. pixel. and a spectral resolution (determined. by the slit width) of9X.," A slit was used with the R400 grating and the detector was binned $\times$ 2, yielding a dispersion of 1.4 /pixel and a spectral resolution (determined by the slit width) of."
+. Ehe useful spectral range was typically5900-9500., The useful spectral range was typically.
+.. An C515 filter was used o remove the second. order for the 2002. observations., An OG515 filter was used to remove the second order for the 2002 observations.
+ A single exposure was obtained for most objects. with exposure imes ranging from. LSOO0s to 2400s (For the faintest targets).," A single exposure was obtained for most objects, with exposure times ranging from 1800s to 2400s (for the faintest targets)."
+ The data were reduced following standard methods with he Gemini LIUAE. to pre-process and mosaic the spectra from the 3 CCD detectors together. followed by extraction and reduction of he 1D spectra withSPECRED.," The data were reduced following standard methods with the Gemini IRAF to pre-process and mosaic the spectra from the 3 CCD detectors together, followed by extraction and reduction of the 1D spectra with."
+. An approximate lux calibration was «determined. from observations of stancard stars with identical instrument. parameters during he same period of observation (contamination [rom second order blue light further reduced. the accuracy of the [lux calibration for the 2001 November observations)., An approximate flux calibration was determined from observations of standard stars with identical instrument parameters during the same period of observation (contamination from second order blue light further reduced the accuracy of the flux calibration for the 2001 November observations).
+ Details of the fifteen. candidate quasars are given in table 3.., Details of the fifteen candidate quasars are given in table \ref{candidates}.
+ No strong emission line objects are detected., No strong emission line objects are detected.
+ Indeed. all of the observed spectra show strong similarities ancl are suggestive of early type M stars.," Indeed, all of the observed spectra show strong similarities and are suggestive of early type M stars."
+ In order to determine the Classification of stars prone to selection using the V7Z colour diagram a composite stellar spectrum is constructed by co-acleling the individual observations of each candidate (Figure ο.," In order to determine the classification of stars prone to selection using the $VIZ$ colour diagram a composite stellar spectrum is constructed by co-adding the individual observations of each candidate (Figure \ref{all spectra
+combined}) )."
+ The composite spectrum is examined following the classification schemes of Ixirkpatrick. Henry. and. AleCarthy (1991) (for IX-M stars) ancl Wirkpatrick (1999) (for later tvpe L stars)," The composite spectrum is examined following the classification schemes of Kirkpatrick, Henry and McCarthy (1991) (for K-M stars) and Kirkpatrick (1999) (for later type L stars)."
+ The presence of pronounced. absorption features. coincident with the expected location of “TiO bands. suggest identification with stars later than ~AI2.," The presence of pronounced absorption features, coincident with the expected location of TiO bands, suggest identification with stars later than $\sim$ M2."
+ The lack of a broad potassium feature at ~TTOOA lis taken as evidence that the spectra are not from L stars., The lack of a broad potassium feature at $\sim$ is taken as evidence that the spectra are not from L stars.
+ Absorption in VO bands becomes increasingly pronounced in later M stars., Absorption in VO bands becomes increasingly pronounced in later M stars.
+ We therefore identify the composite spectra with a classification earlier than M5., We therefore identify the composite spectra with a classification earlier than M5.
+ Examination of the spectra of cach target individually supports the classification range (M2-5) assigned to the composite spectra., Examination of the spectra of each target individually supports the classification range (M2-5) assigned to the composite spectra.
+ TiO. absorption bands are readily identified., TiO absorption bands are readily identified.
+ While the VO. feature is in a region of the spectrum allected by strong OLL sky residuals. there is little evidence for significant. absorption.," While the VO feature is in a region of the spectrum affected by strong OH sky residuals, there is little evidence for significant absorption."
+ A classification for stellar interlopers in the range M2- is consistent with the svathetic photometry prediction shown in Figures 7 and those of Dobbie.Pinfield. Jameson and Lloclekin (2002).," A classification for stellar interlopers in the range M2-M5 is consistent with the synthetic photometry prediction shown in Figures \ref{low mass stars iz} and those of Dobbie,Pinfield, Jameson and Hodgkin (2002)."
+ A tentative quasar space density of 2.96. 10. 7 ⋅ mM ≼⇍∪⊔↕⊓⇂⋖⊾⊔≼⇍⋖⋅↓↓⊔⊔↿⊳∖∶∶≓↽⊳4610 MESESCy. inferredsop for⋅ the redshift∢⋅ range 48<2.<5.2 and absolute magnitude. Mg«/— 23.5(Vega) limit accessible to the CCELII2lIN. VZZ survey.," A tentative quasar space density of $\times$ $^{-7}$ $^{-3}$ confidence limits $\pm_{7.5\times10^{-9}}^{1.6\times10^{-6}}$ ) is inferred for the redshift range $4.8<z<5.2$ and absolute magnitude, $_B<-23.5$ (Vega) limit accessible to the CFH12K $VIZ$ survey."
+ Vhe survey areas ds. estimated via monte-carlo integration., The survey area is estimated via monte-carlo integration.
+ “The survey completeness function is estimated in the manner adopted by Sharp (2001)., The survey completeness function is estimated in the manner adopted by Sharp (2001).
+ Α eric of moclel quasar spectra is generated. over the accessible redshift range. with parameters for the quasar spectral model. cliscussecl in section J) 3.. crawn at random from the clistribution functions Synthetic photometry estimates are computed for each," A grid of model quasar spectra is generated, over the accessible redshift range, with parameters for the quasar spectral model, discussed in section \ref{colour selection}, drawn at random from the distribution functions Synthetic photometry estimates are computed for each"
+To evaluate the possible role of bulk Doppler shilts for radiation transport in hot Jupiter abmospheres. i( is (hus important to evaluate the shapes and widths of radiative lines in (hese atimospheres.,"To evaluate the possible role of bulk Doppler shifts for radiation transport in hot Jupiter atmospheres, it is thus important to evaluate the shapes and widths of radiative lines in these atmospheres."
+ Various static. globally-averagecl radiative transfer studies have established that the raciativelv forced regions of hot Jupiter atmospheres are found. broadly speaking. above the 10LOO bar pressure level 2005).," Various static, globally-averaged radiative transfer studies have established that the radiatively forced regions of hot Jupiter atmospheres are found, broadly speaking, above the 10–100 bar pressure level ."
+. These models also indicate that photospheric vary strongly with wavelength in the opticalIR spectral range. as a result of strong variations in line opacities. [roni a few bars to <107 bars typically2007).," These models also indicate that photospheric vary strongly with wavelength in the optical–IR spectral range, as a result of strong variations in line opacities, from a few bars to $\la 10^{-2}$ bars typically."
+. Regions in the pressure range from cLO to 10.7 bars are (hus of parücular interest [or (he study of radiation transport in dvnanmic hot Jupiter atmosplieres., Regions in the pressure range from $\sim 10$ to $10^{-3}$ bars are thus of particular interest for the study of radiation transport in dynamic hot Jupiter atmospheres.
+ A detailed account of the typical treatment of lines in static radiative transfer caleulations for hot Jupiter atinospheres is provided by(2007)., A detailed account of the typical treatment of lines in static radiative transfer calculations for hot Jupiter atmospheres is provided by.
+. The general Voigt prolile of a pressure- and Doppler-broadened line is givenbv ∖∖⇁↥∐↲↕⋅≼↲∆↗∣∕↙⇂≀↕↴∐≼⇂∆↗∕⊓∐∐↲≀↧⊍∖⊽⋯⋅≼↲⊔∐↲↕↽≻↕⋅≼↲⋟∖∷∖⇁∏↕⋅≼↲≀↧↴↕∐⇂∐∪↕↽≻↕↽≻↥≼↲↕⋅∣↽≻↕⋅∪≀↧↴≼⇂≼↲∐↕∐≸≟∖∖⊽↕≼∐∐⋟∖⊽⋅↕⋅≼↲⋟∖⊽↕↽≻≼↲≺∢∐∖↽≼↲↥⋡∖↽⋅ ≀↕↴∐≼⇂≺↕≻∖⇁↕⋟∖⊽∐∪↕⋅∐," The general Voigt profile of a pressure- and Doppler-broadened line is givenby where $\Delta \nu_p$ and $\Delta \nu_D$ measure the pressure and Doppler broadening widths, respectively, and $\Phi_V$ is normalized to unity."
+↓≀↧↴∐∠≼↲≺⇂↥∪∏∐∐⋡∖↽≼⋡⊂↽⊐∪⋯⇂∡���↽≪↽∖↽↘⊽∏∐↖⊂↽↔↴↓≤∍⋖↽∖⋟∣≤∍∶↴∏↥∪⊔≀↧↪∖⊽≪↽∖↽⊳↔⊲↥≀↧↴∐∐∐↲⋟∖⊽∃∩∩∃∶⊔∪∏ ⊋∪∩∃↕⋝⋅⋅↕↴↕⋅≼↲⋟∖⊽⋟∖⇁⋯⋅≼↲∣↽≻↕⋅∪≀↧↴≼⇂≼↲∐↕∐≸≟↖≺↽↔↴≼↲∐≼↲↕⋅≀↧↴∐∡∖↽≼⇂≼↲↕↽≻≼↲∐≼⇂⋟∖⊽∪∐↕↽≻↕⋅≼↲⋟∖⊽⋟∖⊽⋯⋅≼↲⋅∩↲⊔↕↽≻≼↲↕⋅≀↧↴⊓∐⋅≼↲≀↧↴∐≼⇂⊔∐↲↕⋅≀↧↴≺∐≀↧↴∐∖⇁↩ ≺∢∪∐⋟∖⊽⋯⋯↲∐↥∏∐≺⇂≼↲↕⋅≺∢∪∐⋟∖⊽↕≼⇂≼↲↕⋅≀↧↴∐∪∐∣↽≻⋯⊔∐↲⋅⋅≺∢↥≀↧⊔∖⊽⋟∖⊽↕≺∢≀↧↴↥⋅⋅⋟∖⊽≺∢≀↧↴∐∐↖≺↽↔↴ with a —1/2 power law dependence on temperature. should be sufficient for our orcler- estimates2007).," Pressure broadening generally depends on pressure, temperature and the radiative constituent under consideration but the “classical” scaling with a $-1/2$ power law dependence on temperature, should be sufficient for our order-of-magnitude estimates."
+. Doppler broadening depends on temperature. wavelength (Ay= οι) ancl the mass. m4. Of the radiative constituent under consideration.," Doppler broadening depends on temperature, wavelength $\lambda_0 = c /
+\nu_0$ ) and the mass, $m_{\rm mol}$ , of the radiative constituent under consideration,"
+"integral 1+ (24Q)e/c.yf= (A +9)/2andi""= (5 +37)/8.","integral where $\eta=1+(2+\zeta)c_g/c$ , $\eta^\prime=(1+\eta)/2$ and $\eta^{\prime\prime}= (5+3\eta)/8$."
+"Wethenintegrate the various terms by parts. using the formula [ (do /o"". !)f(s;)=—f(s;)/no""4-(do/na""df(de."," We then integrate the various terms by parts, using the formula $\int (d\sigma /
+\sigma^{n+1}) f(s_a) = -f(s_a)/n\sigma^n + \int (d\sigma /n\sigma^n)
+df/d\sigma$."
+ Each o-derivative df/do raises the order of that term by €F?7: we stop integrating by parts when the resulting (erm is Ofe) relative to the leading term., Each $\sigma$ -derivative $df/d\sigma$ raises the order of that term by $\epsilon^{1/2}$; we stop integrating by parts when the resulting term is $O(\epsilon)$ relative to the leading term.
+" Because σ in the integral paramelrizes the light path. ancl because we are not Tavlor expanding the orbital variables. we must take into account (hat s, depends on σ through Eq. (25))."," Because $\sigma$ in the integral parametrizes the light path, and because we are not Taylor expanding the orbital variables, we must take into account that $s_a$ depends on $\sigma$ through Eq. \ref{sigmaexpand}) ),"
+ so that After integration. we discard terms that depend on negative powers of the emission time σι: lor light[rom a distant quasar. these are clearly negligible.," so that After integration, we discard terms that depend on negative powers of the emission time $\sigma_e$; for lightfrom a distant quasar, these are clearly negligible."
+ The result is, The result is
+"""Tomorrow. a gamma ray burst (GRB) may be observed that originates from the death of one of the first stars. during he epoch of reionisation.","Tomorrow, a gamma ray burst (GRB) may be observed that originates from the death of one of the first stars, during the epoch of reionisation."
+ Despite the great distance to this »urst. it will be the brightest. gamma αν source on the sky or several tens of seconds. one of the brightest cosmological A-ray sources for hours. and its afterglow will be observable for weeks in the near-infrared (and. for the first few hours. xighter than anv 2~6 QSO).," Despite the great distance to this burst, it will be the brightest gamma ray source on the sky for several tens of seconds, one of the brightest cosmological X-ray sources for hours, and its afterglow will be observable for weeks in the near-infrared (and, for the first few hours, brighter than any $z \sim 6$ QSO)."
+ Much of the optical and near-infrared light will be obscured by the Lye forest. and his obscuration will enable the strongest constraint to date on the neutral hydrogen fraction of the intergalactic medium (IGMD at the bursts redshift.," Much of the optical and near-infrared light will be obscured by the $\alpha$ forest, and this obscuration will enable the strongest constraint to date on the neutral hydrogen fraction of the intergalactic medium (IGM) at the burst's redshift."
+ In facet. such an occurrence may already have been realised. Llaislipetal.(2006)... Ixawaietal. (2006).. lTagliaferrictal. (2005).," In fact, such an occurrence may already have been realised. \citet{haislip06}, \citet{kawai05}, \citet{tagliaferri05},"
+. ancl Totanietal.(2006) observed and analysed the optical/near-infrared afterglow of CGI1UD050904. identified to beat z—6.3. possibly during the rejionisation epoch and the GRB with the highest confirmed redshift.," and \citet{totani06} observed and analysed the optical/near-infrared afterglow of GRB050904, identified to be at $z =6.3$ – possibly during the reionisation epoch and the GRB with the highest confirmed redshift."
+ Totanietal.(2006)., \citet{totani06} derived the constraint on the global neutral fraction $\bar{x}_{H} < 0.6$ at $z = 6.3$.
+ derived the constraint on t," In this paper, we discuss the assumptions that went into their analysis, and we investigate how realistic modelling of reionisation can affect constraints on $\bar{x}_H$ from GRB050904 and from future $z>6$ GRBs."
+"he ""dark"" because Lvo absorp", The Swift satellite has greatly increased the sample of GRBs with known redshifts in the last two years \citep{gehrels04}.
+tion from the high-redshift L," Future missions such as EXIST \citep{grindlay06} and JWST \citep{gardner06}
+ will further enhance our ability to detect high-redshift GRBs and will enable more detailed follow-up studies of their near-infrared afterglows."
+AI absorbs the optical emission (e.g..," Interestingly, approximately one-half of Swift bursts are “dark bursts” – bursts that have detected X-ray afterglows, but that have no measurable optical emission (e.g., \citealt{filliatre06}) )."
+ Malesanietal. 2005)).," While it is probable that most dark bursts originate from low-redshift, dust-rich galaxies, a fraction of dark bursts may originate from $z>6$ and are “dark” because $\alpha$ absorption from the high-redshift IGM absorbs the optical emission (e.g., \citealt{malesani05}) )."
+ 1n addition to their extreme [uminositv. there are several other advantages to studying reionisation with GRBs compared to other probes of this epoch.," In addition to their extreme luminosity, there are several other advantages to studying reionisation with GRBs compared to other probes of this epoch."
+ First. the afterglows of high-redshift CRBs are observed. at. earlier. (brighter) times in the source frame than those at. lower redshifts. so the dimming owing to increased Luminosity distance is nearly cancelled. and the observed. Dux is almost independent. of redshift (Lamb&Reichart2001:CiardiLoeb.," First, the afterglows of high-redshift GRBs are observed at earlier (brighter) times in the source frame than those at lower redshifts, so the dimming owing to increased luminosity distance is nearly cancelled, and the observed flux is almost independent of redshift \citep{lamb01, ciardi99}."
+2000).. second. unlike the spectra of galaxies ancl quasars. the intrinsic. afterglow spectrum. of a GIU is a featureless power-law at the relevant wavelengths. allowing a more precise measurement of absorption owing to a neutral IGM (Barkana&Loeb2004).," Second, unlike the spectra of galaxies and quasars, the intrinsic afterglow spectrum of a GRB is a featureless power-law at the relevant wavelengths, allowing a more precise measurement of absorption owing to a neutral IGM \citep{barkana04b}."
+ Finally. since. the theoretical expectation is that most of the star formation at 2=6 occurs in halos with im~LO?AL. and because observations ab 2m6 currently probe only the most massive. galaxies and QSOs (m107AL. ). high-redshift CRB host galaxies should be less massive than galaxies selected. in. another manner.," Finally, since the theoretical expectation is that most of the star formation at $z \gtrsim 6$ occurs in halos with $m \sim 10^9 \; \Msun$ and because observations at $z \gtrsim 6$ currently probe only the most massive galaxies and QSOs $m \gtrsim
+10^{11} ~\Msun$ ), high-redshift GRB host galaxies should be less massive than galaxies selected in another manner."
+ Consequently. CRB host galaxies will sit in smaller LU regions during reionisation (on average) than galaxies," Consequently, GRB host galaxies will sit in smaller HII regions during reionisation (on average) than galaxies"
+will be neglected. because where Z4i4. denotes the duration of disc heating by a single spike of the potential [luctuations. which is onlv cllective close to the peak in Fig.,"will be neglected, because where $T_{\rm acc}$ denotes the duration of disc heating by a single spike of the potential fluctuations, which is only effective close to the peak in Fig."
+ 1l. whereas i is the orbital period of the stars.," 1, whereas $T_{\rm orb}$ is the orbital period of the stars."
+ In. this aspect. the disc heating mechanism described here is rather impulsive., In this aspect the disc heating mechanism described here is rather impulsive.
+ Next ] consider the integration in equation (23) with respect to lh. where the upper row refers to Dj. the middle row to Di»=Da. and the lower row to Doo. respectively.," Next I consider the integration in equation (23) with respect to $l_1$, where the upper row refers to $D_{11}$, the middle row to $D_{12}=D_{21}$, and the lower row to $D_{22}$, respectively."
+" The results are given by delta functions and derivatives thereof. The next step is the integration with respect to the angle variable wy. which leads for the diffusion coellicient Dos to for Di=Po, to and for D, to Phe integration over wj gives simply a factor of 2x for Dy. whereas the integration> over the trigonometricὃν functions in equation (87) leads to the result Integrating equation (38) with respect to wj gives The final step is the integration with respect to fo. where the upper rows refer to the cdillusion coelficien Oy and the lower rows to o». respectively."," The results are given by delta functions and derivatives thereof, The next step is the integration with respect to the angle variable $w_1$ , which leads for the diffusion coefficient $D_{22}$ to for $D_{12}=D_{21}$ to and for $D_{11}$ to The integration over $w'_1$ gives simply a factor of $2 \pi$ for $D_{11}$, whereas the integration over the trigonometric functions in equation (37) leads to the result Integrating equation (38) with respect to $w'_1$ gives The final step is the integration with respect to $l_2$ , where the upper rows refer to the diffusion coefficient $D_{11}$ and the lower rows to $D_{22}$, respectively."
+" There. is a formal divergence at fo = 0 in the second term of Oy, on the ths of equation. (41).", There is a formal divergence at $l_2$ = 0 in the second term of $D_{11}$ on the lhs of equation (41).
+ E have chosen to ignore this ancl have replaced. the integral by a saddle poin approximation., I have chosen to ignore this and have replaced the integral by a saddle point approximation.
+ “Phe reason is that the mocdel of disc heating used here (cf, The reason is that the model of disc heating used here (cf.
+" equation 25) becomes unphysical lor smal circumferential wave numbers ο, because the density waves approach then the WIXB limit and become long lived. so tha they do not heat the disc ellectivelv."," equation 25) becomes unphysical for small circumferential wave numbers $l_2$, because the density waves approach then the WKB limit and become long lived, so that they do not heat the disc effectively."
+ Due to the svmmoetry of the distribution of amplitudes (24) with respect to k the ellect of density waves with negative wave numbers ἂν can be taken into account by multiplying the diffusion coefficients by a factor of two., Due to the symmetry of the distribution of amplitudes (24) with respect to ${\bf k}$ the effect of density waves with negative wave numbers $k_{\rm y}$ can be taken into account by multiplying the diffusion coefficients by a factor of two.
+ Assembling all results leacs a diffusion tensor of the form., Assembling all results leads a diffusion tensor of the form.
+" with D,=Sra[bulezσι,hoofA."," with $D_{\rm 0} = 8 \pi^\frac{7}{2} |\Phi_0|^2\sigma^2_{\rm k_{\rm x}}
+\sigma_{\rm k_{\rm y}} k_{\rm 20}/A $."
+" The diffusion equation takes the form where the overhead tilde means the the J, dependence of Dy, has been written separately.", The diffusion equation takes the form where the overhead tilde means the the $J_1$ –dependence of $D_{11}$ has been written separately.
+ For this kind. of disc heating by transient spiral density waves no correlation between the dillsion in racial action ancl angular momentuni space is found., For this kind of disc heating by transient spiral density waves no correlation between the diffsion in radial action and angular momentum space is found.
+ Phe clilfusion equation (43) is highhy nonlinear. because the cilfusion coellicients. depend. on. the distribution. function. (f£) themselves.," The diffusion equation (43) is highly non--linear, because the diffusion coefficients depend on the distribution function $\langle f \rangle$ themselves."
+" In. particular. the cllectivity. of swing amplification of spiral density. waves depends critically on the value of the “Toone stability parameter Q = go/(3.360%..). where a, denotes the racial velocity dispersion of the stars."," In particular, the effectivity of swing amplification of spiral density waves depends critically on the value of the Toomre stability parameter $Q$ = $\kappa
+\sigma_{\rm u}/(3.36 G \Sigma_{\rm d})$, where $\sigma_{\rm u}$ denotes the radial velocity dispersion of the stars."
+ Thus. when the disc heats up. the amplitudes eu of the density waves and the cilfusion coellicients. (42) will decrease. and the cise heating rate slows down to zero.," Thus, when the disc heats up, the amplitudes $\Phi_0$ of the density waves and the diffusion coefficients (42) will decrease and the disc heating rate slows down to zero."
+ Numerical simulations of the dynamical evolution of galactic disces (Sellwood Carlbere 1984. Fuchs v. Linden 1998) have shown that this can happen on Comparatively short. time scales. if the clises are left uncooled.," Numerical simulations of the dynamical evolution of galactic discs (Sellwood Carlberg 1984, Fuchs v. Linden 1998) have shown that this can happen on comparatively short time scales, if the discs are left uncooled."
+ Only if the clises are cooled dynamically by adding stars on low peculiar velocity orbits. the spiral density wave activity can be maintained atà constant level despite the rising velocity dispersions.," Only if the discs are cooled dynamically by adding stars on low peculiar velocity orbits, the spiral density wave activity can be maintained ata constant level despite the rising velocity dispersions."
+ Inthat case £uz const., Inthat case $D_{\rm 0} \approx$ const.
+ and a simple solution of the dillusion equation (43) is found. by, and a simple solution of the diffusion equation (43) is found by
+board theFermi satellite during the first vear of operation. although possible association for the LAT sources. with racdio-quiet ACGNs are. tentatively proposed. in the first LAT AGN Catalog. GXbedoetal.2010a).,"board the satellite during the first year of operation, although possible association for the LAT sources with radio-quiet AGNs are tentatively proposed in the first LAT AGN Catalog \citep{abdo10a}."
+. Απ ΝΤ continues to collect cata and its sensitivity increases. some of these associations could be confirmed.," As -LAT continues to collect data and its sensitivity increases, some of these associations could be confirmed."
+ The AGN dominating the οταν sky is the blazar population. comprising Lat spectrum radio quasars (ESI) and BL Lac objects.," The AGN dominating the $\gamma$ -ray sky is the blazar population, comprising flat spectrum radio quasars (FSRQ) and BL Lac objects."
+ These sources are characterized by the presence of a compact radio core. apparent superluminal jet speed. extreme Εαν density. variability in all bands. ancl high fraction of polarized optical and radio emission.," These sources are characterized by the presence of a compact radio core, apparent superluminal jet speed, extreme flux density variability in all bands, and high fraction of polarized optical and radio emission."
+ Their observational properties are interpreted. as the result of severe beaming effects due to the orientation of the relativistic jet at very small angles to the line of sight., Their observational properties are interpreted as the result of severe beaming effects due to the orientation of the relativistic jet at very small angles to the line of sight.
+ Thanks to the episodes of enhanced luminosity across he entire. electromagnetic spectrum. ib is possible to set constraints on the physical properties of the region along the jet responsible for the emission at the various wavelengths.," Thanks to the episodes of enhanced luminosity across the entire electromagnetic spectrum, it is possible to set constraints on the physical properties of the region along the jet responsible for the emission at the various wavelengths."
+ tadio monitoring of ΙΣΤ sources suggested: that. the üghest levels of 5-ray emission is observed. close in time o radio [lares (e.g.Làhteenmaki& Valtaoja2003). and connected. with the emission ofà new jet component— (Jorstadetal.2001).," Radio monitoring of EGRET sources suggested that the highest levels of $\gamma$ -ray emission is observed close in time to radio flares \citep[e.g.][]{lathe03}, and connected with the emission ofa new jet component \citep{jorstad01}."
+. 3oth pieces of evidence give support o the idea that the strongest 5-rav emission is strictly related to a shock in the jet that produces the svnchrotron radio [lare (c.g.Marscher&CGoar1985)., Both pieces of evidence give support to the idea that the strongest $\gamma$ -ray emission is strictly related to a shock in the jet that produces the synchrotron radio flare \citep[e.g.][]{marscher85}.
+. 1t must be noted hat the EGRET data were sparse. with large uncertainties and selection ellects. precluding a clear temporal correlation )etween 5-ray activity and the emission at lower frequencies.," It must be noted that the EGRET data were sparse, with large uncertainties and selection effects, precluding a clear temporal correlation between $\gamma$ -ray activity and the emission at lower frequencies."
+ The advent of the AGILE and -ray satellites allowed us to test the results obtained in the EGRET era. providing details on the connection between οταν and racio emission o find out the mechanisms responsible for the high-energy emission.," The advent of the AGILE and $\gamma$ -ray satellites allowed us to test the results obtained in the EGRET era, providing details on the connection between $\gamma$ -ray and radio emission to find out the mechanisms responsible for the high-energy emission."
+ Correlation studies. of the rst. three months of Fermi-LAT cata ancl quasi-simultaneous Very Long Baseline Loterferometry (VLBI) observations showed. that he eamma-ray emitting blazars have faster apparent jet speeds (Listeretal.2009a).. wider apparent opening angles (Pushkarevetal.2009).. and higher variability and Doppler actor (Savolainenetal.2010)... with respect to. blazars with weak 5-rav emission.," Correlation studies of the first three months of -LAT data and quasi-simultaneous Very Long Baseline Interferometry (VLBI) observations showed that the gamma-ray emitting blazars have faster apparent jet speeds \citep{lister09c}, wider apparent opening angles \citep{pushkarev09}, and higher variability and Doppler factor \citep{savolainen10}, with respect to blazars with weak $\gamma$ -ray emission."
+ Variability. studies. from. racio o high energy. bands give important. insight to locate the site of [lares and infer the physical conditions of the jet., Variability studies from radio to high energy bands give important insight to locate the site of flares and infer the physical conditions of the jet.
+ In particular. using the most recent z-rav as well as radio xolarimetric data. a likely close connection between high 5-rav states and the activity in the pe-scale core region has »ven investigated. (seec.g.Agucloetal.2011).," In particular, using the most recent $\gamma$ -ray as well as radio polarimetric data, a likely close connection between high $\gamma$ -ray states and the activity in the pc-scale core region has been investigated \citep[see e.g.][]{agudo11}."
+". In addition. ""ushkarev ct al. ("," In addition, Pushkarev et al. ("
+2010) found. for a sample of 150 sources a non-zero time delay between radio emission. measured »w pe-seale observations at 15 CGllz ancl y-ray radiation detected by Lerni-LAT. suggesting that the delay is most ikely connected with svnchrotron opacity in the core region.,"2010) found for a sample of 186 sources a non-zero time delay between radio emission measured by pc-scale observations at 15 GHz and $\gamma$ -ray radiation detected by -LAT, suggesting that the delay is most likely connected with synchrotron opacity in the core region."
+ Among blazars. 11510-089 is an ideal target to investigate the location of the emitting region and. the physical processes occurring in relativistic jets.," Among blazars, 1510-089 is an ideal target to investigate the location of the emitting region and the physical processes occurring in relativistic jets."
+ This object is a Lat spectrum radio quasar at z=0.361 CFhompsonetal.L990) with highly polarized optical emission., This object is a flat spectrum radio quasar at $z=0.361$ \citep{thompson90} with highly polarized optical emission.
+ The svnchrotron emission has its peak in the Ht band. while the inverse Compton component peaks in the 5-ray regime.," The synchrotron emission has its peak in the IR band, while the inverse Compton component peaks in the $\gamma$ -ray regime."
+ Η1σία flux density variability is observed throughout the electromagnetic spectrum. from the radio to. the high energy bands.," High flux density variability is observed throughout the electromagnetic spectrum, from the radio to the high energy bands."
+ La particular. episodes of high 5-rav activity were recently detected bv both ΟΠΗΣΕΛ (Cutini&Lays2009:Ciprini&Corbel2000:Tramacere2008) and the Comma Rav Imaging Detector (GRID) on board AGILE (Strianietal.2OLO:Vercellone.2009:Pueellaοἱal.2009:D'Xmmandoet2009. 2008).," In particular, episodes of high $\gamma$ -ray activity were recently detected by both -LAT \citep{cutini09,ciprini09,tramacere08}
+ and the Gamma Ray Imaging Detector (GRID) on board AGILE \citep{striani10,vercellone09,pucella09,dammando09a,dammando08}."
+. Description of these events. together with their. possible connection with multiwavelengths emission can be found in. D’Ammancoetal.(2009b)... al.(2010).. Abeloctal.(2010b).. anc D'Xmnmandoοἱ," Description of these events, together with their possible connection with multiwavelengths emission can be found in \citet{pucella08}, \citet{dammando09b}, \citet{marscher10}, \citet{abdo10b}, and \citet{dammando11}."
+ In the radio band the emission is dominated. by the core component., In the radio band the emission is dominated by the core component.
+ Multi-epoch. parsec-scale observations revealec highly superluminal knots with apparent velocity exceeding 20e (e.g.Lomanctal.2001:Jorstadet2005:Lister2009b:Marscheretal. 2010).. ejected along the north-wes direction at an angle of about -307 with respect to the core.," Multi-epoch parsec-scale observations revealed highly superluminal knots with apparent velocity exceeding $c$ \citep[e.g.][]{homan01,jorstad05,lister09b,marscher10}, ejected along the north-west direction at an angle of about $^{\circ}$ with respect to the core."
+ On aresccond scale the jet structure is oriented in he opposite direction. indicating a severe misalignment of almost 1901 between the pe- and kpc-scale jet.," On arcsecond scale the jet structure is oriented in the opposite direction, indicating a severe misalignment of almost $^{\circ}$ between the pc- and kpc-scale jet."
+ Misalignmen oween pc- and kpe-scale structure has been [requently observed in core-dominated. sources (e.g.Pearson&Reacl-reac 1988)., Misalignment between pc- and kpc-scale structure has been frequently observed in core-dominated sources \citep[e.g.][]{pearson88}.
+. However. müisalignment. larecr than 1107 as hose found in 0054|556 and 1652|398 (Listeretal.2001) is very rare.," However, misalignment larger than $^{\circ}$ as those found in 0954+556 and 1652+398 \citep{lister01}
+ is very rare."
+ The extraordinary case of PAS 1510-089 vas been described by Lomanetal.(2002a) assuming a small change of about 127-24 in the intrinsic jet direction. hat appears amplified due to projection elfects. providing a simple explanation for the observed: morphology.," The extraordinary case of PKS 1510-089 has been described by \citet{homan02b} assuming a small change of about $^{\circ}$ $^{\circ}$ in the intrinsic jet direction, that appears amplified due to projection effects, providing a simple explanation for the observed morphology."
+ Both he misalignment and the highly superluminal jet. speed indicate that the jet axis of 1510-0580 forms a very small angle of a few degrees to PINSthe line of sight., Both the misalignment and the highly superluminal jet speed indicate that the jet axis of 1510-089 forms a very small angle of a few degrees to the line of sight.
+ Such an extreme orientation enhances beaming cllects making this source a good target to investigate the possible connection between οταν [ares and the ejection of jet components., Such an extreme orientation enhances beaming effects making this source a good target to investigate the possible connection between $\gamma$ -ray flares and the ejection of jet components.
+ In this paper we present new results of proprietary multi-epoch polarimetric VLBI and Space-VLDI observations of carried out— at 48. S4 and 22 Cllz between 1999 and 11510-0809.2001 and not vet. publishect. with the aim of studying changes in the source structure and investigating their possible connection with other physical properties. such as lux density variability. spectral. index distribution and polarization properties.," In this paper we present new results of proprietary multi-epoch polarimetric VLBI and Space-VLBI observations of 1510-089 carried out at 4.8, 8.4 and 22 GHz between 1999 and 2001 and not yet published, with the aim of studying changes in the source structure and investigating their possible connection with other physical properties, such as flux density variability, spectral index distribution and polarization properties."
+ We then compare our results with multi-epoch Very long Baseline Array (VLBA) data at 15 Cllz from the MO.JAVIZ (Monitoring Of Jets in Active galactie nuclei with VLBA Experiments) (Listeretal.2009€) spanning a larger time interval. between 1995 and 2010.," We then compare our results with multi-epoch Very long Baseline Array (VLBA) data at 15 GHz from the MOJAVE (Monitoring Of Jets in Active galactic nuclei with VLBA Experiments) \citep{lister09a}
+ spanning a larger time interval, between 1995 and 2010."
+ The addition. of high energv information from AXCGILIS and data gives. us a clue to connect 5-ray. emission and radio properties. like Hux density variability and changes in the pe-seale racio structure for the period," The addition of high energy information from AGILE and data gives us a clue to connect $\gamma$ -ray emission and radio properties, like flux density variability and changes in the pc-scale radio structure for the period"
+(Bok-elobules) to highly embedded cores.,(Bok-globules) to highly embedded cores.
+ Tn low pressure regions. a hieher extinction of the fibuneuts leads to au enhanced cussion in the Bavleigh-Jewus duit of the spectrin.," In low pressure regions, a higher extinction of the filaments leads to an enhanced emission in the Rayleigh-Jeans limit of the spectrum."
+ This behavior is respousible for the pressurc-extinction relation of the filaments., This behavior is responsible for the pressure-extinction relation of the filaments.
+ As I lave shown (Sect. 3.1.2)).," As I have shown (Sect. \ref{sect_sed_heating}) ),"
+ this relation significantly decreases the heating rate inside the core and produces above ly~ linaga arecr temperature variation inside the core., this relation significantly decreases the heating rate inside the core and produces above $A_V\sim 1$ mag a larger temperature variation inside the core.
+ The effect disappears if the extinction does not leac to a strong increase iu the central pressure as iu the hieh pressure reeion below Ay-~16mag., The effect disappears if the extinction does not lead to a strong increase in the central pressure as in the high pressure region below $A_V\sim 16~{\rm mag}$.
+ The lower amount of radiation from embedded cores is then largely compensated by the shift to colder dust temperatures., The lower amount of radiation from embedded cores is then largely compensated by the shift to colder dust temperatures.
+ For direct comparisons with observations. the theoretical dust cussion spectrum at long waveleueths is fitted by a inodified black-body spectriuu having Zi. ancl Hp as free paracters.," For direct comparisons with observations, the theoretical dust emission spectrum at long wavelengths is fitted by a modified black-body spectrum having $T_{\rm dust}$, $\beta$, and $\kappa_0$ as free parameters."
+ The fit is achieved using a non-linear Ht with Fy and Sy refer to the mean theoretical flux aud the mean fitted fux of the modified black body spectra (Eq. 13)), The fit is achieved using a non-linear $\chi^2$ -fit with where $F_{\lambda_i}$ and $S_{\lambda_i}$ refer to the mean theoretical flux and the mean fitted flux of the modified black body spectrum (Eq. \ref{eq_modbb}) )
+ over wavelength bins ἐν respectively.," over wavelength bins $i$, respectively."
+ Fie. As can be seen in 7... the dust," As can be seen in Fig. \ref{fig_sedfit},"
+ shorter enudsson spectrin at wavelengths from the mau deviates considerably from) the simple modified black-ον spectrmu., the dust emission spectrum at shorter wavelengths from the maximum deviates considerably from the simple modified black-body spectrum.
+ The fit is therefore restricted to the interval The derived paranmieters for the modified dlack-body Happroximation are listed in Table 1.. where I also provide jc DIunuinositv of he modified black-body. function. aud 1e inteerated huninositv.," The fit is therefore restricted to the interval The derived parameters for the modified black-body approximation are listed in Table \ref{table_fitparameters}, where I also provide the luminosity of the modified black-body function and the integrated luminosity."
+ In the model. at low extinction values. a lareer fraction of the light is absorbed by PATI -iolecules aud small dust particles aud subsequently re-enütted at shorter wavelengths.," In the model, at low extinction values, a larger fraction of the light is absorbed by PAH molecules and small dust particles and subsequently re-emitted at shorter wavelengths."
+ This causes the mocditied dack-hody function to underestimate the total cust ClUSSION., This causes the modified black-body function to underestimate the total dust emission.
+ The PAID/small dust particle contziibution disappears gradually towards highly embedded: cores. as 1ο UV aud optical helt becomes almost entirely absorbed inside the flaments heating the small dust particles., The PAH/small dust particle contribution disappears gradually towards highly embedded cores as the UV and optical light becomes almost entirely absorbed inside the filaments heating the small dust particles.
+ The last two οςinus give the FYIAL aud the total iutegrated fux of an Caussian source approximation for the brightuess profile (sec Sect. 3.3))., The last two columns give the $FWHM$ and the total integrated flux of an Gaussian source approximation for the brightness profile (see Sect. \ref{sect_mbb_profile}) ).
+" For Jj. ry. L/M, aud the FMWIAL. I provide polvnouual fits as Tuictions of the dust temperature (Table 2))."," For $\beta$, $\kappa_0$, $L/M_{\rm core}$, and the $FWHM$, I provide polynomial fits as functions of the dust temperature (Table \ref{table_coeff}) )."
+ I emphasize tha the derived. relatiis depend ou the specific dust properties. and to an extent. on the radiation field.," I emphasize that the derived relations depend on the specific dust properties, and to an extent, on the radiation field."
+ The relations provided are appropriate for dust with the properties of diffuse ISM dust., The relations provided are appropriate for dust with the properties of diffuse ISM dust.
+" As an estimate of the accuracy of the achieved ft. list: Wilyfyopyg,oD where Vig2=WEN- and Nawτ=A3. is the ummber of free parameters."," As an estimate of the accuracy of the achieved fit, I list $\sqrt{1/\chi^2_{\rm red}}$ where $\chi^2_{\rm red}=\chi^2/N_{\rm free}$ and $N_{\rm free}=N-3$ is the number of free parameters."
+ The accuracy is within to with the largest deviation occuring around the peak of the dust re-eniüssion spectrum., The accuracy is within to with the largest deviation occuring around the peak of the dust re-emission spectrum.
+ The ft becomes worse for higher extinction values because of the lareer teiiperatfure variations in the core., The fit becomes worse for higher extinction values because of the larger temperature variations in the core.
+ Closer agreement is achieved again at extinction values larecr than Ay=Glmas., Closer agreement is achieved again at extinction values larger than $A_V=64~{\rm mag}$ .
+ Because of the larecr temperature variations inside the core for a given central extinction of the filament. the agreement is also less good in high pressure regions.," Because of the larger temperature variations inside the core for a given central extinction of the filament, the agreement is also less good in high pressure regions."
+ I note that in part the disagreement is also caused by the optical properties used in the radiative transfer caleulations. whose properties differ. as those of eraplüte erains. strouglv from a simple power-law behavior iu the FIR.," I note that in part the disagreement is also caused by the optical properties used in the radiative transfer calculations, whose properties differ, as those of graphite grains, strongly from a simple power-law behavior in the FIR."
+ Tn low pressure regions. the filaments and the condensed cores become less optically thin.," In low pressure regions, the filaments and the condensed cores become less optically thin."
+ In the liuüt of a very low pressure region. the dust cussion from the cores is identical to the one in the diffuse ISM.," In the limit of a very low pressure region, the dust emission from the cores is identical to the one in the diffuse ISM."
+" To obtain the correspouding effective values of the modified black-body. approximation for the diffuse dust euission. T considered a spherical ποοταναπο cloud with 0.01mag (assuming p/h=2«LO! feu),"," To obtain the corresponding effective values of the modified black-body approximation for the diffuse dust emission, I considered a spherical self-gravitating cloud with $A_V=0.01~{\rm mag}$ (assuming $p/k=2\times 10^4~{\rm K/cm^3}$ )."
+ The fit of the dust re-cuussion spectruüuni is uot as good as for cunission from the deuse cores because the contribution of stochastically heated all erains at the blue side of the cuuission spectrum is strouger., The fit of the dust re-emission spectrum is not as good as for emission from the dense cores because the contribution of stochastically heated small grains at the blue side of the emission spectrum is stronger.
+ The derived effective temperature is also at μι=19.86Is relatively high., The derived effective temperature is also at $T_{\rm dust}=19.86~{\rm K}$ relatively high.
+ The paramcters of the cmissivity ave 3=1.53 aud Ky=3.10cu?fe., The parameters of the emissivity are $\beta=1.83$ and $\kappa_0=3.40~{\rm cm^2/g}$.
+ If the erains are heated by the noi-atteuuated ISRF. the ratio of the luminosity of the dust cuussion to the total mass of eas and dust is close to 1 (in solar unita).," If the grains are heated by the non-attenuated ISRF, the ratio of the luminosity of the dust emission to the total mass of gas and dust is close to `1' (in solar units)."
+ The ratio obtained with the modified black-body fit is 0.736L../M..., The ratio obtained with the modified black-body fit is $0.736~L_{\odot}/M_\odot$.
+ For the total dust cussion spectra. the vatio is higher at LeoraMoo=1.119L./M..," For the total dust emission spectrum, the ratio is higher at $L_{\rm total}/M_{\rm core}=1.149~{L_{\odot}/M_{\odot}}$."
+ The propertics of the effective cuuissivity. ie uuuositv. aud the teniperaure iro Ομ in he following subsections.," The properties of the effective emissivity, the luminosity, and the temperature are described in the following subsections."
+ It ds found overall tha he dependence of the effective parameters on dus Clτσovrature is not strougly affected by the shape of he filameuts., It is found overall that the dependence of the effective parameters on dust temperature is not strongly affected by the shape of the filaments.
+ For the salle overpressure. the coutra extinction of a evlinder is dudeed lower.," For the same overpressure, the central extinction of a cylinder is indeed lower."
+ However. when averaging over the orientations the different eeoimetrv xoduces In the mean au extinction close to the one of a spherical filament.," However, when averaging over the orientations the different geometry produces in the mean an extinction close to the one of a spherical filament."
+ The dust temperature for a given overpressure is therefore closely the same., The dust temperature for a given overpressure is therefore closely the same.
+ As shown im Sect., As shown in Sect.
+ 2.1.2 at hiel overpressure. as for lughly cmibedded cores. the radiative transfer problem. does not strongly depend on the exterual pressure.," \ref{sect_pressextrel} at high overpressure, as for highly embedded cores, the radiative transfer problem does not strongly depend on the external pressure."
+ Therefore. the effect of the external pressure on the paranetors of the modified black-body spectrum weakens towards colder dust temperatures from cores.," Therefore, the effect of the external pressure on the parameters of the modified black-body spectrum weakens towards colder dust temperatures from cores."
+objects in which it remains uncertain whether (a) the reddening is related to the torus or is caused. by. dust further from the nucleus. for example in the host galaxy of the quasar: ancl (b) whether. particularly given the inhomogencous spectrophotometric dataset. for the 3CRR sample. all such objects have vet to be discovered.,"objects in which it remains uncertain whether (a) the reddening is related to the torus or is caused by dust further from the nucleus, for example in the host galaxy of the quasar; and (b) whether, particularly given the inhomogeneous spectrophotometric dataset for the 3CRR sample, all such objects have yet to be discovered."
+ Taking OScyS5 as a working definition of light reddening. we have chosen to count all such objects known to us as quasars.," Taking $0 \ltsimeq A_{V} \ltsimeq 5$ as a working definition of `light' reddening, we have chosen to count all such objects known to us as quasars."
+ We will return to this important point in Section 4.2., We will return to this important point in Section 4.2.
+ In Fig., In Fig.
+" E. (left). we plot low-frequeney radio Luminosity. Lis, against redshift z for the combined sample of ETUI sources."," \ref{fig:qf1} (left), we plot low-frequency radio luminosity $L_{151}$ against redshift $z$ for the combined sample of FRII sources."
+ The τι> plane has been binned and the »ercentage of quasars in each bin shown. along with the associated. Poisson errors.," The $L_{151}-z$ plane has been binned and the percentage of quasars in each bin shown, along with the associated Poisson errors."
+" The first thing to note is that here are very Low quasars at logi(L,21,/ W Lap i) 26.5. I", The first thing to note is that there are very few quasars at $\log_{10} (L_{151} /$ W $^{-1}$ $^{-1}) < 26.5$ .
+EExcluding this region. the quasar fraction appears Oo increase slightly as a function of both redshift and radio luminosity. ranging from 0.3 to 0.5.," Excluding this region, the quasar fraction appears to increase slightly as a function of both redshift and radio luminosity, ranging from 0.3 to 0.5."
+ However. with Poisson errors of σ~O.OS these clifferences can at best »e called marginallysignificant. and we will not consider hem further.," However, with Poisson errors of $\sigma \sim 0.08$ these differences can at best be called marginally–significant, and we will not consider them further."
+ Note also that for the hieh-luminosity bins. the median luminosity in cach bin changes with redshift. so the cHeets of redshift and luminosity have not been completely separatec.," Note also that for the high-luminosity bins, the median luminosity in each bin changes with redshift, so the effects of redshift and luminosity have not been completely separated."
+ Laing et al. (, Laing et al. (
+1994). proposed. that. by excluding lowexcitation radio galaxies (LECGs: Line Longa 1079) which have very weak or absent emission lines and lowionization narrow lines. the quasar fraction in 3CRR. was not a function of luminosity (or recshift).,"1994) proposed that by excluding low--excitation radio galaxies (LEGs; Hine Longair 1979) which have very weak or absent emission lines and low--ionization narrow lines, the quasar fraction in 3CRR was not a function of luminosity (or redshift)."
+ Pheir classification for LEGs was that they have OLLI] to Πα ratios of «0.2 and ΟΠΗ equivalent widths of 3A., Their classification for LEGs was that they have [OIII] to $\alpha$ ratios of $<0.2$ and [OIII] equivalent widths of $<3$.
+ Since the radio ealaxies in the combined samples here span a large range of redshift. Balmer and/or OLI] lines are often not observed in optical spectra and a classification scheme such as this cannot be applied.," Since the radio galaxies in the combined samples here span a large range of redshift, Balmer and/or [OIII] lines are often not observed in optical spectra and a classification scheme such as this cannot be applied."
+ Instead. we separate sources into those with high and low OU} emission line luminosities., Instead we separate sources into those with high and low [OII] emission line luminosities.
+ The division between these classes we adopt is logy(Leow/W)= 35.1., The division between these classes we adopt is $\log_{10} (L_{\rm [OII]} / {\rm W}) = 35.1$ .
+ Below this line luminosity there are very few. broad ine objects in the complete samples (Willott et al., Below this line luminosity there are very few broad line objects in the complete samples (Willott et al.
+ 1999)., 1999).
+ ‘This cut in narrow emission line luminosity was chosen to e equivalent to a rest-frame ΟΠΗ] equivalent width of 10 fora quasar with Ale=23'., This cut in narrow emission line luminosity was chosen to be equivalent to a rest-frame [OII] equivalent width of 10 for a quasar with $M_{B}=-23$.
+. Note that this division »uts more sources in the low-Iuminosity category than those vpically classified as LEGs., Note that this division puts more sources in the low-luminosity category than those typically classified as LEGs.
+ On the right-hand. side. of Fig., On the right-hand side of Fig.
+ 1 we plot only sources with OL] luminosities οσοι)235.1., \ref{fig:qf1} we plot only sources with [OII] luminosities $\log_{10} (L_{\rm [OII]} / {\rm W}) \ge 35.1$.
+ The differences between the quasar fractions in. the bins is reduced. somewhat here., The differences between the quasar fractions in the bins is reduced somewhat here.
+ Now we find that (with. the exception of the poorlypopulated: low radio luminosity. jn) the quasar fraction is 0.40 in all the bins at the le evel., Now we find that (with the exception of the poorly–populated low radio luminosity bin) the quasar fraction is 0.40 in all the bins at the $\sigma$ level.
+ The largest change has occurred in the low-redshilt. intermediate luminosity bin which contains many weak-lined BCRR. galaxies.," The largest change has occurred in the low-redshift, intermediate luminosity bin which contains many weak-lined 3CRR galaxies."
+ Note that the 1.z<2. intermediate uminosity bin has the lowest quasar fraction.," Note that the $1 \le z < 2$, intermediate luminosity bin has the lowest quasar fraction."
+ X possible reason for this is that this bin contains most of the galaxies without lines in their optical spectra which may not all ruby lie within this redshift range., A possible reason for this is that this bin contains most of the galaxies without lines in their optical spectra which may not all truly lie within this redshift range.
+ Considering all 216 high emission line luminosity sources. plotted on the right-hand side of Fig. l.," Considering all 216 high emission line luminosity sources plotted on the right-hand side of Fig. \ref{fig:qf1},"
+ we find a quasar fraction of 0.40+0.03 implving a mean torus half-opening angle of 53°+3 for the luminous population., we find a quasar fraction of $0.40 \pm 0.03$ implying a mean torus half-opening angle of $53^{\circ}\pm 3 ^{\circ}$ for the luminous population.
+ There is likely to be quite a range of torus opening angles and the small error presented here on the mean angle should not be interpreted as the dispersion in opening angles present in the population., There is likely to be quite a range of torus opening angles and the small error presented here on the mean angle should not be interpreted as the dispersion in opening angles present in the population.
+ We have repeated the analysis of this section excluding CSO sources (projected linear sizes <30 kpc)., We have repeated the analysis of this section excluding CSO sources (projected linear sizes $\leq 30$ kpc).
+ We find that this reduces the quasar fraction by =0.04 in all the bins., We find that this reduces the quasar fraction by $\approx 0.04$ in all the bins.
+ The quasar fraction of all luminous CSO sources is 0.08. different at the 27 level from that of non-CSO sources.," The quasar fraction of all luminous CSO sources is $0.56 \pm 0.08$ , different at the $\sigma$ level from that of non-CSO sources."
+ We defer cliscussion of possible reasons for a higher quasar fraction for CSOs in. low-f[requeney. selected. samples to a future paper., We defer discussion of possible reasons for a higher quasar fraction for CSOs in low-frequency selected samples to a future paper.
+ We first investigate whether the decrease in quasar fraction with decreasing narrow-line and radio Luminosity is a simple selection. effect., We first investigate whether the decrease in quasar fraction with decreasing narrow-line and radio luminosity is a simple selection effect.
+ 1£ the Dow-Iuminosity objects have lower unminosity narrow emission lines it is natural to expect that heir broad line and continuum luminosities should be lower oo (e.g. Miller et al., If the low-luminosity objects have lower luminosity narrow emission lines it is natural to expect that their broad line and continuum luminosities should be lower too (e.g. Miller et al.
+ 1992)., 1992).
+ Lt follows that their spectra may on average have lower signal-to-noise anc weak broad ines may be missed., It follows that their spectra may on average have lower signal-to-noise and weak broad lines may be missed.
+ Given the strong correlation. between narrow-line and radio luminosities (e.g. Willott et al., Given the strong correlation between narrow-line and radio luminosities (e.g. Willott et al.
+ 1999) his could. explain the deficit of broad. line objects at racio uminositics logy(List tse +)«26.5. and also the apparent constaney of the quasar fraction for objects with uminous narrow Lines discussed in Section 3.," 1999) this could explain the deficit of broad line objects at radio luminosities $\log_{10} (L_{151}$ / $^{-1}$ $^{-1}) <26.5$, and also the apparent constancy of the quasar fraction for objects with luminous narrow lines discussed in Section 3."
+ Vo test whether selection ellects such as these cou cause the apparent lack of low-luminosity broad line objectga we first caleulate the expected broad. line Huxes of objects in the BCRR anc 7€ samples as a function of Iuminositv.," To test whether selection effects such as these could cause the apparent lack of low-luminosity broad line objects, we first calculate the expected broad line fluxes of objects in the 3CRR and 7C samples as a function of luminosity."
+ To do this we make use of the fact that the narrow emission line luminosities of radio galaxies and quasars are positively correlated with Llow-frequencey. radio luminosity as shown in Willott et al. (, To do this we make use of the fact that the narrow emission line luminosities of radio galaxies and quasars are positively correlated with low-frequency radio luminosity as shown in Willott et al. (
+1999).,1999).
+ ‘This correlation has a slope of O.79+0.04 in the sense that LeoxLYS? and is most likely," This correlation has a slope of $0.79 \pm
+0.04$ in the sense that $L_{\rm [OII]} \propto L_{151}^{0.79}$ and is most likely"
+more sophisticated buoyancy mechanism (han was used by Dikpati&Charbonneau(1999).,more sophisticated buoyancy mechanism than was used by \citet{dc99}.
+. While both surface Hux-transport models and flux-transport dvnamo models include the acdvective-diffusive transport of magnetic flix. and both can consistently explain polar field patterns. there are inherent differences — surface transport models simulate (he evolution of racial fields on the photospheric Iatitude-lIongitude surface. whereas Εαν transport dvnamo models solve the axisvimietric dynamo equations For the toroidal field. D. aud (he vector potential. A (V.xA represents poloidal fields) in (he meridian plane in the convection zone.," While both surface flux-transport models and flux-transport dynamo models include the advective-diffusive transport of magnetic flux, and both can consistently explain polar field patterns, there are inherent differences – surface transport models simulate the evolution of radial fields on the photospheric latitude-longitude surface, whereas flux transport dynamo models solve the axisymmetric dynamo equations for the toroidal field, $B_T$, and the vector potential, $A$ $\nabla \times A$ represents poloidal fields) in the meridian plane in the convection zone."
+ There are additional physical effects operating in (hie evolution of magnetic fields in a dvnamo model. due to the presence of the radial How in the circulation pattern. racial diffusion. aud depth-dependent diffusivitv.," There are additional physical effects operating in the evolution of magnetic fields in a flux-transport dynamo model, due to the presence of the radial flow in the circulation pattern, radial diffusion, and depth-dependent diffusivity."
+ Similarly additional physical effects. are captured in surface transport models. namely their more realistic treatment. of longitude dependence and hence estimation of the Babcock-Leighton surface source term for poloidal and therefore polar fields.," Similarly additional physical effects are captured in surface transport models, namely their more realistic treatment of longitude dependence and hence estimation of the Babcock-Leighton surface source term for poloidal and therefore polar fields."
+ When poleward surface flow is increased in a flux-transport dvnanmo. if there is no change in the profile of meridional flow. the upward flow near the equator aud the downwad flow near the high-Jatitudes will also increase. causing (he poloidal fields produced Irom the leader polaritv to move up (o a slightly higher diffusivity region compared to where they would have been if the flow would not have increased.," When poleward surface flow is increased in a flux-transport dynamo, if there is no change in the profile of meridional flow, the upward flow near the equator and the downward flow near the high-latitudes will also increase, causing the poloidal fields produced from the leader polarity to move up to a slightly higher diffusivity region compared to where they would have been if the flow would not have increased."
+ The poloidal fields from (he Iollower spots. on the other haud. do the opposite they sink down towards the lower diffusivity region (see Figure 1c).," The poloidal fields from the follower spots, on the other hand, do the opposite — they sink down towards the lower diffusivity region (see Figure 1c)."
+ As a consequence. the equatorwearel side of the polarity division line of those poloidal fields uuclergoes Laster ciffusive decay. while (he poleward side uudergoes less decay and Cherefore remains more frozen.," As a consequence, the equatorward side of the polarity division line of those poloidal fields undergoes faster diffusive decay, while the poleward side undergoes less decay and therefore remains more frozen."
+ Since the dynamo equations solve for the vector potential. A. (the chanees in A described above get reflected in the radial component of the poloidal field. given by ———75751ὃCYrsin0).," Since the dynamo equations solve for the vector potential, $A$, the changes in $A$ described above get reflected in the radial component of the poloidal field, given by ${1 \over r^2 \sin\theta}{\partial
+\over \partial\theta}\left(A\,r\,\sin\theta\right)$."
+ Thus the presence of cepth-clepencence in the diffusivity profile ancl the radial component in (he meridional flow prolile both contribute to (he increase in the poloidal fields on the poleward side of the bipoles., Thus the presence of depth-dependence in the diffusivity profile and the radial component in the meridional flow profile both contribute to the increase in the poloidal fields on the poleward side of the bipoles.
+ This effect is not present in surface (iransport models., This effect is not present in surface transport models.
+ Nevertheless. we have shown through physical descripüon and numerical calculations that the surface flux-transport models aud flux-transport. cdviamo models. despite some differences in (heir physics. produce similar results when run wilh a very similar latitudinal profile of surface poleward meridional flow. allhough surface-transport models can better reproduce polar field patterus in latitude and time compared to any dvnamo moclel.," Nevertheless, we have shown through physical description and numerical calculations that the surface flux-transport models and flux-transport dynamo models, despite some differences in their physics, produce similar results when run with a very similar latitudinal profile of surface poleward meridional flow, although surface-transport models can better reproduce polar field patterns in latitude and time compared to any dynamo model."
+The classification of small bodies in the inner solar svstem as either asteroids or comets has historically been attempted bv different scientists using different techniques and emploving,The classification of small bodies in the inner solar system as either asteroids or comets has historically been attempted by different scientists using different techniques and employing
+atinosphere.,atmosphere.
+ The second explanation. which we prefer. is that the uou-zero value of μα is a measure of the amount of numerical dissipation in our model runs.," The second explanation, which we prefer, is that the non-zero value of $q_{\mathrm{rad}}$ is a measure of the amount of numerical dissipation in our model runs."
+ As can be clearly seen in Equation 8.. any amount of dissipation (physical or numerical). d. requires a non-zero average global heating rate. d. in order for there to be a net conversion from APE to kinetic energy to balance its loss.," As can be clearly seen in Equation \ref{eq:balance}, any amount of dissipation (physical or numerical), $\hat{\mathrm{d}}$, requires a non-zero average global heating rate, $\hat{q}$, in order for there to be a net conversion from APE to kinetic energy to balance its loss."
+ While (he uou-zero global heating rate will be supplied by the global frictional heating (gy in Equation 9)). in our numerical models the frictional heating is not included (yp=0) aud so the elobal radiative heating. «μα. must be non-zero to compensate for numerical losses (cu).," While the non-zero global heating rate will be supplied by the global frictional heating $\hat{q}_f$ in Equation \ref{eq:aveq}) ), in our numerical models the frictional heating is not included $q_f=0$ ) and so the global radiative heating, $q_{\mathrm{rad}}$, must be non-zero to compensate for numerical losses $\mathrm{d}_{\mathrm{num}}$ )."
+ This iuterpretation is supported V Lieine that αμα is higher for the ruus with drag applied. because iu (hese cases the increased amount. of dissipation (now both numerical aud physical. duidaas) iust be balanced by an increased amount of heating (grad. given that yp remalus zero).," This interpretation is supported by noting that $q_{\mathrm{rad}}$ is higher for the runs with drag applied, because in these cases the increased amount of dissipation (now both numerical and physical, $\mathrm{d}_{\mathrm{num}}+\mathrm{d}_{\mathrm{drag}}$ ) must be balanced by an increased amount of heating $q_{\mathrm{rad}}$, given that $q_f$ remains zero)."
+ In order to estimate the kinetic enereyC loss in our models. we can compare the [weelobal dissipation (assuming d= qq) to the tuput heating [rom incident stellar flux.," In order to estimate the kinetic energy loss in our models, we can compare the global dissipation (assuming $\mathrm{d}=q_{\mathrm{rad}}$ ) to the input heating from incident stellar flux."
+ Our radiative forcing is applied through a Newtonian relaxation scheme (fordetailsseeRauscher&Menou2010).. so that there is no well-defined stellar flux iucideut ou the mocel atmospliere.," Our radiative forcing is applied through a Newtonian relaxation scheme \citep[for details see][]{RM10}, so that there is no well-defined stellar flux incident on the model atmosphere."
+" However. the values chosen [or the forciug are derived from a one-dimensional radiative trausler model (Iroetal.2005) for the hot Jupiter HD 2091258b aud so it is reasonable to compare our dissipation to a stellar heating rate appropriate for that Assuming an incident stellar heating rate of 3x1077 W. our drag-free"" modelsB quac-1.5x107> W- (=dig) implies. a nunerical. dissipation. rate ofJ ~55.-4* "," However, the values chosen for the forcing are derived from a one-dimensional radiative transfer model \citep{Iro05} for the hot Jupiter HD 209458b and so it is reasonable to compare our dissipation to a stellar heating rate appropriate for that Assuming an incident stellar heating rate of $\sim3 \times 10^{22}$ W, our drag-free model's $q_{\mathrm{rad}} \simeq 1.5 \times 10^{21}$ W $=\mathrm{d}_{\mathrm{num}}$ ) implies a numerical dissipation rate of $\sim5$ ."
+For the models with an increasing amouut of physical drag applied (PMRa. b. ο) the global 1[usuuerical dissipation is dawn=dead—D aud apparently dominates over the pliysical drag loss CD). with a magnitude evaluated as αμα—D)/(3x1072)~1315'4.," For the models with an increasing amount of physical drag applied (PMRa, b, c) the global numerical dissipation is $\mathrm{d}_{\mathrm{num}}=q_{\mathrm{rad}}-\mathcal{D}$ and apparently dominates over the physical drag loss $\mathcal{D}$ ), with a magnitude evaluated as $(q_{\mathrm{rad}}-\mathcal{D})/(3 \times 10^{22}) \simeq 13-15$."
+. Without a more detailed lalysis it is difficult to know where this numerical loss of kinetic energy occurs in our inocel. be it evenly throughout the atmosphere or preferentially at regions of stroug flow convergence or shear.," Without a more detailed analysis it is difficult to know where this numerical loss of kinetic energy occurs in our model, be it evenly throughout the atmosphere or preferentially at regions of strong flow convergence or shear."
+ The amount of kinetic energy. lost may also be depeucdent ou the order aud maguitude . hyperdissipation These results imply that our nunerical scheme leads to a siguilicant tuount of kinetic energy loss for the hot Jupiter atinospheric regime. au issue that will require further attention outside of tliis paper.," The amount of kinetic energy lost may also be dependent on the order and magnitude of hyperdissipation These results imply that our numerical scheme leads to a significant amount of kinetic energy loss for the hot Jupiter atmospheric regime, an issue that will require further attention outside of this paper."
+ As expected. the amount of (physical) kinetic energy loss (D) increases with the amount of draga applied to the atiuosphere ancl it also becomes equal to a greater percentage of the radiative heating (although this is not linear. with shorter drag timescales. since the increased drag will also work to decrease wind speeds aud reduce the kinetic energy).," As expected, the amount of (physical) kinetic energy loss $\mathcal{D}$ ) increases with the amount of drag applied to the atmosphere and it also becomes equal to a greater percentage of the radiative heating (although this is not linear with shorter drag timescales, since the increased drag will also work to decrease wind speeds and reduce the kinetic energy)."
+ From the formalisin presented above. we kuow that it is not the ratio of qr/qyaq that matters in determining whether the energetics of the circulation will be altered. but the contributionthat the frietioual heating would make to the," From the formalism presented above, we know that it is not the ratio of $q_f/q_{\mathrm{rad}}$ that matters in determining whether the energetics of the circulation will be altered, but the contributionthat the frictional heating would make to the"
+our attention to the DL Lac source1255+244.. present in the +OFF Ποια.,"our attention to the BL Lac source, present in the +OFF field."
+ Luckily. this sanie source was observed by on May. 1998 in the lramework of a spectral survey ol BL Lacs by Beckmann (2002). who. somewhat surprisingly. state “lor LES 12554244 there are no PDS data.," Luckily, this same source was observed by on May 1998 in the framework of a spectral survey of BL Lacs by Beckmann (2002), who, somewhat surprisingly, state ""for 1ES 1255+244 there are no PDS data""."
+ Indeed. we retrieved the raw data from the ASI Scientific Data Center. extracted the PDS spectrum (the background has to be evaluated only on one offset [ield because the other is pointed exactly on Coma —— (he two sources are contaminating each other!)," Indeed, we retrieved the raw data from the ASI Scientific Data Center, extracted the PDS spectrum (the background has to be evaluated only on one offset field because the other is pointed exactly on Coma — the two sources are contaminating each other!)"
+ and found that the source is quite faint. consistent with zero flux being detected.," and found that the source is quite faint, consistent with zero flux being detected."
+ Because of the very short exposure (ime (~ 3 ksec) it is only possible (o give a 26 upper limit of 0.26 c(s/s in 15.100 keV. corresponding to 1.6 mCrab. however compatible with the background excess measured in ODS2.," Because of the very short exposure time $\sim$ 3 ksec) it is only possible to give a $\sigma$ upper limit of 0.26 cts/s in 15–100 keV, corresponding to 1.6 mCrab, however compatible with the background excess measured in OBS2."
+ It is worth noticing (hat it is possible to derive a more stringent upper limit on the Xταν [lux from 1ES 1255-44 at the time of our Coma ODS2 bv assuming that all the contamination in the +OFF field comes [rom this source., It is worth noticing that it is possible to derive a more stringent upper limit on the X–ray flux from 1ES 1255+44 at the time of our Coma OBS2 by assuming that all the contamination in the +OFF field comes from this source.
+ It is straightforward to show that the 0.0642E0.021 cts/s excess translates into 428+424 net counts when one takes into account ihe much shorter BL Lac observation (3340 sec) compared to our ODS2 and an upward factor (wo correction due to the ~40! off-axis position of the source. corresponding to a intensity reduction because of the triangular collimator response (Frontera 1997b).," It is straightforward to show that the $0.064\pm 0.021$ cts/s excess translates into $428\pm 424$ net counts when one takes into account the much shorter BL Lac observation (3340 sec) compared to our OBS2 and an upward factor two correction due to the $\sim 40'$ off-axis position of the source, corresponding to a intensity reduction because of the triangular collimator response (Frontera 1997b)."
+ At ils [ace value (his is consistent with the run of PDS counts as compared to the MECS counts in the BL Lac sample studied by Beckmann (2002). barring time variability effects.," At its face value this is consistent with the run of PDS counts as compared to the MECS counts in the BL Lac sample studied by Beckmann (2002), barring time variability effects."
+ Returning to the Coma observations. to remain on the sale side we decided (o exclude the +OFF field in the background evaluation. and consider only the OFF field as the background lor both the Coma observations.," Returning to the Coma observations, to remain on the safe side we decided to exclude the +OFF field in the background evaluation, and consider only the –OFF field as the ""un-contaminated"" background for both the Coma observations."
+ Moreover. just in the center of ihe +OFF field is also present the extremely weak ROSAT sourceJI25847.," Moreover, just in the center of the +OFF field is also present the extremely weak ROSAT source."
+12-242741.. llowever. in section 3 we also report the level of confidence of the excess considering the average of the measured backgrounds in the (wo positions.," However, in section 3 we also report the level of confidence of the excess considering the average of the measured backgrounds in the two positions."
+ The observed count rate of ODS1 is 0.782z:0.02 cts/s in the 15100 keV energy range. αἱ a confidence level of ~26e.," The observed count rate of OBS1 is $\pm$ 0.03 cts/s in the 15–100 keV energy range, at a confidence level of $\sim
+26\sigma$."
+ In the first analvsis the excess will respect to the thermal bremsstrahlung emission reported by the PDS was al the confidence level of 4.59., In the first analysis the excess with respect to the thermal bremsstrahlung emission reported by the PDS was at the confidence level of $\sim 4.5\sigma$.
+ The derived [lux was ~2.2xLOHergem7s! in the 2080 keV energy band (assuming a photon index Dy = 1.5)., The derived flux was $\sim 2.2\times 10^{-11}\erg$ in the 20–80 keV energy band (assuming a photon index $\Gamma_X$ = 1.5).
+ By relating the radio ancl the fluxes it was then possible to estimate a volume-averaged intracluster magnetic field By70.1546. using onlv observables (Fusco-Femiano 1999).," By relating the radio and the fluxes it was then possible to estimate a volume-averaged intracluster magnetic field $B_X\sim 0.15\mu G$, using only observables (Fusco-Femiano 1999)."
+ It should be pointed out. Chat a re-analysis of these data has evidenced a (rivial mistake in (he previous data analvsis (summing three, It should be pointed out that a re-analysis of these data has evidenced a trivial mistake in the previous data analysis (summing three
+available.,available.
+ For LBV 1806-20. two other ACIS datasets are available.," For LBV 1806-20, two other ACIS datasets are available."
+ We thus only checked that the limit derived in the same regions using the task on the HRC data was compatible with the ACIS results., We thus only checked that the limit derived in the same regions using the task on the HRC data was compatible with the ACIS results.
+ For MN53. the HRC data are the sole available. and the HRC limit on the count rate found with was transformed into actual flux using WebPIMMs and the above model.," For MN53, the HRC data are the sole available, and the HRC limit on the count rate found with was transformed into actual flux using WebPIMMs and the above model."
+ Table 3 lists our results., Table \ref{ch} lists our results.
+ As forXMM-Newton. we consider to have a secure detection when the source Is detected with at least c significance.," As for, we consider to have a secure detection when the source is detected with at least $\sigma$ significance."
+ Finally. a search was made for the cataloged (c)LBVs in the Chandra Source Catalog(CSC).," Finally, a search was made for the cataloged (c)LBVs in the Chandra Source Catalog."
+.. X-ray sources were only found close to GAL 026.47+00.02. W243 and Sher 25.," X-ray sources were only found close to GAL 026.47+00.02, W243 and Sher 25."
+ For the former. the X-ray source lies at 0.27. implying detection. but for the last two objects. the CSC X-ray sources lie at > 5°. casting doubt on their identification with the LBVs.," For the former, the X-ray source lies at 0.2”, implying detection, but for the last two objects, the CSC X-ray sources lie at $>5$ ”, casting doubt on their identification with the LBVs."
+ Sher 25 is a special case in our survey since 5 observations exist but 4 of these were still proprietary at the time of the analysis., Sher 25 is a special case in our survey since 5 observations exist but 4 of these were still proprietary at the time of the analysis.
+ These additional data greatly enhance the detection limit. though. as the total exposure time is then multiplied by 10.," These additional data greatly enhance the detection limit, though, as the total exposure time is then multiplied by 10."
+ The PI of these data. L. Townsley. kindly made them available to us.," The PI of these data, L. Townsley, kindly made them available to us."
+ The full dataset (the public exposure and the 4 private observations) was reduced and combined using theExtract software by P. Broos: we use here an extraction specifically made at the position of 225 on the full dataset., The full dataset (the public exposure and the 4 private observations) was reduced and combined using the software by P. Broos: we use here an extraction specifically made at the position of 25 on the full dataset.
+ Similarly. à deep analysis of the X-ray emission of MCW 930 was possible because the PI of the ddata. L. Bassani. kindly made the observation available to us before the end of her proprietary period.," Similarly, a deep analysis of the X-ray emission of MCW 930 was possible because the PI of the data, L. Bassani, kindly made the observation available to us before the end of her proprietary period."
+ These data were reduced as our other ddata (see Sect., These data were reduced as our other data (see Sect.
+ 3.1)., 3.1).
+ Note that there is also à, Note that there is also a
+Claim10.3. Thereholds |75| € f. Claim 10.4.,Lemma \ref{lem:2r2} ensures that in the job removal step removes at most $7$ jobs from any machine in $A$.
+ For anyJ) €T3. Mj's ," For any machine in $B$, one job is removed."
+load immediate, Hence the total number of migrations is at most $7\lfloor m/2 \rfloor + \lceil m/2 \rceil \leq 4m$ .
+ly before th, This concludes the proof of Theorem \ref{th:3}.
+e assignment at, We next turn to the algorithm .
+In Paper I we concluded that in the voung Galactic disk (here are (wo spatially different svslems. (he GB and the LGD.,"In Paper I we concluded that in the young Galactic disk there are two spatially different systems, the GB and the LGD."
+ Now we have classifiel only bv. spatial criteria a sample of O-BG stars from the catalogue. and found that the GB's elobal kinematics is essentially different [rom the kinematies of the LGD.," Now we have classified only by spatial criteria a sample of O-B6 stars from the catalogue, and found that the GB's global kinematics is essentially different from the kinematics of the LGD."
+ Not only that. multidimensional tests prove that the GB and the LGD are separated systems in the phase space.," Not only that, multidimensional two-samples tests prove that the GB and the LGD are separated systems in the phase space."
+ This does not necessarily imply that the GB is a coherent structure born from a single source such as a great molecular cloud., This does not necessarily imply that the GB is a coherent structure born from a single source such as a great molecular cloud.
+ We just note that in its present state. (he GB is a local system (whose size is well represented by our sample) showing very different. kinematic properties than the larger svstem in whieh it is embedded. (he Galactic disk.," We just note that in its present state, the GB is a local system (whose size is well represented by our sample) showing very different kinematic properties than the larger system in which it is embedded, the Galactic disk."
+ The fact that the GB is mainlv composed of certain moving groups challenges the idea of (his svstem coming from a sinele origin. and rises the question of whether we are witnessing a coherent structure wilh physical entity or we are just observing a (ransitory. picture of several smaller svstems with no common origin al all.," The fact that the GB is mainly composed of certain moving groups challenges the idea of this system coming from a single origin, and rises the question of whether we are witnessing a coherent structure with physical entity or we are just observing a transitory picture of several smaller systems with no common origin at all."
+ The answer to this question can only be sought in the dynamical study of (he moving groups that form the GB (which would require a deep knowledge of the Galactic potential ancl its asviumetries). in order (o (race back (heir trajectories in the past and discover whether they come from a single protostellar cloud or they are the (rausilory result of some clvnamical traps.," The answer to this question can only be sought in the dynamical study of the moving groups that form the GB (which would require a deep knowledge of the Galactic potential and its asymmetries), in order to trace back their trajectories in the past and discover whether they come from a single protostellar cloud or they are the transitory result of some dynamical traps."
+ We have also proved (hat the classic problem of the negative vertex deviation of voung stars in (he solar neighborhood is caused by the contribution of the GB., We have also proved that the classic problem of the negative vertex deviation of young stars in the solar neighborhood is caused by the contribution of the GB.
+ This ellect disappears once (his stellar svstem is removed from the sample. leaving the LGD -defined by its spatial cistribution- as the only remaining structure.," This effect disappears once this stellar system is removed from the sample, leaving the LGD -defined by its spatial distribution- as the only remaining structure."
+ Finally. we have observed how the presence of what is called the GB introduces disturbances in (he estimation of the Oort constants that describe the kinematics of the voung Galactic disk. making it necessary to discard its contribution by identilving and removing the GB menbers.," Finally, we have observed how the presence of what is called the GB introduces disturbances in the estimation of the Oort constants that describe the kinematics of the young Galactic disk, making it necessary to discard its contribution by identifying and removing the GB members."
+ Once the voung Disk is pruned by eliminating (he GB members. we estimate a flat rotation curve with a local velocity νουν close to the values ealeulated by several authors in the last decade for a wide range of ages.," Once the young Disk is pruned by eliminating the GB members, we estimate a flat rotation curve with a local velocity very close to the values calculated by several authors in the last decade for a wide range of ages."
+ Thus. although we find that the kinematic analvsis is not enough to decipher (he origin of the GB. it is indeed fundamental to characterize (he complexity of the voung Galactic disk and to better understaud the different moving groups that form the bulk of the GB stellar component.," Thus, although we find that the kinematic analysis is not enough to decipher the origin of the GB, it is indeed fundamental to characterize the complexity of the young Galactic disk and to better understand the different moving groups that form the bulk of the GB stellar component."
+ A comprehensive explanation of the origin of the GB will require a dynamical analvsis of (hese moving groups., A comprehensive explanation of the origin of the GB will require a dynamical analysis of these moving groups.
+ F.E. wants to thank the Departamento de Fissica Atómnmiea. Molecular v Nuclear of the," F.E. wants to thank the Departamento de sica Atómmica, Molecular y Nuclear of the"
+The galactic bulge source SLX 1735269 was discovered i] 1985 bv Skinner ct al. (,The galactic bulge source SLX 1735–269 was discovered in 1985 by Skinner et al. (
+1987) durius the2 iuission.,1987) during the mission.
+ Although observed on several occasions with other N-ray lustrunieuts (6.8.. GRANAT/SIGMA: CGoldwiurn et al.," Although observed on several occasions with other X-ray instruments (e.g., /SIGMA: Goldwurm et al."
+ 1996 and references therein:ASCA: David et al., 1996 and references therein;: David et al.
+ 1997). little is kuown about this source.," 1997), little is known about this source."
+ Coldwiurn et al., Goldwurm et al.
+ 1996 detectec the source up to about 150 keV with a spectral iudex above 30 keV of ~3., 1996 detected the source up to about 150 keV with a spectral index above 30 keV of $\sim-3$.
+ This is steeper than usually observed for black-hole candidates and herefore they teutatively suggested hat the compact object in the system is a neutrou star., This is steeper than usually observed for black-hole candidates and therefore they tentatively suggested that the compact object in the system is a neutron star.
+ observations of this source below 10 keV also could not uniquelv identify the nature of the compact object (David ct al., observations of this source below 10 keV also could not uniquely identify the nature of the compact object (David et al.
+ 1997) but they were consisqnt with the neutron star hypothesis., 1997) but they were consistent with the neutron star hypothesis.
+ The issue of the nature of the compact object in SLN 1735269 was finally settle by the discovery of a type I NX-vav burst from this source using the Wide Field Cameras onboardan (Bazzano ct al., The issue of the nature of the compact object in SLX 1735–269 was finally settled by the discovery of a type I X-ray burst from this source using the Wide Field Cameras onboard (Bazzano et al.
+ 19972: Bazzano ct al., 1997a; Bazzano et al.
+ 1997b: Cocchi et al., 1997b; Cocchi et al.
+ L998}. demonstrating that SLX 1735269 is a low-nass N-ray binary (LMXD) coutaiueg a neutron star.," 1998), demonstrating that SLX 1735–269 is a low-mass X-ray binary (LMXB) containing a neutron star."
+ So far. he rapid X-ray variability of this source has not been studied iu detail.," So far, the rapid X-ray variability of this source has not been studied in detail."
+ The neutron star nature of this svstemi motivated us to analyze the timiug behavior of this source as observed by theEvplorer (RNTE)., The neutron star nature of this system motivated us to analyze the timing behavior of this source as observed by the ).
+ We searched for. quasi-periodic oscillations (QPOs) between 300 and 1200 Iz. which are often observed in neutron star LAINBs (sce viui der Isis 19958. 1999 for reviews). aud coherent pulsatious such as observed in the accretiou-driven millisecond X-rav pulsar SAN J1s08.1: 1055(Wiajnands van der Isis 19982).," We searched for quasi-periodic oscillations (QPOs) between 300 and 1200 Hz, which are often observed in neutron star LMXBs (see van der Klis 1998, 1999 for reviews), and coherent pulsations such as observed in the accretion-driven millisecond X-ray pulsar SAX J1808.4–3658 (Wijnands van der Klis 1998a)."
+" Although those phenomena were not detected. we discovered. one characteristic of the timing behavior which. so fay. has ouly been observed. for SAN Jis,13658. increasing the similazity between SAX JIs(0s.13658 and the other neutron star LAINBs."," Although those phenomena were not detected, we discovered one characteristic of the timing behavior which, so far, has only been observed for SAX J1808.4--3658, increasing the similarity between SAX J1808.4–3658 and the other neutron star LMXBs."
+ SLX 1735269 was observed usingRXTE ou several OCCaSIOLS (seo Table 1. for a log of the observations) for a total of 11 ksec of ou-source data., SLX 1735–269 was observed using on several occasions (see Table \ref{obslog} for a log of the observations) for a total of 11 ksec of on-source data.
+ Data were collected sinultauneouslv with 16 s time resolution iu 129 photon energv channels (effective cherey range 260 keV). aud with Lys time resolution iu 256 channels (260 keV).," Data were collected simultaneously with 16 s time resolution in 129 photon energy channels (effective energy range 2–60 keV), and with 1 time resolution in 256 channels (2–60 keV)."
+" A liebt curve, au X-ray color-color diagram (CD). aud au X- harducss-iutensity diagram (IIID) were created using"," A light curve, an X-ray color-color diagram (CD), and an X-ray hardness-intensity diagram (HID) were created using"
+lit of Eq. 2..,"limit of Eq. \ref{eq:mbh},"
+ the deusitv of (active) black. hole mass inferred is pio<15s10!AZ.Mpe5 svehich is uel. sanaller than the prescut-day density of all black holes. Polo=2.9«10°AL.Mpe (ys," the density of (active) black hole mass inferred is $\rho_{\rm bh,6}<1.5\times10^{4}\,M_{\odot}\,{\rm
+Mpc}^{-3}$, which is much smaller than the present-day density of all black holes, $\rho_{\rm bh,0}=2.9\times10^{5}\,M_{\odot}\,{\rm
+Mpc}^{-3}$ \citep{yu:02}."
+ Therefore. the present lit is plausible. while allowing ercat latitude for black holes to continue to erow (and shine).," Therefore, the present limit is plausible, while allowing great latitude for black holes to continue to grow (and shine)."
+ For full Ciavdrogen) ionization. the ionization rate (Eq. 3)) i," For full (hydrogen) ionization, the ionization rate (Eq. \ref{eq:nion}) )"
+imst exceed the recombination rate. Pu=Capay(.)?. where C=G2)/In? is the ebunptuess of the ionized eas. ap=2.6.10Pans+ is the (case D) recombination coefficient for a temperature Tz:10! KK. and mys) is the particle deusity of hydrogen.," must exceed the recombination rate, $\dot{n}_{\rm
+recomb}=C\,\alpha_{\rm B}\,n_{\rm H}(z)^2$ , where $C=\langle
+n_e^2\rangle/\langle n_e\rangle^2$ is the clumpiness of the ionized gas, $\alpha_{\rm B}\approx2.6\times10^{-13}\,{\rm cm}^3\,{\rm
+s}^{-1}$ is the (case B) recombination coefficient for a temperature $T\approx10^4$ K, and $n_{\rm H}(z)$ is the particle density of hydrogen."
+" Then. where mpg=2.7&10""enu| (?) is the present-day density of barvous. aud Vy.τί0.2 Lis the primordial Ie fraction."," Then, where $n_{\rm b,0}=2.7\times10^{-7}\,{\rm cm}^{-3}$ \citep{spergel:03}
+ is the present-day density of baryons, and $Y_{\rm He}\approx0.24$ is the primordial He fraction."
+ Estimates and simulations suggest that the chunpiness is Cz1 30 (??7)..," Estimates and simulations suggest that the clumpiness is $C\approx1$ $30$ \citep{madau:99, cen:03}."
+ Therefore. even at the upper liit of Eq. 3..," Therefore, even at the upper limit of Eq. \ref{eq:nion},"
+ active black holes associated with huninous galaxies cannot be responsible for but a siuall fraction of the total ionizing flux at :z6., active black holes associated with luminous galaxies cannot be responsible for but a small fraction of the total ionizing flux at $z\approx6$.
+ Most changes to the calculations would be in the seuse of the Moniz/Mrecoml ratio., Most changes to the calculations would be in the sense of the $\dot{n}_{\rm ioniz}/\dot{n}_{\rm recomb}$ ratio.
+ Cosmic variance (77) could affect these results by a factor of two.," Cosmic variance \citep{barkana:03, rss:04b} could affect these results by a factor of two."
+ The surface deusity of the 2%6 ealaxies used here is. however. cousisteut with results from independent fields (?)..," The surface density of the $z\approx6$ galaxies used here is, however, consistent with results from independent fields \citep{bouwens:03}."
+" It is worth noting that if the SED of an accreting Ay,XLW!AZ. black hole ds than the ? teiiplate. as las been argued e.g. bv ?7.. the N-rav luit derived above would nuplv ionizing UV photons."," It is worth noting that if the SED of an accreting $M_{\rm
+bh}\lsim10^{6}\,M_{\odot}$ black hole is than the \citet{sazonov:04} template, as has been argued e.g. by \citet{madau:04}, the X-ray limit derived above would imply ionizing UV photons."
+ Such sources may still be plaving an inrportaut role with partial ionization (especially if they are not explicitly associated with huuinous galaxies). or lav be iustz!ueutal in achieving the carly reionization sugeested by WALAP (eg.?7)..," Such sources may still be playing an important role with partial ionization (especially if they are not explicitly associated with luminous galaxies), or may be instrumental in achieving the early reionization suggested by WMAP \citep[e.g.][]{rss:03b, madau:04}."
+ Our results are consistent with 7.. who use the observed soft N-rav. background to imfer the maximi contribution to the rest-frame hard ταν background at >6. and based on strong observational limits. rule out a significant contribution to the ionizing flux bv active accreting black holes.," Our results are consistent with \citet{dijkstra:04}, who use the observed soft X-ray background to infer the maximum contribution to the rest-frame hard X-ray background at $z>6$, and based on strong observational limits, rule out a significant contribution to the ionizing flux by active accreting black holes."
+ Based ou a selfreeulated prescription for the erowth of black holes in à ACDAL framework. ? motivate the local ? relation. and make predictions for massive black holes as a function of redshift.," Based on a self-regulated prescription for the growth of black holes in a $\Lambda$ CDM framework, \citet{wyithe:03b} motivate the local \citet{magorrian:98} relation, and make predictions for massive black holes as a function of redshift."
+ Their predicted uunuber counts of active black holes at high redshift are shown in Figure 2.. with the data from this work.," Their predicted number counts of active black holes at high redshift are shown in Figure \ref{fig:numdens}, with the data from this work."
+ This shows that useful constraints ou such AC'DALT models will be possible with ouly a few times larger survev areas. even at sensitivities comparable to the present CDE-S. The ? black hole mass function (see their fig.," This shows that useful constraints on such $\Lambda$ CDM models will be possible with only a few times larger survey areas, even at sensitivities comparable to the present CDF-S. The \citet{wyithe:03b} black hole mass function (see their fig."
+" 5) predicts a cumulative black hole ummber density of Hus09«10.1Mpe* down to Afy,z3:109AL... which is about five times smaller than the upper lit from our siunuple."," 5) predicts a cumulative black hole number density of $n_{\rm bh}\approx9\times10^{-4}\,{\rm Mpc}^{-3}$ down to $M_{\rm bh}\approx3\times10^{6}\,M_{\odot}$, which is about five times smaller than the upper limit from our sample."
+ This suggests that. distinct from the census of ionizing photons from such sources. larger deep survevs in fields with large high-redshift calaxy samples lav achieve «ποσο constraints on the models.," This suggests that, distinct from the census of ionizing photons from such sources, larger deep surveys in fields with large high-redshift galaxy samples may achieve stronger constraints on the models."
+" Future N-rav iuissious. such as aud will casily reach the relevant sensitivitics if auv black holes diiviug ""undimdequasus atom6 are there to be found at all."," Future X-ray missions, such as and will easily reach the relevant sensitivities – if any black holes driving “mini-quasars” at $z\approx6$ are there to be found at all."
+" We are erateful to BR. Somerville. D. Steru. M. Livio. T. Abel. and C. Vignali for discussions aud colmments, and to S. Wwithe and A. Loeb for sharing their model predictions."," We are grateful to R. Somerville, D. Stern, M. Livio, T. Abel, and C. Vignali for discussions and comments, and to S. Wyithe and A. Loeb for sharing their model predictions."
+ L.AAL acknowledges. support by NASA through contract umuber 1221666 issued by the Jet Propulsion Laboratory. California Lhustitute of Technology under NASA coutract L107.," L.A.M. acknowledges support by NASA through contract number 1224666 issued by the Jet Propulsion Laboratory, California Institute of Technology under NASA contract 1407."
+? investigated the properties of high-order g-modes in MS stellar models and show that the sharp variation in the N profile at the limit of the CC in an MS model lets clear asteroseismic signature: the oscillation of the period spacing around its mean value.,\cite{miglio2008} investigated the properties of high-order $g$ -modes in MS stellar models and show that the sharp variation in the $N$ profile at the limit of the CC in an MS model lets clear asteroseismic signature: the oscillation of the period spacing around its mean value.
+" They define oP, as the difference between the periods of a star showing such a sharp variation in N and the periods of an otherwise fictitious smooth model with the same value of L 2r.", They define $\delta P_k$ as the difference between the periods of a star showing such a sharp variation in $N$ and the periods of an otherwise fictitious smooth model with the same value of $\int_{r_0}^{r_1}{\frac{N}{r}dr}$ .
+" Assuming the ?.. JWKB (see e.g. ?)) and asymptotic approximations. they derive the following expression of 6P, for aN profile modelled by a step function where @ is a phase constant. v=(+) with N_ and N_ respectively the values of the Brunt-Viinsillaé frequency at the outer and inner borders of the j-gradient region."," Assuming the \cite{cowling1941}, JWKB (see $e.g.$ \citealt{gough2007}) ) and asymptotic approximations, they derive the following expression of $\delta P_k$ for a $N$ profile modelled by a step function where $\Phi$ is a phase constant, $\nu = \left(\frac{N_+}{N_-}\right)$ with $N_+$ and $N_-$ respectively the values of the Brunt-Väiisällä frequency at the outer and inner borders of the $\mu$ -gradient region."
+ The local buoyancy radius is defined as Then the total buoyancy radius is, The local buoyancy radius is defined as Then the total buoyancy radius is
+can thus be explained by the wide jet component which »oconies nonrelativistic.,can thus be explained by the wide jet component which becomes non–relativistic.
+ This value of ὄνῃ gives a slope of the electixn distribution p=2.37£0.13. correspouding oan Xrav spectrum with P—(p/2|1)=21st 07.," This value of $\delta_{NR}$ gives a slope of the electron distribution $p=2.37\pm0.13$, corresponding to an X–ray spectrum with $\Gamma=(p/2+1)=2.18\pm 0.07$ ."
+ The spectra are consistent wih thisvalue?., The spectra are consistent with this.
+. At these late times. the cooling Ίοςuenceyv ds ale:uly low the optical baud.," At these late times, the cooling frequency is already below the optical band."
+ Therefore. the optical axd he Xrav dieht curves should have the same decay sloTo," Therefore, the optical and the X–ray light curves should have the same decay slope."
+ After accomine for the contribution fron the superucVa. Bergeretal.(2003) and Lipkinctal.(2001) (erived Ou 2.35 between dav 10 aud day O (ic. after he jet break of he wide jet).," After accounting for the contribution from the supernova, \cite{berger} and \cite{lipkin} derived $\delta_{opt}\sim$ 2.35 between day 10 and day 40 (i.e. after the jet break of the wide jet)."
+ This is iu agreement. with he expected vaue Ope=p (Sari.Piran&Ualpern1990€ }))., This is in agreement with the expected value $\delta_{opt}=p$ \cite{sari}) ).
+ Note that he optical light curve ater LO days is dominated w the supernova light. which masX he predicted flatome.," Note that the optical light curve after 40 days is dominated by the supernova light, which masks the predicted flattening."
+ Finally. we would like to caution the reader that there lay exist oher possible interpretations of the observed fiatcling of the Xrav light-cirve.," Finally, we would like to caution the reader that there may exist other possible interpretations of the observed flattening of the X–ray light-curve."
+ Iu the cawronball acciuro (e.g. Dado. Dar De Rujula 2002) a flatome of the light curve can be associated to the afterglow contribution of a new caunonball: indeed. in Dado. Dar De Rujula (2001) the optical afterglow of GRB 030329 is explained with two cannonballs (plus the SN contribution) up to 70 days after the trigger.," In the cannonball scenario (e.g. Dado, Dar De Rujula 2002) a flattening of the light curve can be associated to the afterglow contribution of a new cannonball: indeed, in Dado, Dar De Rujula (2004) the optical afterglow of GRB 030329 is explained with two cannonballs (plus the SN contribution) up to 70 days after the trigger."
+" A flattening of the Ν΄ray light-curve could then be possible if a third cauuouball starts to contibute in the Xrav band (but not in the optical. which is dominated bv the SN Πο),"," A flattening of the X–ray light-curve could then be possible if a third cannonball starts to contibute in the X–ray band (but not in the optical, which is dominated by the SN light)."
+ IIowever. the muuber of caunouballs should be equal to the unnber— of main pulses in the prompt 5 ray enission. which are onulv two in CRB 030329.," However, the number of cannonballs should be equal to the number of main pulses in the prompt $\gamma$ –ray emission, which are only two in GRB 030329."
+ Iu the dvadosphere model (c.g. Ruffini et al., In the dyadosphere model (e.g. Ruffini et al.
+ 2003). a flatteniug ⋅↜of the light-curve↴⋅ i8 ⋅predicted. as ↑↕∐↴∖↴↸⊳∪∐∏⋯∐↸∖∐↑↴⋝↸∖↸⊳⋜⋯∐∖∐∪∐↥⋅↸∖↕⋜↧⊓↖⇁↕↴∖↴⊓↸⊳⋜↧↑∿↓∩≼↧⋜∏⇁↴∖↴⋜↧⇈↸∖↥⋅iu the more conventional fireball scenario. when a transition to the ⋀∖⊽↸∖↖↖↽↑∪∐⋜⋯↸∖⊼↻⋜⋯↴∖↴↕∪∐∪↸⊳↸⊳↿∐⋅↴∖↴∙↴⋝∏↑∪↑∐↸∖↥⋅↕⋟↸∖⋜↧⊓∐⋅↸∖↴∖↴↴∖↴⋉∖↸⊳↕∐↸⊳ to GRBD 030329noms (such as the break at ~0.5 days) would remain⋅ unexplained.," 2003), a flattening of the light-curve is predicted, as in the more conventional fireball scenario, when a transition to the Newtonian expansion occurs, but other features specific to GRB 030329 (such as the break at $\sim 0.5$ days) would remain unexplained."
+↽⋅ Thehe combinationcombinaty ofthe the exceptionalexceptioual brielituessbrightness of ; > We ousitivity ∩≱↸∣⋟−⋟⋞⋯↸↧∏⊔↕∐⊓≔↕↕↴∖↸∐↴∖↕↾↕∏↾⋅↸∪↕⋪∖⋎∏⋎∏⋎∖⊔⊔↓↙⊔⋯↴∖∪↕⊔∪↥⋅↸∖↥⋅⋜∥∐∪≼↧⋜↧↑⋜↧↸⊳∪∏↕≼↧↴∖↴↸∖⇈↕↸∖↑↕∐↴∖↴↕↴∖∷∖↴⋯∖∙↽∕∏∐∖↕⋜↧↸⊳↨↘↽∪↕⋜↧X as allowed us to study a GRB Xray afterglow up to very late the. leading toa 37 detection of the afterglow 258 days after the burst.," The combination of the exceptional brightness of and the high sensitivity of has allowed us to study a GRB X–ray afterglow up to very late time, leading to a $\lsim$ $\sigma$ detection of the afterglow 258 days after the burst."
+ Although other interpretations are possible. the data at f =37. GL and 258 days are consistent with the Xrav lieht curve being dominated bv the wide jet component cuvisaged by Berger et al. ," Although other interpretations are possible, the data at $t=$ 37, 61 and 258 days are consistent with the X–ray light curve being dominated by the wide jet component envisaged by Berger et al. ("
+Lipkin ct al. ,"2003), Lipkin et al. ("
+) and Sheth et (,2004) and Sheth et al. (
+2003).,2003).
+ In particular.(2003). aud more iportautlv.(200 we have al.evidence that . ∙∙∙ . the GRB explosion. in remarkable agrecment with what ↻↥⋅↸∖≼∐↸⊳↑↸∖≼⋜↧↴∖↴↴∖↴⋯⊔↕∐∶↴∙⊾ ↑↕∐∖↻⋜∐⋅⋜⊔⊔↸∖↑⋜∖↥∷∖↴∐↘↽∐∐∖↑↕↸⊳↸∖∐↸∖↥⋅∶↴∙⊾⋅↖↽⋜⋯≺⇂ ∙∙ densitylusit of £nthe external'1 ⋠∙ουiπα). derivedlevived byby. tlthe ..above authors by fittiug⋅∙ the radio and optical. light.∙⋅ curves.," In particular, and more importantly, we have evidence that this component became non–relativistic at $\sim$ 40 days after the GRB explosion, in remarkable agreement with what predicted assuming the parameters (kinetic energy and density of the external medium) derived by the above authors by fitting the radio and optical light curves."
+ The iuniediate‘∙⋅ predictione of this ≯∙∙interpretation∙ is. that also the radio liebt curve should show a flateuiie from 6=2 to Ove~1.5 after f ~10 davs., The immediate prediction of this interpretation is that also the radio light curve should show a flattening from $\delta=2$ to $\delta_{NR}\sim 1.5$ after $t\sim$ 40 days.
+" The publicly available radio data at these late times have large uncertainties aad Gunot exclude such a behavior,.", The publicly available radio data at these late times have large uncertainties and cannot exclude such a behavior.
+ On↽v further. analysis. : : lo 1 ↴ ⋟, Only further analysis of more radio data could settle this issue.
+ flattening iu the radio lieht curve would imply that the iG has uot yet reached the nou.relativistic phase. due o a larger total kinetic energv of the wide jet and/or a gualler deusity for the interstellar medimm. ∶≩∪↕≯↕↘⊽⋯↕↖↽↸∖∐∪↑∪∏," The lack of a flattening in the radio light curve would imply that the fireball has not yet reached the non–relativistic phase, due to a larger total kinetic energy of the wide jet and/or a smaller density for the interstellar medium. \cite{kouveliotou})"
+ flattening iu the radio lieht curve would imply that the iG has uot yet reached the nou.relativistic phase. due o a larger total kinetic energv of the wide jet and/or a gualler deusity for the interstellar medimm. ∶≩∪↕≯↕↘⊽⋯↕↖↽↸∖∐∪↑∪∏↸," The lack of a flattening in the radio light curve would imply that the fireball has not yet reached the non–relativistic phase, due to a larger total kinetic energy of the wide jet and/or a smaller density for the interstellar medium. \cite{kouveliotou})"
+ flattening iu the radio lieht curve would imply that the iG has uot yet reached the nou.relativistic phase. due o a larger total kinetic energv of the wide jet and/or a gualler deusity for the interstellar medimm. ∶≩∪↕≯↕↘⊽⋯↕↖↽↸∖∐∪↑∪∏↸∖," The lack of a flattening in the radio light curve would imply that the fireball has not yet reached the non–relativistic phase, due to a larger total kinetic energy of the wide jet and/or a smaller density for the interstellar medium. \cite{kouveliotou})"
+ flattening iu the radio lieht curve would imply that the iG has uot yet reached the nou.relativistic phase. due o a larger total kinetic energv of the wide jet and/or a gualler deusity for the interstellar medimm. ∶≩∪↕≯↕↘⊽⋯↕↖↽↸∖∐∪↑∪∏↸∖↑," The lack of a flattening in the radio light curve would imply that the fireball has not yet reached the non–relativistic phase, due to a larger total kinetic energy of the wide jet and/or a smaller density for the interstellar medium. \cite{kouveliotou})"
+ flattening iu the radio lieht curve would imply that the iG has uot yet reached the nou.relativistic phase. due o a larger total kinetic energv of the wide jet and/or a gualler deusity for the interstellar medimm. ∶≩∪↕≯↕↘⊽⋯↕↖↽↸∖∐∪↑∪∏↸∖↑⋜," The lack of a flattening in the radio light curve would imply that the fireball has not yet reached the non–relativistic phase, due to a larger total kinetic energy of the wide jet and/or a smaller density for the interstellar medium. \cite{kouveliotou})"
+ flattening iu the radio lieht curve would imply that the iG has uot yet reached the nou.relativistic phase. due o a larger total kinetic energv of the wide jet and/or a gualler deusity for the interstellar medimm. ∶≩∪↕≯↕↘⊽⋯↕↖↽↸∖∐∪↑∪∏↸∖↑⋜↧," The lack of a flattening in the radio light curve would imply that the fireball has not yet reached the non–relativistic phase, due to a larger total kinetic energy of the wide jet and/or a smaller density for the interstellar medium. \cite{kouveliotou})"
+ flattening iu the radio lieht curve would imply that the iG has uot yet reached the nou.relativistic phase. due o a larger total kinetic energv of the wide jet and/or a gualler deusity for the interstellar medimm. ∶≩∪↕≯↕↘⊽⋯↕↖↽↸∖∐∪↑∪∏↸∖↑⋜↧↕," The lack of a flattening in the radio light curve would imply that the fireball has not yet reached the non–relativistic phase, due to a larger total kinetic energy of the wide jet and/or a smaller density for the interstellar medium. \cite{kouveliotou})"
+ flattening iu the radio lieht curve would imply that the iG has uot yet reached the nou.relativistic phase. due o a larger total kinetic energv of the wide jet and/or a gualler deusity for the interstellar medimm. ∶≩∪↕≯↕↘⊽⋯↕↖↽↸∖∐∪↑∪∏↸∖↑⋜↧↕∙," The lack of a flattening in the radio light curve would imply that the fireball has not yet reached the non–relativistic phase, due to a larger total kinetic energy of the wide jet and/or a smaller density for the interstellar medium. \cite{kouveliotou})"
+ flattening iu the radio lieht curve would imply that the iG has uot yet reached the nou.relativistic phase. due o a larger total kinetic energv of the wide jet and/or a gualler deusity for the interstellar medimm. ∶≩∪↕≯↕↘⊽⋯↕↖↽↸∖∐∪↑∪∏↸∖↑⋜↧↕∙⊇," The lack of a flattening in the radio light curve would imply that the fireball has not yet reached the non–relativistic phase, due to a larger total kinetic energy of the wide jet and/or a smaller density for the interstellar medium. \cite{kouveliotou})"
+ flattening iu the radio lieht curve would imply that the iG has uot yet reached the nou.relativistic phase. due o a larger total kinetic energv of the wide jet and/or a gualler deusity for the interstellar medimm. ∶≩∪↕≯↕↘⊽⋯↕↖↽↸∖∐∪↑∪∏↸∖↑⋜↧↕∙⊇∩," The lack of a flattening in the radio light curve would imply that the fireball has not yet reached the non–relativistic phase, due to a larger total kinetic energy of the wide jet and/or a smaller density for the interstellar medium. \cite{kouveliotou})"
+ flattening iu the radio lieht curve would imply that the iG has uot yet reached the nou.relativistic phase. due o a larger total kinetic energv of the wide jet and/or a gualler deusity for the interstellar medimm. ∶≩∪↕≯↕↘⊽⋯↕↖↽↸∖∐∪↑∪∏↸∖↑⋜↧↕∙⊇∩∩," The lack of a flattening in the radio light curve would imply that the fireball has not yet reached the non–relativistic phase, due to a larger total kinetic energy of the wide jet and/or a smaller density for the interstellar medium. \cite{kouveliotou})"
+ flattening iu the radio lieht curve would imply that the iG has uot yet reached the nou.relativistic phase. due o a larger total kinetic energv of the wide jet and/or a gualler deusity for the interstellar medimm. ∶≩∪↕≯↕↘⊽⋯↕↖↽↸∖∐∪↑∪∏↸∖↑⋜↧↕∙⊇∩∩⊔," The lack of a flattening in the radio light curve would imply that the fireball has not yet reached the non–relativistic phase, due to a larger total kinetic energy of the wide jet and/or a smaller density for the interstellar medium. \cite{kouveliotou})"
+ flattening iu the radio lieht curve would imply that the iG has uot yet reached the nou.relativistic phase. due o a larger total kinetic energv of the wide jet and/or a gualler deusity for the interstellar medimm. ∶≩∪↕≯↕↘⊽⋯↕↖↽↸∖∐∪↑∪∏↸∖↑⋜↧↕∙⊇∩∩⊔⋟," The lack of a flattening in the radio light curve would imply that the fireball has not yet reached the non–relativistic phase, due to a larger total kinetic energy of the wide jet and/or a smaller density for the interstellar medium. \cite{kouveliotou})"
+ flattening iu the radio lieht curve would imply that the iG has uot yet reached the nou.relativistic phase. due o a larger total kinetic energv of the wide jet and/or a gualler deusity for the interstellar medimm. ∶≩∪↕≯↕↘⊽⋯↕↖↽↸∖∐∪↑∪∏↸∖↑⋜↧↕∙⊇∩∩⊔⋟∙," The lack of a flattening in the radio light curve would imply that the fireball has not yet reached the non–relativistic phase, due to a larger total kinetic energy of the wide jet and/or a smaller density for the interstellar medium. \cite{kouveliotou})"
+"the article reports that a pseudobulge “is a concentration of starlight near the centre of a] galaxy. but in the disk. not extending above it as do stars in a bulge"".","the article reports that a pseudobulge ""is a concentration of starlight near the centre of [a] galaxy, but in the disk, not extending above it as do stars in a bulge""."
+ While the detection of rotating bulges. pseudo or not. goes back a long time BBabcock 1939: Rubin. Ford Ixumar 1973: Pellet. 1976: Bertola Capacecioli 1977: 'eterson. LOTS: Moebold et 110979: Rubin 1950). some classical bulges. just like low-luminosity elliptical galaxies. are expected to have significant rotation (Naab. Ixhochfar Burkert 2006: Bekki 2010).," While the detection of rotating bulges, pseudo or not, goes back a long time Babcock 1939; Rubin, Ford Kumar 1973; Pellet 1976; Bertola Capaccioli 1977; Peterson 1978; Mebold et 1979; Rubin 1980), some classical bulges, just like low-luminosity elliptical galaxies, are expected to have significant rotation (Naab, Khochfar Burkert 2006; Bekki 2010)."
+ Classical bulges can also appear o rotate due to the presence of bar DBabusiaux ct 22010)., Classical bulges can also appear to rotate due to the presence of a bar Babusiaux et 2010).
+ Furthermore. classical bulgesa can have Sérrsic (LOGS) indices (a measure of how centrally concentrated their stellar ight is. see Graham Driver 2005r and references therein) ess than 2. just as low-Iuminosity elliptical galaxies clo CCaon et 11993: Young Currie 1994: Scannapieco et 22010: Graham 2010 and references therein).," Furthermore, classical bulges can have Sérrsic (1968) indices (a measure of how centrally concentrated their stellar light is, see Graham Driver 2005 and references therein) less than 2, just as low-luminosity elliptical galaxies do Caon et 1993; Young Currie 1994; Scannapieco et 2010; Graham 2010 and references therein)."
+ To further complicate matters. classical απ pseudobulges can exist within the same galaxy such as the case of NGC 2787 (Erwin et 22003) used in the Nature article.," To further complicate matters, classical and pseudobulges can exist within the same galaxy, such as the case of NGC 2787 (Erwin et 2003) used in the Nature article."
+ Ehe classical bulec in this galaxv has a magnitude roughly 1 mag fainter than the pseudobulge. vet no distinction is mace in the Nature article’s Figure l which would have revealed. at. odds: with the Nature articles premise. that the classical bulge. rather than the pseudobulge. is the outlying Figure 1 in the Nature article reveals that only 3 or 4 alleged pseudobulges appear to depart from the AM -L relation defined by the other galaxies. while the remaining 7 or S alleged. pseudobulges. i.e. the bulk of them. agree with the main relation.," The classical bulge in this galaxy has a magnitude roughly 1 mag fainter than the pseudobulge, yet no distinction is made in the Nature article's Figure 1 which would have revealed, at odds with the Nature article's premise, that the classical bulge, rather than the pseudobulge, is the outlying Figure 1 in the Nature article reveals that only 3 or 4 alleged pseudobulges appear to depart from the $M_{\mathrm{bh}\,}$ $L$ relation defined by the other galaxies, while the remaining 7 or 8 alleged pseudobulges, i.e. the bulk of them, agree with the main relation."
+ This would appear to support the notion of the coevolution of black holes ancl pseudobulges. at odds with the article's conclusion.," This would appear to support the notion of the coevolution of black holes and pseudobulges, at odds with the article's conclusion."
+ This observation might instead refleet that the actual number of pseudobulges has been stenificanthy over-cstimated due to the use of questionable selection criteria and/or reveal a problem with the assigned bulge luminosities., This observation might instead reflect that the actual number of pseudobulges has been significantly over-estimated due to the use of questionable selection criteria and/or reveal a problem with the assigned bulge luminosities.
+ One of the few ollset galaxies is NGC 1068. an SB galaxy with a large 3 kpe bar oriented at a position angle of approximately 45 degrees. (Scoville et 11988: Thronson ct 11989).," One of the few offset galaxies is NGC 1068, an SB galaxy with a large 3 kpc bar oriented at a position angle of approximately 45 degrees (Scoville et 1988; Thronson et 1989)."
+ Given the connection between pseudobulges and the presence of large-scale. bars (ancl nuclear bars). additionally modelling the bar light. along with the Sérrsic-bulge plus the exponential-disc light. would be appropriate given that the galaxy samples now contain many barred. galaxies.," Given the connection between pseudobulges and the presence of large-scale bars (and nuclear bars), additionally modelling the bar light, along with the Sérrsic-bulge plus the exponential-disc light, would be appropriate given that the galaxy samples now contain many barred galaxies."
+ This. however. has not been clone.," This, however, has not been done."
+ Figure Lin the Nature article has used a bulge-to-total ratio of 0.11 for the barred (pseudobulge) galaxy NGC 3227 while not assigning the bar light to the bulge is known to eive à ratio of 0.068 (Cacotti 2008). the bulge is actually half a magnitude fainter than assumed.," Figure 1 in the Nature article has used a bulge-to-total ratio of 0.11 for the barred (pseudobulge) galaxy NGC 3227, while not assigning the bar light to the bulge is known to give a ratio of 0.068 (Gadotti 2008), the bulge is actually half a magnitude fainter than assumed."
+ Similarly. the barred galaxy NGC 4596 has been assigned a ratio of 0.3 in the article. where as Laurikainen et ((2005) revealed that the bulge-to-total ratio is 0.13 when the bar light is excluded. this amounts to a dillerence of nearly 1 mag in Figure 1 of the Nature article. and explains much of the apparent olfset nature of these galaxies in the Mia -Lias diagram.," Similarly, the barred galaxy NGC 4596 has been assigned a ratio of 0.3 in the article, where as Laurikainen et (2005) revealed that the bulge-to-total ratio is 0.13 when the bar light is excluded, this amounts to a difference of nearly 1 mag in Figure 1 of the Nature article, and explains much of the apparent offset nature of these galaxies in the $M_{\mathrm{bh}\,}$ $L_{\rm bulge}$ diagram."
+ A proper decomposition of the light is important. to prevent over-estimating the bulge Hux from not only large-scale bar light that is assigned to the bulge. but. [rom unmocielled nuclear bars and star clusters which can increase the apparent Sérrsic index of the bulge and consequently its derived luminosity BBaleells et 22008).," A proper decomposition of the light is important to prevent over-estimating the bulge flux from not only large-scale bar light that is assigned to the bulge, but from unmodelled nuclear bars and star clusters which can increase the apparent Sérrsic index of the bulge and consequently its derived luminosity Balcells et 2003)."
+ Moreover. eiven the coexistence of massive black holes in dense nuclear star clusters at the low mass end of the Mi -0 diagram (Seth 2008. Gonzalez-Delgado 2008: Graham Spitler 2009 and references therein). one is not looking at the full picture if one ignores these compact nuclei with masses up to 10AZ..," Moreover, given the coexistence of massive black holes in dense nuclear star clusters at the low mass end of the $M_{\mathrm{bh}\,}$ $\sigma$ diagram (Seth 2008, Gonzalez-Delgado 2008; Graham Spitler 2009 and references therein), one is not looking at the full picture if one ignores these compact nuclei with masses up to $10^7 M_{\odot}$."
+ Accounting for their mass may well partially explain. the apparent olfset nature of some (pseudobulge) galaxies at the low-mass end of the Mig -0 diagram.," Accounting for their mass may well partially explain the apparent offset nature of some (pseudobulge) galaxies at the low-mass end of the $M_{\mathrm{bh}\,}$ $\sigma$ diagram."
+ ‘Traditionally. all bulges were assumed to have the same concentration: the same Sérrsic index with a value of 4 was used to describe their radial (lux clistribution.," Traditionally, all bulges were assumed to have the same concentration; the same Sérrsic index with a value of 4 was used to describe their radial flux distribution."
+ We now know this was wrong AAndrecakis 1994: de Jong 1996: Dalcells ct 22003)., We now know this was wrong Andredakis 1994; de Jong 1996; Balcells et 2003).
+ Unfortunately the literature remains full of bulec-to-total (2/7) Dux ratios which are too high., Unfortunately the literature remains full of bulge-to-total $B/T$ ) flux ratios which are too high.
+ The average ratio For (classical bulges abundant among) the SO galaxy population is ~0.25 BBalcells et 22003: Laurikeünen et 22005. 2007. 2010: Craham Worley 2008).," The average ratio for (classical bulges abundant among) the S0 galaxy population is $\sim$ 0.25 Balcells et 2003; Laurikainen et 2005, 2007, 2010; Graham Worley 2008)."
+ Lt is therefore somewhat concerning to find half of the classical bulges in Table 1 ofthe article with ο70.5., It is therefore somewhat concerning to find half of the classical bulges in Table 1 of the article with $B/T > 0.5$.
+ At the other extreme is the unaccounted for excess central Lux seen in the late-type spiral galaxy NGC 4395 (see Graham Spitler 2009)., At the other extreme is the unaccounted for excess central flux seen in the late-type spiral galaxy NGC 4395 (see Graham Spitler 2009).
+" ""hat is. at least one of the two allegedly pure-dise galaxies with a black hole mass that is not consistent with a value of zero appears to have a bulge. pseudo or otherwise. of the size expected. for. its morphological type."," That is, at least one of the two allegedly pure-disc galaxies with a black hole mass that is not consistent with a value of zero appears to have a bulge, pseudo or otherwise, of the size expected for its morphological type."
+ Only with reliable bulge luminosities will we know if pseudobulge. ancl barred. galaxies follow the same (black hole mass)-(bulge luminosity) relation as elliptical galaxies ancl non-barred. galaxies., Only with reliable bulge luminosities will we know if pseudobulge and barred galaxies follow the same (black hole mass)-(bulge luminosity) relation as elliptical galaxies and non-barred galaxies.
+ Importantly. if no clear offset. of the former relative to the distribution of the latter in this ciagram is found. it would suggest that these disc galaxies. or at least the majority of them. have bulges which coevolve with their black holes.," Importantly, if no clear offset of the former relative to the distribution of the latter in this diagram is found, it would suggest that these disc galaxies, or at least the majority of them, have bulges which coevolve with their black holes."
+ Phe already established offset nature of (some of the) galaxies with bars ancl pseudobulges in the AMigy-o0 diagram would then suggest an issue with the measurecl dynamics. ie. the stellar “velocity dispersions? of the host bulges.," The already established offset nature of (some of the) galaxies with bars and pseudobulges in the $M_{\mathrm{bh}\,}$ $\,\sigma$ diagram would then suggest an issue with the measured dynamics, i.e. the stellar ""velocity dispersions"" of the host bulges."
+ This may arise due to dillering levels of interference from the dynamics of the bars whose light signal is mixed. with that of the bulge (an issue explore by Graham et 22010)., This may arise due to differing levels of interference from the dynamics of the bars whose light signal is mixed with that of the bulge (an issue explore by Graham et 2010).
+ Furthermore. if bar dynamics cleviate significantly. from the often assumed axisvnimetric. rather than triaxial. stellar orbits used to constrain the black hole mass. then this could also contribute to potential olfsets of barred galaxies in both the Miy-60 and Ai -L diagram.," Furthermore, if bar dynamics deviate significantly from the often assumed axisymmetric, rather than triaxial, stellar orbits used to constrain the black hole mass, then this could also contribute to potential offsets of barred galaxies in both the $M_{\mathrm{bh}\,}$ $\,\sigma$ and $M_{\mathrm{bh}\,}$ $\,L$ diagram."
+ABCC 196 is located.,ABCG 496 is located.
+ As usual for Schinidt plates the PSF inav indeed be quite poor at the borders. which randomly increases the fuzziness of otherwise poiut-like objects aud thereby leads to parameter estimates closer o those of diffuse objects than to geuuine star-like ones (cf.," As usual for Schmidt plates the PSF may indeed be quite poor at the borders, which randomly increases the fuzziness of otherwise point-like objects and thereby leads to parameter estimates closer to those of diffuse objects than to genuine star-like ones (cf."
+ Fic., Fig.
+ 1 where au overdensity is clearly visible at the edees of the plate)., 1 where an overdensity is clearly visible at the edges of the plate).
+ More generally. oue cau also questiou he efficieucy of the classification procedure itself when the ΠΡΟ deusitv of stars is vorv hieh.," More generally, one can also question the efficiency of the classification procedure itself when the number density of stars is very high."
+ Iu fact. a selection uostlv based on a surface brightness criterion nmuplies seepiug z of the total uuniber of stars mi order to select inost of the ealaxies.," In fact, a selection mostly based on a surface brightness criterion implies keeping $\simeq$ of the total number of stars in order to select most of the galaxies."
+ Hence. im most cases. the contamination level of a quite complete galaxy sample unavoidably increases with the star umber deusitv.," Hence, in most cases, the contamination level of a quite complete galaxy sample unavoidably increases with the star number density."
+ The ealactic latitude of ABCC: 196 is bj= ‘which is quite close to the Galactic plane aud may explain the high absolute umber of uvisclassified single stars.," The galactic latitude of ABCG 496 is $_{\rm II}=-$ $^{\circ}$, which is quite close to the Galactic plane and may explain the high absolute number of misclassified single stars."
+ But. for the present data. blended stars identified as a sinele diffuse object is the main explanation. as first noticed during the spectroscopic run aud partly expected from the involved detection software.," But, for the present data, blended stars identified as a single diffuse object is the main explanation, as first noticed during the spectroscopic run and partly expected from the involved detection software."
+ This is confined by a check of our list of bright (5;« 17.5) galaxy candidates against the APM list for the #6621 SRC-J field., This is confirmed by a check of our list of bright $b_{\rm J}<$ 17.5) galaxy candidates against the APM list for the 621 SRC-J field.
+ Amone our 119 candidates within the same celestial zone. 389 objects are described as uchular in the APM catalogue. 52 are classified as like. l is related to a plate flaw while 10 have no close counterpart.," Among our 449 candidates within the same celestial zone, 389 objects are described as nebular in the APM catalogue, 52 are classified as star-like, 1 is related to a plate flaw while 10 have no close counterpart."
+ As evidenced by a visual check using the DSS. znone these 62 discrepant objects. here are 8 eenuine ealaxies. 2 asvinetrical objects. LO blended stars. 8 sinele bright stars Gauportant saturation. diffraction spikes). au | star-like objects.," As evidenced by a visual check using the DSS, among these 62 discrepant objects, there are 8 genuine galaxies, 2 asymetrical objects, 40 blended stars, 8 single bright stars (important saturation, diffraction spikes), and 4 star-like objects."
+ So. it appears that ouly 52 objects are misclassified of the total umber of candidates). out of which are mereed inages. plus wo asviuetrica objects.," So, it appears that only 52 objects are misclassified of the total number of candidates), out of which are merged images, plus two asymetrical objects."
+ On oue haud. this ~ contanünation level is disturbing for studies involving individual objects pickec αλλος galaxy candidates.," On one hand, this $\simeq$ contamination level is disturbing for studies involving individual objects picked among galaxy candidates."
+ Ou the other hand. its net effect for statistical studies is only a decrease of the contrast for the signal of interest providing that the misclassifie stars are randoilv distributed.," On the other hand, its net effect for statistical studies is only a decrease of the contrast for the signal of interest providing that the misclassified stars are randomly distributed."
+ So. for such applications. the preseut photometric catalogue of ABCC 196 certainly remains valuable.," So, for such applications, the present photometric catalogue of ABCG 496 certainly remains valuable."
+" Table 1 lists the catalogue of galaxy cauclidates obtained from the SRC-J 621 plate in the 2.5° 42.5"" field of ABCC 196.", Table 1 lists the catalogue of galaxy candidates obtained from the SRC-J 621 plate in the $^{\circ}\times$ $^{\circ}$ field of ABCG 496.
+ Note that the LOL misclassified stars we were able to identify during our spectroscopic follow-up have been rejected. as well as the 51 objects selected by a visual check (the 52. iunisclassified objects and 2 asvinctrical objects described above). vielding 3.879 cutrics.," Note that the 101 misclassified stars we were able to identify during our spectroscopic follow-up have been rejected, as well as the 54 objects selected by a visual check (the 52 misclassified objects and 2 asymetrical objects described above), yielding 3,879 entries."
+ The meaning of the οπας is the following (1) running (2) to (1) right ascension (equinox (5) to (7) declination (equinox (5) halfanajor axis (9) exceutricitv e defined as yt(iy.D where e aud bare the major aud minor axes (10) position angle of the major axis (from North to (11) bi (12) and (13) N aud Y positions iu arcsecoud relative to the ceutre defined as that of the diffuse ο eiiission of the cluster (see (11) distance to cluster center iu (15) MAMA catalogue refereuce number.," The meaning of the columns is the following (1) running (2) to (4) right ascension (equinox (5) to (7) declination (equinox (8) half-major axis (9) excentricity $e$ defined as $\sqrt{1-({b\over a})^2}$, where $a$ and $b$ are the major and minor axes (10) position angle of the major axis (from North to (11) $b_{\rm J}$ (12) and (13) X and Y positions in arcsecond relative to the centre defined as that of the diffuse X-ray emission of the cluster (see (14) distance to cluster center in (15) MAMA catalogue reference number."
+ The observations were performed with the Danish Lin telescope at ESO La Silla durus 2 uiehts on November 2 and 3. 1991.," The observations were performed with the Danish 1.5m telescope at ESO La Silla during 2 nights on November 2 and 3, 1994."
+ Α sketch of the observed fields with the positions of the ealaxics appeuriug i the CCD catalogue is displaved in Fig. 3.., A sketch of the observed fields with the positions of the galaxies appearing in the CCD catalogue is displayed in Fig. \ref{champsccd}.
+ The ceutral field was centered on the coordinates of the cluster center. defined both bv the position of the cD aud bv the ceutroid of the N-ray cinission (Pislar 1998): 0 τις” 13715LF(equinox 2000.0).," The central field was centered on the coordinates of the cluster center, defined both by the position of the cD and by the centroid of the X-ray emission (Pislar 1998): $^{\rm h}$ $^{\rm mn}$ $^{\rm s}$, $-13^\circ$ 15'47""(equinox 2000.0)."
+ There was almost no overla) hetween the various fields (oulv a few aresecouds)., There was almost no overlap between the various fields (only a few arcseconds).
+ Joliison V and B. filters were used., Johnson V and R filters were used.
+ Exposure, Exposure
+tthat we have conducted in order to determine its orbital and spectral parameters.,that we have conducted in order to determine its orbital and spectral parameters.
+" In Section 4,, we discuss the implications of our findings for the nature of the unseen companion of1257+5428."," In Section \ref{sec:Companion}, we discuss the implications of our findings for the nature of the unseen companion of."
+. We summarize our conclusions and outline the future prospects for our survey in Section 5.., We summarize our conclusions and outline the future prospects for our survey in Section \ref{sec:Conclusions}.
+ The first version of the SDSS WD catalog was published by Kleinmanetal.(2004).., The first version of the SDSS WD catalog was published by \citet{kleinman04:SDSS_WDs}.
+" The most recent version, presented in Eisensteinetal.(2006) (henceforth, corresponds to Data Release 4 of the SDSS E06),(Adelman-McCarthyetal. and it is currently the largest catalog of its kind, 2006),,with 9316 spectroscopically confirmed WDs."," The most recent version, presented in \citet{eisenstein06:WD_SDSS_DR4} (henceforth, E06), corresponds to Data Release 4 of the SDSS \citep{adelman06:SDSS_DR4}, and it is currently the largest catalog of its kind, with 9316 spectroscopically confirmed WDs."
+ This number should be roughly doubled in the forthcoming DR7 version., This number should be roughly doubled in the forthcoming DR7 version.
+" Many publications have analyzed the SDSS spectra of WDs and related objects like subdwarf stars and cataclysmic variables (see E06 for a list of references), but so far all this work has focused on the final co-added spectra."," Many publications have analyzed the SDSS spectra of WDs and related objects like subdwarf stars and cataclysmic variables (see E06 for a list of references), but so far all this work has focused on the final co-added spectra."
+" All the SDSS spectra are in fact taken in three or more exposures, approximately 15 minutes long (Stoughtonetal.2002),, in order to facilitate cosmic ray rejection."," All the SDSS spectra are in fact taken in three or more exposures, approximately 15 minutes long \citep{stoughton02:SDSS_EDR}, in order to facilitate cosmic ray rejection."
+" The exact number of exposures and the time lapsed between them varies from plate to plate, depending on the observing conditions at the time, but exposures are usually taken back to back during the same night."," The exact number of exposures and the time lapsed between them varies from plate to plate, depending on the observing conditions at the time, but exposures are usually taken back to back during the same night."
+" Starting with DR7 (Abazajianetal. the individual exposures are available separately, 2009),,which opens the domain of time resolved spectroscopy in the SDSS data base."," Starting with DR7 \citep{abazajian09:SDSS_DR7}, the individual exposures are available separately, which opens the domain of time resolved spectroscopy in the SDSS data base."
+" The goal of iis to take advantage of this situation in the context of the WD spectra, looking for RV shifts among different exposures of the same object that may be the hallmark of a massive companion."," The goal of is to take advantage of this situation in the context of the WD spectra, looking for RV shifts among different exposures of the same object that may be the hallmark of a massive companion."
+" To facilitate the search, we have started by considering only WDs with ‘dominant type’ DA or DB in the E06 catalog, i.e., objects that have a relatively simple spectrum with absorption lines from either H or He."," To facilitate the search, we have started by considering only WDs with `dominant type' DA or DB in the E06 catalog, i.e., objects that have a relatively simple spectrum with absorption lines from either H or He."
+" From this list, we have discarded any objects with excess flux in the red part of the spectrum, which are usually classified as detached WD--M star binaries (Silvestrietal.2006)."," From this list, we have discarded any objects with excess flux in the red part of the spectrum, which are usually classified as detached WD+M star binaries \citep{silvestri06:WD+M}."
+". This initial triage yields 7956 objects, which we have examined for RV shifts using several automated techniques supplemented by visual examination."," This initial triage yields 7956 objects, which we have examined for RV shifts using several automated techniques supplemented by visual examination."
+" A detailed description of our search strategy and the implications for completeness, the derivation of merger rates, and other related issues will be the subject of a forthcoming publication."," A detailed description of our search strategy and the implications for completeness, the derivation of merger rates, and other related issues will be the subject of a forthcoming publication."
+" Before going on to discuss specific results, it is instructive to outline the fundamental differences between aand the SPY survey (Napiwotzkietal.2001).."," Before going on to discuss specific results, it is instructive to outline the fundamental differences between and the SPY survey \citep{napiwotzki01:SPY_survey}."
+" At a very basic level, both surveys are complementary, because the SPY objects were examined for RV shifts from the VLT telescopes in Paranal (Chile) and have therefore mostly southern declinations, while SDSS data are taken from the Apache Point Observatory (APO) in New Mexico and have mostly northern declinations."," At a very basic level, both surveys are complementary, because the SPY objects were examined for RV shifts from the VLT telescopes in Paranal (Chile) and have therefore mostly southern declinations, while SDSS data are taken from the Apache Point Observatory (APO) in New Mexico and have mostly northern declinations."
+" The SPY survey achieved an excellent RV accuracy of ~2kms7! by using the high-resolution UVES spectrograph (R— 18500), but the demands of high-resolution spectroscopy imposed a stringent magnitude cutoff of B<16.5 on the SPY WDs, even for the large collecting area (8.2 m) of the VLT Kueyen telescope (Napiwotzkietal.2001).."," The SPY survey achieved an excellent RV accuracy of $\sim2\,\mathrm{km\,s^{-1}}$ by using the high-resolution UVES spectrograph $R=18500$ ), but the demands of high-resolution spectroscopy imposed a stringent magnitude cutoff of $B\leq16.5$ on the SPY WDs, even for the large collecting area (8.2 m) of the VLT Kueyen telescope \citep{napiwotzki01:SPY_survey}."
+" By comparison, the resolution that can be achieved by wwith the spectrographs on the SDSS 2.5m telescope (R=1800,Yorketal.2000) is much lower, around 170kms-!."," By comparison, the resolution that can be achieved by with the spectrographs on the SDSS 2.5m telescope \citep[$R=1800$,][]{york00:SDSS_Technical} is much lower, around $170\,\mathrm{km\,s^{-1}}$."
+ But RV resolution is not the most important factor to find potential DDWD SN Ia progenitors., But RV resolution is not the most important factor to find potential DDWD SN Ia progenitors.
+" The most interesting systems are those with relatively massive components, or tight orbits, or both, which are expected to have RV shifts higher than ~150kms7! (Napiwotzkietal.2004)."," The most interesting systems are those with relatively massive components, or tight orbits, or both, which are expected to have RV shifts higher than $\sim150\,\mathrm{km\,s^{-1}}$ \citep{napiwotzki04:SPY_04}."
+". Moreover, by foregoing the need for high-resolution spectroscopy we can examine dimmer WDs and hence many more objects."," Moreover, by foregoing the need for high-resolution spectroscopy we can examine dimmer WDs and hence many more objects."
+" Through a combination of centroid fits and cross-correlation techniques, we have been able to achieve an accuracy of ~120kms-!, detecting RV shifts between consecutive SDSS exposures in a WD with a g magnitude of 18.9 (Mullallyetal.2009)."," Through a combination of centroid fits and cross-correlation techniques, we have been able to achieve an accuracy of $\sim120\,\mathrm{km\,s^{-1}}$, detecting RV shifts between consecutive SDSS exposures in a WD with a $g$ magnitude of $18.9$ \citep{mullally09:DDWDs}."
+". Thus, ccan examine roughly an order of magnitude more objects than SPY, albeit at a much lower resolution."," Thus, can examine roughly an order of magnitude more objects than SPY, albeit at a much lower resolution."
+" We can summarize our comparison by saying that SPY is better suited for a systematic study of the properties of DDWD systems, but iis better positioned to identify interesting objects, particularly the pre-mergers and SN Ia progenitors."," We can summarize our comparison by saying that SPY is better suited for a systematic study of the properties of DDWD systems, but is better positioned to identify interesting objects, particularly the pre-mergers and SN Ia progenitors."
+" To illustrate the kind of objects that can be found bySWARMS,, we present here results for1257+5428,, the first CWDB identified by the survey."," To illustrate the kind of objects that can be found by, we present here results for, the first CWDB identified by the survey."
+" In the E06 catalog, wwas listed as a DA WD with a g magnitude of 16.8 (see Table 1))."," In the E06 catalog, was listed as a DA WD with a $g$ magnitude of $16.8$ (see Table \ref{tab-1}) )."
+ There are three individual exposures of this WD in the DR7 SDSS data base., There are three individual exposures of this WD in the DR7 SDSS data base.
+ The WD is identified as a binary by the large shift in the centroid of the H lines between exposures 0 and 1 (taken on the night of October 3 2003) and exposure 2 (taken on the following night; see Figure , The WD is identified as a binary by the large shift in the centroid of the H lines between exposures 0 and 1 (taken on the night of October 3 2003) and exposure 2 (taken on the following night; see Figure \ref{fig-1})).
+"The shift has a value of 7.9A at the Hg line, corresponding1)). to a RV difference of 487kms !."," The shift has a value of $7.9\,\mathrm{\AA}$ at the $_{\beta}$ line, corresponding to a RV difference of $487\,\mathrm{km\,s^{-1}}$ ."
+is (minus) (he integral of the Gace of the time delay.,is (minus) the integral of the trace of the time delay.
+ This expression exhibits our choice ol orientation in the caleulation of the winding number. namely it goes [rom energv 0 to enerev x: along the side Pa of (he square.," This expression exhibits our choice of orientation in the calculation of the winding number, namely it goes from energy $0$ to energy $\infty$ along the side $B_2$ of the square."
+ Comparing with we see therelore thatthe correction term 7» arises now on the oof the equality [rom the possible contribution of Py and Ps to knowthe winding number., Comparing with we see therefore thatthe correction term $\nu$ arises now on the of the equality from the possible contribution of $\Gamma_1$ and $\Gamma_3$ to the winding number.
+" Whereas for point interactions.implies icneed not to be 0 INI]... the n fact that lor potential scattering οσο)=1 that P4601)=1 and i,=0."," Whereas for point interactions, $w_3$ need not to be $0$ \cite{KR1}, the known fact that for potential scattering $S(+\infty) = 1$ implies that $\Gamma_3(A)=1$ and $w_3=0$."
+ We now determine the contribution coming from i4., We now determine the contribution coming from $w_1$.
+ For that we could use the known results of the literature about the form of 5(0)but we will give an independent argument., For that we could use the known results of the literature about the form of $S(0)$but we will give an independent argument.
+ Its only ingredients are the unitarity of TyCA). 5(0) and RCA). whieh follows from the fact that O is a isometry. andthe explicit form of ἐς1).," Its only ingredients are the unitarity of $\Gamma_1(A)$, $S(0)$ and $R(A)$, which follows from the fact that $\Omega$ is a isometry, andthe explicit form of $R(A)$ ."
+ More precisely. in the decomposition ol J£ into 2E.p A. the operator ἐπ1) takes (he form (Dnd) with," More precisely, in the decomposition of $\H$ into $\H_e \oplus \H_o$ , the operator $R(A)$ takes the form $\left(\begin{smallmatrix}
+r_e(A) & 0 \\
+0 & r_o(A)
+\end{smallmatrix}\right)$ , with"
+surroundings. and (hen create a mask lor each magnetogram.,"surroundings, and then create a mask for each magnetogram."
+ These masks include many small clusters and isolated pixels. so only the islands with area lareer 9x9 pixels are defined as (he active regions (Hagenaar et al.," These masks include many small clusters and isolated pixels, so only the islands with area larger $\times$ 9 pixels are defined as the active regions (Hagenaar et al."
+ 2003)., 2003).
+ Considering the active regions close to the edge οἱ GO degree. in order (to avoid missing them in the automatic procedure. we always search the active regions in the solar disk with angle a less than 10 degree first. as that shown in (he left panels of Fig.," Considering the active regions close to the edge of 60 degree, in order to avoid missing them in the automatic procedure, we always search the active regions in the solar disk with angle $\alpha$ less than 70 degree first, as that shown in the left panels of Fig."
+ 1., 1.
+ Thusly. the islands with area less (han 31 pixels within GO degree disk are still defined as the active regions if thev have more than 51 pixels searched within TO degree disk.," Thusly, the islands with area less than 81 pixels within 60 degree disk are still defined as the active regions if they have more than 81 pixels searched within 70 degree disk."
+ Two magnetograms within the 70deg from disk center at approximately Che solar maximum and minimum phases. respectively. are displaved in the left panels of Fig.," Two magnetograms within the $\deg$ from disk center at approximately the solar maximum and minimum phases, respectively, are displayed in the left panels of Fig."
+ 1., 1.
+ On these retrieved magnelograms the selected ARs are masked by red curves., On these retrieved magnetograms the selected ARs are masked by red curves.
+ The criterion of selecting ARs appears to work well rom a visually examination for (he given cases., The criterion of selecting ARs appears to work well from a visually examination for the given cases.
+ In the right panels two selected sub-windows of the magnetograms are shown with contours outlining the network elements which are selected by a procedure of automatic feature selection., In the right panels two selected sub-windows of the magnetograms are shown with contours outlining the network elements which are selected by a procedure of automatic feature selection.
+ The vellow and ereen contours outline the selected netsvork elements that are belong to the components of correlated and anti-correlated with sunspots in (he solar evcle. respectively (see Section 3.2).," The yellow and green contours outline the selected network elements that are belong to the components of correlated and anti-correlated with sunspots in the solar cycle, respectively (see Section 3.2)."
+ In order to compare the evelic variations of magnetic [lux of active regions wilh that of «quiet regions. we calculate (heir magnetic flux. respectively. which is shown in the left panel of Fig.," In order to compare the cyclic variations of magnetic flux of active regions with that of quiet regions, we calculate their magnetic flux, respectively, which is shown in the left panel of Fig."
+ 2., 2.
+ At the same time. the area ratio of quiet regions is also computed. which is shown bv purple svimbols in the right panel of Fig.," At the same time, the area ratio of quiet regions is also computed, which is shown by purple symbols in the right panel of Fig."
+ 2., 2.
+ It is found that the quiel Sun contributed (0.94—1.44)x107? Mx flux from approximately the solar minimum to maxinimum in Cvele 23., It is found that the quiet Sun contributed $(0.94-1.44) \times 10^{23}$ Mx flux from approximately the solar minimum to maximum in Cycle 23.
+ The fractional area of quiet regions always exceeds in the entire solar evcle 23. and decreased. [rom the evele minimum to maximum by a factor of 1.2. although their total flux increased by a factor of 1.53: as a comparison. the active region [αν increased bv several orders of magnitude.," The fractional area of quiet regions always exceeds in the entire solar cycle 23, and decreased from the cycle minimum to maximum by a factor of 1.2, although their total flux increased by a factor of 1.53; as a comparison, the active region flux increased by several orders of magnitude."
+ The measurements confirm the global behavior of the «quiet Sun fields (see Meunier 2003 and Hagenaar et al., The measurements confirm the global behavior of the quiet Sun fields (see Meunier 2003 and Hagenaar et al.
+ 2003)., 2003).
+ Dining the 12.25 vears of Cvele 23. from October 1996 to December 2008 (see http://www.ips.gov.au). the (quiet Sun dominated the Sun'ss magnetic [lux for 7.92 vears.," During the 12.25 years of Cycle 23, from October 1996 to December 2008 (see http://www.ips.gov.au), the quiet Sun dominated the s magnetic flux for 7.92 years."
+ The monthly average magnetic flux of quiet Sun is 1.12 limes (hat of active regions., The monthly average magnetic flux of quiet Sun is 1.12 times that of active regions.
+ The magnetic fields on the quiet Sun. indeed. are a fundamental component of (he Sun'ss activity. evcle which maintains the Sun'ss magnetic οποιον and Povnting flux at a certain level.," The magnetic fields on the quiet Sun, indeed, are a fundamental component of the s activity cycle which maintains the s magnetic energy and Poynting flux at a certain level."
+that many aspects of GRBs are constant. despite the chaotic appearance of (heir light curves.,"that many aspects of GRBs are constant, despite the chaotic appearance of their light curves."
+ One implication of this constancy is thal GRBs might prove useful as tools for cosmology., One implication of this constancy is that GRBs might prove useful as tools for cosmology.
+ The lag/Iuminosity. ancl variability/huminosity relations will be the primary methods of eetting Iuminosity distances., The lag/luminosity and variability/luminosity relations will be the primary methods of getting luminosity distances.
+ Unfortunately. these relations are not as üghlit as we mieht hope.," Unfortunately, these relations are not as tight as we might hope."
+ It is possible that ρω might be used as a third parameter to improve the distance indicators. much as the Cepheid luminosities are improved with a period/color/huninosity relation.," It is possible that $E_{peak}$ might be used as a third parameter to improve the distance indicators, much as the Cepheid luminosities are improved with a period/color/luminosity relation."
+" Thus. lag/ E,,,,/Iuninosity. and variabilitv/E; /buninositv. relations might. substantially improve (he accuracy of GRBs as standard candles."," Thus, $E_{peak}$ /luminosity and $E_{peak}$ /luminosity relations might substantially improve the accuracy of GRBs as standard candles."
+ I thank Robert Mallozzi and the DATSE Team for providing the Γρ values for the 84 bursts., I thank Robert Mallozzi and the BATSE Team for providing the $E_{peak}$ values for the 84 bursts.
+ ⋡⋅↼ ∙ ⋅ ⋅⋅↽≽,density of $\sim$ 27 $^{-2}$ \citep{williams2005var}.
+ ⊔∐↲↕⋅≼↲↕⊳∖⇁∪∐↥∡∖↽≀↧↴∾⊥↸∕⊓∣↙ probability of such. à variable randomly falling in our error circle.," As our error circle covers only $\times$ $^{-4}$ $^2$, there is only a $\sim$ probability of such a variable randomly falling in our error circle."
+ In addition. the lishteurve of 13-127 (see Figure 2)) did not exhibit a simple monotonic decay over one to several mouths.," In addition, the lightcurve of r3-127 (see Figure \ref{lc}) ) did not exhibit a simple monotonic decay over one to several months."
+ At the very least. 13-127 exhibited a double-peak reminiscent ol NTE 1550-564 (Jainetal.2001b).," At the very least, r3-127 exhibited a double-peak reminiscent of XTE 1550-564 \citep{jain2001}."
+. The low sample rate of our lighteurve and the dillering decay (mes of the peaks highlight the possibility that 13-127 was exhibiting fast flaring. as has been seen in some Galactic (ransient events such as. lor example. GRO 1655-40 (see Fig.," The low sample rate of our lightcurve and the differing decay times of the peaks highlight the possibility that r3-127 was exhibiting fast flaring, as has been seen in some Galactic transient events such as, for example, GRO 1655-40 (see Fig."
+ 1 of Remillardetal. 1999)) and GRS 19154105 (see Fig., 1 of \citealp{remillard1999}) ) and GRS 1915+105 (see Fig.
+ 1 of Bauetal.2003))., 1 of \citealp{rau2003}) ).
+ If the X-ray lighteurve was complex. it is nol so surprising that the candidate counterpart did not vary as expected.," If the X-ray lightcurve was complex, it is not so surprising that the candidate counterpart did not vary as expected."
+" For example. in NRN ATE J1550-564 the optical ""rellare"" was brighter than the optical fIux during the initial X-ray outburst by more (han half a magnitude (Jainetal.2001b)."," For example, in XRN XTE J1550-564 the optical “reflare” was brighter than the optical flux during the initial X-ray outburst by more than half a magnitude \citep{jain2001}."
+. Another flaring event in this source exhibited a secondary peak in the optical (hat was not observed in A-ravs at all (Jainetal.2001a)., Another flaring event in this source exhibited a secondary peak in the optical that was not observed in X-rays at all \citep{jain2001l}.
+". source r3-12T could be a good candidate for such a reflare as the observed 0.37 keV luminosity of the NRN 10"" erg 4) is only ~LO0% of the Eddington Iuminositv of a neutron star and only ~1% of the Eddington huninosity of a typical stellar-mass black hole.", Source r3-127 could be a good candidate for such a reflare as the observed 0.3–7 keV luminosity of the XRN $\times$ $^{37}$ erg $^{-1}$ ) is only $\sim$ of the Eddington luminosity of a neutron star and only $\sim$ of the Eddington luminosity of a typical stellar-mass black hole.
+ These numbers suggest Chat only a small fraction of the mass of the disk was consumed by the oulburst. allowing the potential for another accretion event Chat could have resulted in a high optical flix in the second epoch of ACS data.," These numbers suggest that only a small fraction of the mass of the disk was consumed by the outburst, allowing the potential for another accretion event that could have resulted in a high optical flux in the second epoch of ACS data."
+ All of these arguments point to our candidate. the most clearly variable star in (he 13-127 error ellipse. as the most likely optical counterpart to r3-127.," All of these arguments point to our candidate, the most clearly variable star in the r3-127 error ellipse, as the most likely optical counterpart to r3-127."
+ With a strong candidate counterpart. we can predict the orbital period of 13-127.," With a strong candidate counterpart, we can predict the orbital period of r3-127."
+ There is an empirical relation between (he X-ray Iuminosity. optical Iuminosity. and orbital period of Galactic LAINBs curing outburst (vanParaclijs&MeClintock1994).," There is an empirical relation between the X-ray luminosity, optical luminosity, and orbital period of Galactic LMXBs during outburst \citep{vanparadijs1994}."
+. The relation suggests that outbursts that are fainter in (he optical have smaller accretion disks due to closer binary separation., The relation suggests that outbursts that are fainter in the optical have smaller accretion disks due to closer binary separation.
+ Therefore (he fainter the optical counterpart. the shorter the orbital period of the binary.," Therefore the fainter the optical counterpart, the shorter the orbital period of the binary."
+ More recently studied LAINBs have also been found (o fit this relation., More recently studied LMXBs have also been found to fit this relation.
+ applied the relation to the brightest observed. X-ray and. optical Iuminositües in a single outburst [or several recentlv-discovered transient LMXDs., \citet{williams2005bh1} applied the relation to the brightest observed X-ray and optical luminosities in a single outburst for several recently-discovered transient LMXBs.
+ For example. the brightest optical and X-ray Iuminosities observed for 4U 1543-47 during the 1983 outburst were applied to the relation regardless of the specific timing of the observations.," For example, the brightest optical and X-ray luminosities observed for 4U 1543-47 during the 1983 outburst were applied to the relation regardless of the specific timing of the observations."
+ In most cases. (le sources followed (he relation.," In most cases, the sources followed the relation."
+ Furthermore. (Williamsetal.2005€).‘) performed. detailed checks of the application of the relation when applied to optical and X-ray observations separated by," Furthermore, \citep{williams2005bh4} performed detailed checks of the application of the relation when applied to optical and X-ray observations separated by"
+" ACDAL sOXAL. ocLm ACDM This firmly established theoretical result has led to a uunuber of observational challenges. such as the “miussing satellites problemi (777?7).. the ""void. phenomenon(27)... as well as the discrepancy between the sizes of nunui-volds observed iu the local universe and those 1 Linm CDALoiumlatiAL simulationsE"," $\Lambda$ $n \propto M^\alpha$ $\alpha \approx -1.9$ $\Lambda$ This firmly established theoretical result has led to a number of observational challenges, such as the “missing satellites problem” , the “void phenomenon”, as well as the discrepancy between the sizes of mini-voids observed in the local universe and those produced in CDM simulations."
+" οAdditiοι,1 again procneedclosely related to the distribution of halo masses predicted.""(9eο by7 CDM.1 arenvay raised)Saleο bv7 theο flatnessnITUR of1 tlicB galactic luLULOSITN fiction 77). 11 LASS fiction ancl galactic stellar 1nass function the uM: low;.n»"," Additional concerns, again closely related to the distribution of halo masses predicted by CDM, are raised by the flatness of the galactic luminosity function , HI mass function and galactic stellar mass function at their faint/low-mass end."
+" A aiobs Oue ποat""Cs piu... ἂνcud.3 itha MLsationsmiea 1ornis ani he vcline 2n d-- universe Ruuverse poscan narvo ΟΠΩΣ donerweeon More mass and luminosity abarvonicnearmass."," These observational distributions display power-laws with $\alpha \approx -1.3$, much shallower than expected from the combination of a CDM universe plus a naive linear relationship between halo mass and luminosity/baryonic mass."
+" Despite their ? diversity. all: statements described abovedeseribed are ] different""n aspects. of the. same fundamental issue:. asdefied CDAI structure formation predicts laree παος of low nass haos. sccluinely in contradiction with the relative xuicitv of visible low-mass galaxies."," Despite their apparent diversity, all statements described above are just different aspects of the same fundamental issue: CDM structure formation predicts large numbers of low mass halos, seemingly in contradiction with the relative paucity of visible low-mass galaxies."
+ IHereafter. we refer o this ⋅⋅«iscrepancey as theproblem).," Hereafter, we refer to this discrepancy as the."
+.. The main (knowncaveat reeards the proper interpretation of. these observational results., The main caveat regards the proper interpretation of these observational results.
+ All phenomena mentioned so far rely on the measurement of quantities indirectly related to the amass of the hosting DAL halo (c.g. Iuninositv or HI/stellar iiass) aud. as a result. do uot provide a direct scans of comparing the ME expected for CDAL with the MIF realized in nature.," All phenomena mentioned so far rely on the measurement of quantities indirectly related to the mass of the hosting DM halo (e.g. luminosity or HI/stellar mass) and, as a result, do not provide a direct means of comparing the MF expected for CDM with the MF realized in nature."
+ Iu fact. a ummber of enviromental aud feedback effects (see refsubsecisolutioux)) are expoectec toaffect the barvonic," In fact, a number of environmental and feedback effects (see \\ref{subsec:solutions}) ) are expected toaffect the baryonic"
+These galaxies should preferentially exhibit hieh optical depth in other atomic and/or molecular transitions. including ID and CO.,"These galaxies should preferentially exhibit high optical depth in other atomic and/or molecular transitions, including H and CO."
+ There are only a lew candidates known to date., There are only a few candidates known to date.
+ The three best candidate lenses which meet (hese criteria are. D02182-357. PINS1830-211. and PIXS0201--113.," The three best candidate lenses which meet these criteria are B0218+357, PKS1830-211, and PKS0201+113."
+" Table 1. lists the source redshift. and red-shifted rest. [reeuencies of our search lines (C Lil. Lill. and CO),"," Table \ref{tab:mols} lists the source redshift, and red-shifted rest frequencies of our search lines $^7$ LiH, $^6$ LiH, and $^{13}$ CO)."
+ lis a gravitational lens with an Einstein ring., is a gravitational lens with an Einstein ring.
+ The source is à BL Lac object at a recshilt 2Q.94. and appears. al arcsecond resolution. as two distünet point sources in the ring. with a separation of 335 mas. aud a weak jet called the “hot (Patnaiketal.1993:ODeaοἱWiklind&Combes1995:Bieesοἱal.2003:Cohenet 2003).," The source is a BL Lac object at a redshift $\sim$ 0.94, and appears, at arcsecond resolution, as two distinct point sources in the ring, with a separation of 335 mas, and a weak jet called the “hot \citep{patnaik93,odea92,wiklind95,biggs03,cohen03}."
+. The point sources. A& D. have a flat spectrum. while the Einstein ring has a steep spectrum.," The point sources, A B, have a flat spectrum, while the Einstein ring has a steep spectrum."
+ They are highly variable on a timescale of a few days and each is separated into at least two subcomponents separated by a few mas (Diggsοἱal.2001:Patnaiket1995).," They are highly variable on a timescale of a few days and each is separated into at least two subcomponents separated by a few mas \citep{biggs01,patnaik95}."
+. The A source is brighter bv a factor of ~3 in (the radio ancl both have jets 2001).," The A source is brighter by a factor of $\sim$ 3 in the radio and both have jets \citep{wiklind95,biggs01}."
+. The lensing source is a spiral galaxy at a redshift of 20.685 (Patnaik 1995)..," The lensing source is a spiral galaxy at a redshift of $\sim$ 0.685 \citep{patnaik95, wiklind95}. ."
+" Atomic and molecular lines includingthe 21 em II line 1993).. II3CO (Menten&Reid 1996).. Ποσο. CS (Combes&Wiklind1997).. CO. TON (Wiklind&Combes1995).. NIL, (Henkeletal.2005).. OIL (IXanekaretal.2003).. II5CÓO anti inversion (Zeiger&Darling2010).. and Lill (tentative) (Combes&Wiklind1993) have been detected toward the A component."," Atomic and molecular lines includingthe 21 cm H line \citep{carilli93}, , $_2$ CO \citep{menten96}, , $_2$ O, CS \citep{combes97}, CO, $^+$, HCN \citep{wiklind95}, $_3$ \citep{henkel05}, OH \citep{kanekar03}, $_2$ CO anti inversion \citep{zeiger10}, and LiH (tentative) \citep{combes98} have been detected toward the A component."
+ The line center of the Lill tentative detection is off by 5 km/s compared to other detected lines. but within the CO velocity profile 1993).," The line center of the LiH tentative detection is off by $\sim$ 5 km/s compared to other detected lines, but within the CO velocity profile \citep{combes98}."
+. wwas one of the first quasars that was found to be lensed., was one of the first quasars that was found to be lensed.
+ The source blazar is al a redshilt of 272.5 and appears as two bright images in a faint Einstein ring al sub-arcsecond resolution (Lidinanetal.1999:Subrahmanvan1990).," The source blazar is at a redshift of $\sim$ 2.5 and appears as two bright images in a faint Einstein ring at sub-arcsecond resolution \citep{lidman99,subra90}."
+. The two images are NW and SE of the lens center and are separated by —1: the NW component is the brighter of the two 1999)., The two images are NW and SE of the lens center and are separated by $\sim$; the NW component is the brighter of the two \citep{jin99}.
+. The images vary in brightness. separation. and sizedue to a helical jet emanating from the core (Jinetal.2003:Nair 2005)..," The images vary in brightness, separation, and sizedue to a helical jet emanating from the core \citep{jin03,nair05}. ."
+ These variations causenotablechanges in absorption line features in the lensing galaxy (Muller&Guélin 2008).., These variations causenotablechanges in absorption line features in the lensing galaxy \citep{muller08}. .
+ The lens galaxy. is, The lens galaxy is
+et al.,et al.
+ showed that for the 6 non-detections in their DLA sample. the strength of the DDIB [which shows one of the tightest correlations with iin the MW] is often at least 3 times weaker in DLAs for a given ccompared with Galactic sight-lines.," showed that for the 6 non-detections in their DLA sample, the strength of the DIB [which shows one of the tightest correlations with in the MW] is often at least 3 times weaker in DLAs for a given compared with Galactic sight-lines."
+ The DDIB is even more under-abundant in DLAs for a givenD: j-10 times weaker than towards Galactic sight-lines., The DIB is even more under-abundant in DLAs for a given: 4–10 times weaker than towards Galactic sight-lines.
+ A similar result has been found for DIBs in the Large and Small Magellanic Cloud (LMCandSMC:Weltyetal.2006) where the DDIB is typically 10—30 times weaker than expected from the Galactic relation.," A similar result has been found for DIBs in the Large and Small Magellanic Cloud \citep[LMC and
+SMC;][]{WeltyD_06a} where the DIB is typically 10–30 times weaker than expected from the Galactic relation."
+ On the other hand. the DDIB strength correlates well with iin both Galactic and Magellanic Cloud. sight-lines. and. the detection towards AO | [64 also lies on the same relationship (Yorketal.2006a).," On the other hand, the DIB strength correlates well with in both Galactic and Magellanic Cloud sight-lines, and the detection towards AO $+$ 164 also lies on the same relationship \citep{YorkB_06a}."
+. These results hint that DIB formation/survival and high dust content are closely linked and that DIBs are therefore most likely to be detected in galaxies with high reddening., These results hint that DIB formation/survival and high dust content are closely linked and that DIBs are therefore most likely to be detected in galaxies with high reddening.
+ Wild.Hewett&Pettini(2006) have recently suggested that absorbers identified via high equivalent widths (EWs) of may select the highest aand highest aabsorbers., \citet*{WildV_06a} have recently suggested that absorbers identified via high equivalent widths (EWs) of may select the highest and highest absorbers.
+ For example. whereas DLAs have been constrained to have 0.04 (Murphy&Liske2004:Ellisonetal. 2005).. find that absorberswith A3934 EWs 70.7 hhave vvalues up to ~0.1 mag.," For example, whereas DLAs have been constrained to have $\le$ 0.04 \citep{MurphyM_04c,EllisonS_05a}, , \citet{WildV_06a} find that absorberswith $\lambda$ 3934 EWs $>$ have values up to $\sim$ 0.1 mag."
+ absorbers may therefore be promising sites for the detection of DIBs., absorbers may therefore be promising sites for the detection of DIBs.
+ Wild&Hewett(2005) presented a sample of 0.5<218 absorbers selected from the Sloan Digital Sky Survey (SDSS)., \citet{WildV_05a} presented a sample of $0.8 < z < 1.3$ absorbers selected from the Sloan Digital Sky Survey (SDSS).
+ However. the typical rest wavelengths of the strong DIB features (approximately AA)) makes the Wild&Hewett(2005) sample unsuitable for an optical search for the diffuse bands.," However, the typical rest wavelengths of the strong DIB features (approximately ) makes the \citet{WildV_05a} sample unsuitable for an optical search for the diffuse bands."
+ We have recently conducted an independent search for absorbers in the SDSS at 2<0.6 (seeZychetal.2007.fordetails) and found over 40 new absorbers., We have recently conducted an independent search for absorbers in the SDSS at $z<0.6$ \citep[see][for details]{ZychB_07a} and found over 40 new absorbers.
+ We selected 9 high-EW A3934 0.35 AA)) systems whose redshift places at least one of the strong DIBs (specitically. the 4428. 5705. 5780. 5797. 6284 and bbands were targeted) in regions of the spectrum free from atmospheric absorption and night sky emission.," We selected 9 high-EW $\lambda$ 3934 $\ge$ ) systems whose redshift places at least one of the strong DIBs (specifically, the 4428, 5705, 5780, 5797, 6284 and bands were targeted) in regions of the spectrum free from atmospheric absorption and night sky emission."
+ We note that due o the small impact parameters between the QSO and the galaxy causing absorption (typically<10kpe.e.g.Zychetal.2007) and he relatively large contribution by galactic light in the SDSS fibre. it is probable that some of the EW is contributed by galactic shotospheric absorption.," We note that due to the small impact parameters between the QSO and the galaxy causing absorption \citep[typically $<$ 10 kpc, e.g.][]{ZychB_07a} and the relatively large contribution by galactic light in the SDSS fibre, it is probable that some of the EW is contributed by galactic photospheric absorption."
+ The SDSS EWs may therefore not be an accurate measure of the interstellar content (this effect may be minimised when the EW can be measured from the FORS long slit spectrum: see Table 1)., The SDSS EWs may therefore not be an accurate measure of the interstellar content (this effect may be minimised when the EW can be measured from the FORS long slit spectrum; see Table \ref{dib_table}) ).
+ However. the small impact yurameters also mean that these galaxies are likely to produce significant absorption from their ISM in the QSO spectrum and the potential contamination from photospheric absorption does not alter the conclusions of this work [e.g. inferred aand Vl].," However, the small impact parameters also mean that these galaxies are likely to produce significant absorption from their ISM in the QSO spectrum and the potential contamination from photospheric absorption does not alter the conclusions of this work [e.g. inferred and ]."
+" Table 1. lists the in our sample. as well as their emission and absorption redshifts. 2,4, and 2,1. (the latter being derived from a Gaussian fit to the Ca ITA 3934 lline). and SDSS i-band magnitudes. n1,"," Table \ref{obs_table} lists the in our sample, as well as their emission and absorption redshifts, $z_{\rm em}$ and $z_{\rm
+ abs}$ (the latter being derived from a Gaussian fit to the Ca II $\lambda$ 3934 line), and SDSS $r$ -band magnitudes, $m_r$."
+ The 9 targets in Table |) were observed in long slit mode with the FORS2 spectrograph on the Very Large Telescope (VET) in Chile during ESO's Period 77 (1 April 2006 — 30 September 2006)., The 9 targets in Table \ref{obs_table} were observed in long slit mode with the FORS2 spectrograph on the Very Large Telescope (VLT) in Chile during ESO's Period 77 (1 April 2006 – 30 September 2006).
+ Observations were obtained through a | aresee slit with the CCD binned 2., Observations were obtained through a 1 arcsec slit with the CCD binned $\times$ 2.
+ Grism selection depended on absorber redshift: he exposure time. choice of grism and the resulting FWHM and signal-to-noise (S/N) ratios per pixel are listed in Table Ι..," Grism selection depended on absorber redshift; the exposure time, choice of grism and the resulting FWHM and signal-to-noise (S/N) ratios per pixel are listed in Table \ref{obs_table}."
+ The data reduction procedure followed standard steps for ong slit spectra using TRAF: a median bias frame was subtracted rom each science frame. followed by division by an average amp flat field.," The data reduction procedure followed standard steps for long slit spectra using IRAF: a median bias frame was subtracted from each science frame, followed by division by an average lamp flat field."
+ The spectra were optimally extracted. wavelength calibrated by use of a CuAr lamp and converted toa vacuum-reliocentric scale.," The spectra were optimally extracted, wavelength calibrated by use of a CuAr lamp and converted toa vacuum-heliocentric scale."
+ We experimented with different methods of combining individual exposures. including the usual SCOMBINE ask in IRAF with weightings according to S/N. and also using ppopler. as described in Zychetal.(2007)..," We experimented with different methods of combining individual exposures, including the usual SCOMBINE task in IRAF with weightings according to S/N, and also using popler, as described in \citet{ZychB_07a}."
+ Both gave very similar results., Both gave very similar results.
+ We searched the final spectra for absorption associated with the 4428. 5705. 5780. 5797. 6284 and ddiffuse bands.," We searched the final spectra for absorption associated with the 4428, 5705, 5780, 5797, 6284 and diffuse bands."
+ The first of these bands is intrinsically broad with an average (rest-frame) FWHM measured from 4 Galactic stellar sight-lines ofFWHM ~ (Jenniskens&Desert1994).., The first of these bands is intrinsically broad with an average (rest-frame) FWHM measured from 4 Galactic stellar sight-lines ofFWHM $\sim$ \citep{JenniskensP_94a}.
+ The other 5 DIBS are narrower. with FWHM values of —2.2. 2.1. 1.0. 2.6 and rrespectively in Galactic sight-lines Jenniskens&Desert1994) A.," The other 5 DIBS are narrower, with FWHM values of $\sim$ 2.2, 2.1, 1.0, 2.6 and respectively in Galactic sight-lines \citep{JenniskensP_94a} ."
+.. Comparing these values with the resolution of our spectra in Table |. it can be seen that the DDIB is always resolved in our spectra., Comparing these values with the resolution of our spectra in Table \ref{obs_table} it can be seen that the DIB is always resolved in our spectra.
+ Weusually do not resolve the narrower DIBs: taking into account the 1|2 broadening. the expected FWHM values of the DIBs isΑΑ.. compared with our typical spectral resolution ofAA.," Weusually do not resolve the narrower DIBs; taking into account the $1+z$ broadening, the expected FWHM values of the DIBs is, compared with our typical spectral resolution of."
+. Our search yielded one DIB detection: the bband at z=0.1556 towards 0024. the absorber with the highestapparent EW in our sample.," Our search yielded one DIB detection: the band at $z =
+0.1556$ towards $-$ 0024, the absorber with the highestapparent EW in our sample."
+ Figure 1. shows this detection together with the corresponding (from the SDSS spectrum) and absorption., Figure \ref{dib_fig} shows this detection together with the corresponding (from the SDSS spectrum) and absorption.
+ The redshifts of the 5780 DDIB and lines are in excellent agreement with the, The redshifts of the 5780 DIB and lines are in excellent agreement with the
+bar.,bar.
+ Once more there is a small density peak at the centre of the bar in the first frame (3a: Poeak7ferit ). which then quickly. disperses.," Once more there is a small density peak at the centre of the bar in the first frame (3a; $\rho\pea \simeq 
+\rho\cri$ ), which then quickly disperses."
+ The only significant. dillerence is a fanning-out of the ow close to the binary components. right at the end of he simulation.," The only significant difference is a fanning-out of the bar close to the binary components, right at the end of the simulation."
+ This is a real effect. due to the tidal shearing of the bar by the inspiralling binary components.," This is a real effect, due to the tidal shearing of the bar by the inspiralling binary components."
+" Lt only appears in this simulation because the evolution has »en followed⋅ to somewhat higher. density. (3..)10‘""gem5 instead of 210""gem. 7): hence the binary components iive gotten further out of alignment with the bar. where hey can exert a disruptive shear on the ends of the bar."," It only appears in this simulation because the evolution has been followed to somewhat higher density $3 \times 10^{-9} \, 
+{\rm g} \, {\rm cm}^{-3}$ instead of $2 \times 10^{-9} \, {\rm g} \, 
+{\rm cm}^{-3}$ ); hence the binary components have gotten further out of alignment with the bar, where they can exert a disruptive shear on the ends of the bar."
+ At the end. less than one fifth of the original particles rave been split. so there are ~ 140.000 particles in total.," At the end, less than one fifth of the original particles have been split, so there are $\sim$ 140,000 particles in total."
+ The simulation with On-Phe-Fly Splitting therefore requires less han one quarter the memory and less than of the CPU used. by the standard. simulation (with 600.000. particles. ERIT).," The simulation with On-The-Fly Splitting therefore requires less than one quarter the memory and less than of the CPU used by the standard simulation (with 600,000 particles, 8.1)."
+ lt appears that Particle Splitting is a viable. option for increasing the resolution. locally. during an SPIL simulation of self-gravitating collapse.," It appears that Particle Splitting is a viable option for increasing the resolution, locally, during an SPH simulation of self-gravitating collapse."
+ In this context. the most stringent and. apposite test. of. the algorithm is the BB79 test. where extra resolution is required to avoicl violating the Jeans Condition (and hence to avoid artificial. fragmentation).," In this context, the most stringent and apposite test of the algorithm is the BB79 test, where extra resolution is required to avoid violating the Jeans Condition (and hence to avoid artificial fragmentation)."
+ Acceptable results are. obtained for the DD79 test. even when the ogas is programmed to stay isothermal up lo peg5107 g *.," Acceptable results are obtained for the BB79 test, even when the gas is programmed to stay isothermal up to $\rho\cri = 5 \times 10^{-12}$ g $^{-3}$."
+ Perturbations due to the splitting process are transient. ancl are clleetively ancl quickly damped.," Perturbations due to the splitting process are transient, and are effectively and quickly damped."
+ Moreover. the simulation is accomplished. with much. less memory. ancl CPU than a standard simulation. particularly with the On-'The-Ely. implementation.," Moreover, the simulation is accomplished with much less memory and CPU than a standard simulation, particularly with the On-The-Fly implementation."
+ We note also that the savings in memory and CPU are likely to be far greater in realistic simulations of star formation. where only a fraction of the matter in a cloud. collapses to form stars. and therefore a [ar smaller fraction of particles needs to be split and to be followed with a small time-step.," We note also that the savings in memory and CPU are likely to be far greater in realistic simulations of star formation, where only a fraction of the matter in a cloud collapses to form stars, and therefore a far smaller fraction of particles needs to be split and to be followed with a small time-step."
+ Other tests have been performed using the two Particle Splitting algorithms., Other tests have been performed using the two Particle Splitting algorithms.
+ In particular we note the following. (, In particular we note the following. (
+i) We have evolved a static. stable isothermal sphere. contained by an external pressure 2Z4. and shown (a) that it remains stable. and (b) that its density. profile is not. corrupted.,"i) We have evolved a static, stable isothermal sphere, contained by an external pressure $P_{\mbox{\small ext}}$, and shown (a) that it remains stable, and (b) that its density profile is not corrupted."
+" The configuration treated has mass M... radius 0.01. pc. and temperature 7.9 Ix. H is quite centrally condensed. with Pecntral3edge and Lost230LO"" erg 7. ("," The configuration treated has mass $M_\odot$, radius $0.01$ pc, and temperature 7.9 K. It is quite centrally condensed, with $\rho_{\mbox{\small central}} \simeq 3 \, 
+\rho_{\mbox{\small edge}}$ and $P_{\mbox{\small ext}} \simeq 
+3 \times 10^{-12}$ erg $^{-3}$. ("
+ii) We have also followed the collapse of an isothermal cloud which initially bas uniform density. and. quickly evolves towards an inverse-square density profile (pxr 2).,"ii) We have also followed the collapse of an isothermal cloud which initially has uniform density, and quickly evolves towards an inverse-square density profile $\rho \propto 
+r^{-2}$ )."
+ Phe cloud. has mass AJ. and temperature 7.9 Ix. Initially it has raclius 0.016 pe and uniform density ~45100D & 7.," The cloud has mass $M_\odot$ and temperature 7.9 K. Initially it has radius 0.016 pc and uniform density $\sim 
+4 \times 10^{-18}$ g $^{-3}$."
+ Ehe elfect of Particle Splitting on the ensuing collapse is negligible., The effect of Particle Splitting on the ensuing collapse is negligible.
+ Particle Splitting might be applied in other situations where increased resolution is required. locally., Particle Splitting might be applied in other situations where increased resolution is required locally.
+ Ehe Nested implementation requires prior knowledge of where extra resolution will be required., The Nested implementation requires prior knowledge of where extra resolution will be required.
+ The On-The-FIy implementation requires that the condition for requiring extra resolution can be formulated as a local function of state., The On-The-Fly implementation requires that the condition for requiring extra resolution can be formulated as a local function of state.
+ However. more," However, more"
+uncertainty on a statistical (see e.g. Cordes Lazio 2003). while the errors on (he supershell distances are specifically indicated in (he MeClIure-Griffiths et al.,"uncertainty on a statistical (see e.g. Cordes Lazio 2003), while the errors on the supershell distances are specifically indicated in the McClure-Griffiths et al."
+ sample. aud most of them are in the range 10-20%.," sample, and most of them are in the range 10-20."
+. Unfortunately no distance errors are provided in the Ieiles (1979) sample. ancl for those we assume a more conservative uncertainty.," Unfortunately no distance errors are provided in the Heiles (1979) sample, and for those we assume a more conservative uncertainty."
+ We (hen consider a pulsar a possible association with the supershell if it falls within DxR alter accounting for the distance uncertainties in both the pulsar and the supershell., We then consider a pulsar a possible association with the supershell if it falls within $D\pm R$ after accounting for the distance uncertainties in both the pulsar and the supershell.
+ As Table | shows. in a few cases. due to a combination of the large distances of the supershells and (he comparatively poor sensilivily level αἱ which the corresponding region was surveved. no pulsars are expected to have been detected in association with the supershell. ancl therefore no meaningful constraints can be derived.," As Table 1 shows, in a few cases, due to a combination of the large distances of the supershells and the comparatively poor sensitivity level at which the corresponding region was surveyed, no pulsars are expected to have been detected in association with the supershell, and therefore no meaningful constraints can be derived."
+ For the (wo supershells which lie in the region covered by the Parkes A\lultibeam survey (GSII 285024-86. and GSI 29201455: MeChire-Grilliths et al., For the two supershells which lie in the region covered by the Parkes Multibeam survey (GSH 285–02+86 and GSH 292–01+55; McClure-Griffiths et al.
+ 2002). the predictions of the multiple SN scenario al the detection level of that survey are consistent with the number of cancidate associations. although they do not rule out other interpretations. such as GIDs or collisions with clouds. as in both these scenarios one would not expect any pulsar enhancement correlated wilh the supershell.," 2002), the predictions of the multiple SN scenario at the detection level of that survey are consistent with the number of candidate associations, although they do not rule out other interpretations, such as GRBs or collisions with high-velocity clouds, as in both these scenarios one would not expect any pulsar enhancement correlated with the supershell."
+ Deeper surveys of (hose regions are needed to test the multiple SN scenario., Deeper surveys of those regions are needed to test the multiple SN scenario.
+ Ilowever. an interesting constraint can already be derived for GS1I 242034-37 (lleiles 1979).," However, an interesting constraint can already be derived for GSH 242–03+37 (Heiles 1979)."
+ Our results (see Table 1) show that there is a probability that the supershell was not the result of multiple SNe., Our results (see Table 1) show that there is a probability that the supershell was not the result of multiple SNe.
+ Deeper survevs are necessary {ο set much lighter constraints., Deeper surveys are necessary to set much tighter constraints.
+ Our current. results. in fact. are rather dependent on the luminosity [function (Lorimer et al.," Our current results, in fact, are rather dependent on the luminosity function (Lorimer et al."
+ 1993) that we have adopted here: however. recent new discoveries of low-humninosityv radio pulsars (Camilo et al.," 1993) that we have adopted here; however, recent new discoveries of low-luminosity radio pulsars (Camilo et al."
+ 2002a. 20025) might indicate (hat voung pulsars might be fainter (han previously realized. and (his could be an alternate explanation [or the lack of pulsars in the supershell at the current sensitivity threshold.," 2002a, 2002b) might indicate that young pulsars might be fainter than previously realized, and this could be an alternate explanation for the lack of pulsars in the supershell at the current sensitivity threshold."
+ However. when the sensitivity is sulliciently high that even the faintest pulsars can be detected (as it will be with ihe planned SIXÀ instrument). then the particular details of the Iuminositv finetion will be irrelevant [or the proposed experiment.," However, when the sensitivity is sufficiently high that even the faintest pulsars can be detected (as it will be with the planned SKA instrument), then the particular details of the luminosity function will be irrelevant for the proposed experiment."
+ Also note that. while current pulsar survevs are able to probe mostly the pulsars within our Galaxy. the SIVA survey will be able to detect a large fraction of pulsars also in the LMC. which has several giant supershells (kim et al.," Also note that, while current pulsar surveys are able to probe mostly the pulsars within our Galaxy, the SKA survey will be able to detect a large fraction of pulsars also in the LMC, which has several giant supershells (Kim et al."
+ 1999)., 1999).
+ The analvsis (hal we are proposing here can therelore be mace also [ον the LAIC supershells in the SINA era., The analysis that we are proposing here can therefore be made also for the LMC supershells in the SKA era.
+ Belore concluding we should note that. besides depending on the total," Before concluding we should note that, besides depending on the total"
+et al. (,et al. (
+2000) consider the breakup of a horizontal maguetic laver iuto a nunber of flux tubes. thereby including the effect of rotation.,"2000) consider the breakup of a horizontal magnetic layer into a number of flux tubes, thereby including the effect of rotation."
+ For an iniposed maenetic field in the vertical directio- convection is increasingly hampered with increasing field streneth. as is iudicated by a decrease of the riis turbulent velocity (Fig.1," For an imposed magnetic field in the vertical direction, convection is increasingly hampered with increasing field strength, as is indicated by a decrease of the rms turbulent velocity (Fig.,"
+0. top) aud of the couvective cucrey flix., top) and of the convective energy flux.
+ Since the mmaenetic energy is cuhanced by advection of the iuposed field. while the kinetic cucrey decreases as a result of the Lorentz force. after a while the rims magnetic field can be much larger than the equipartition feld streneth (Fig.1," Since the magnetic energy is enhanced by advection of the imposed field, while the kinetic energy decreases as a result of the Lorentz force, after a while the rms magnetic field can be much larger than the equipartition field strength (Fig.,"
+0. bottom)., bottom).
+ Moreover. if the iuifial magnetic field is too strong. then convection dies out completely.," Moreover, if the initial magnetic field is too strong, then convection dies out completely."
+" This iu fact occured dunug run Te (Bo/Ba,z 1.7).", This in fact occured during run 7e $B_0/B_{\rm eq}\approx 1.7$ ).
+ After 100 time uuits. the volunc-averaged cherey density of the velocity and maeuetic-field perturbations had both fallen byabout 2 orders of magnitude.," After 100 time units, the volume-averaged energy density of the velocity and magnetic-field perturbations had both fallen byabout $2$ orders of magnitude."
+ The decrease of ay is more rapid han that of tuys. because advection of the imposed field w convection is stifled: this explains why alpha decreases also when measured in dvuamical units (.e.. when divided N Unas).," The decrease of $\alpha_{\rm V}$ is more rapid than that of $u_{\rm rms}$, because advection of the imposed field by convection is stifled; this explains why alpha decreases also when measured in dynamical units (i.e., when divided by $u_{\rms}$ )."
+ A debate is ongoing iu the literature about low and whether the magnetic Revnolds wuuber affectsthe naenetic quenching of the à. effect., A debate is ongoing in the literature about how and whether the magnetic Reynolds number affectsthe magnetic quenching of the $\alpha$ effect.
+ The dependence of o ou the maguetic field strength is often schematically represented in the form amoogl|Rin’B?/Be)+., The dependence of $\alpha$ on the magnetic field strength is often schematically represented in the form $\alpha\approx\alpha_0[1+\mb{Rm}^p B^2/B_{\rm eq}^2]^{-1}$.
+ Wile some muimerical and analytical results sugeest that p0(Usraichnuan1979a.1979b. Draudeuburg & Douner 1997). others sugeest pzc1 (Cattaneo & Vainshtein 1991. Vainshtein & Cattaneo 1992. Tao et al. 1," While some numerical and analytical results suggest that $p\approx 0$(Kraichnan, Brandenburg & Donner ), others suggest $p\approx 1$ (Cattaneo & Vainshtein , Vainshtein & Cattaneo , Tao et al. ,"
+ Wile some muimerical and analytical results sugeest that p0(Usraichnuan1979a.1979b. Draudeuburg & Douner 1997). others sugeest pzc1 (Cattaneo & Vainshtein 1991. Vainshtein & Cattaneo 1992. Tao et al. 19," While some numerical and analytical results suggest that $p\approx 0$(Kraichnan, Brandenburg & Donner ), others suggest $p\approx 1$ (Cattaneo & Vainshtein , Vainshtein & Cattaneo , Tao et al. ,"
+ Wile some muimerical and analytical results sugeest that p0(Usraichnuan1979a.1979b. Draudeuburg & Douner 1997). others sugeest pzc1 (Cattaneo & Vainshtein 1991. Vainshtein & Cattaneo 1992. Tao et al. 199," While some numerical and analytical results suggest that $p\approx 0$(Kraichnan, Brandenburg & Donner ), others suggest $p\approx 1$ (Cattaneo & Vainshtein , Vainshtein & Cattaneo , Tao et al. ,"
+ Wile some muimerical and analytical results sugeest that p0(Usraichnuan1979a.1979b. Draudeuburg & Douner 1997). others sugeest pzc1 (Cattaneo & Vainshtein 1991. Vainshtein & Cattaneo 1992. Tao et al. 1993," While some numerical and analytical results suggest that $p\approx 0$(Kraichnan, Brandenburg & Donner ), others suggest $p\approx 1$ (Cattaneo & Vainshtein , Vainshtein & Cattaneo , Tao et al. ,"
+ Wile some muimerical and analytical results sugeest that p0(Usraichnuan1979a.1979b. Draudeuburg & Douner 1997). others sugeest pzc1 (Cattaneo & Vainshtein 1991. Vainshtein & Cattaneo 1992. Tao et al. 1993.," While some numerical and analytical results suggest that $p\approx 0$(Kraichnan, Brandenburg & Donner ), others suggest $p\approx 1$ (Cattaneo & Vainshtein , Vainshtein & Cattaneo , Tao et al. ,"
+ deVaucouleurs(1961)... Corredoiraetal.2005).. (Eughuaier1999:Fux1999)," \citet{DeVaucouleurs64}. \citep{Dwek+95,Binney+97}, \citep{Stanek+97,LopezCorredoira+05}, \citep{Englmaier+Gerhard99,Fux99},"
+... (Ibunadache2006) shear (Dehuen2000:Minchevetal.," \citep{Hamadache+06} \citep{Dehnen00, Minchev+07}."
+"2007).. ~1.5 (e.g.Binneyetal.L997) (Athanassoula2005).. (e.g.€'onibesetRahaeal.1991).. (ο,ο,Lyuden-Be]&WwalnajsMagztinez-Valpuestaetal."," $\sim 1.5$ \citep[e.g.,][]{Binney+97} \citep{Athanassoula05}, \citep[e.g.,][]{Combes+90,Raha+91}. \citep[e.g.][MV06]{Lynden-Bell+Kalnajs72, Athanassoula03,
+ Debattista+Sellwood00, Martinez-Valpuesta+06}."
+2006.NIVOG6).. à—15730° (Cerhard2002).. (Deujunin Tanunersleyetal.(2000):," $\alpha\sim15^\circ-30^\circ$ \citep{Gerhard02}. \citep[][hereafter B05,C07,C08,C09]{Benjamin+05,
+ Cabrera-Lavers+07b,Cabrera-Lavers+08,Churchwell+09}, \citet{Hammersley+00}:"
+ ofa!~~13°. |o217 to d—107 (L5 pe to 1.5 kpc ealactoceutric radius). aud therefore coexisti1ο with the conventional Galacticbar over a range of radi.," $\alpha'\sim 43^\circ$ $l\simeq 27^\circ$ to $l\simeq 10^\circ$ $4.5$ kpc to $1.5$ kpc galactocentric radius), and therefore coexisting with the conventional Galactic over a range of radii."
+ This interpretation. 1f correct. would dynamically be quite puzzliie: two separate rotating bars should aligi with each other through cwnamical coupling iu at mos a few rotatiou periods," This interpretation, if correct, would dynamically be quite puzzling: two separate rotating bars should align with each other through dynamical coupling in at most a few rotation periods."
+ Iutus letter we show that a separate inferredbar does not necessariv follow from the star cout data. aud we suggest a plauxible model to explain these data with a siuele barred structure whose iuner parts represent the boxy bulge.," In this letter we show that a separate inferred does not necessarily follow from the star count data, and we suggest a plausible model to explain these data with a single barred structure whose inner parts represent the boxy bulge."
+ We also show some predictions for radial volociv distributions that could be used to test this model in the necw future., We also show some predictions for radial velocity distributions that could be used to test this model in the near future.
+" The smiulation used iu tus work is similar to that published iun MX""06 aud has uot been run to match the ΑΣ structure.", The simulation used in this work is similar to that published in MV06 and has not been run to match the MW structure.
+ The code used is FTAL tL (updated version) from 1Jeller&Shlosmman(1991)., The code used is FTM 4.4 (updated version) from \citet{Heller+Shlosman94}.
+. The total uunmnber of partic desis d&109. distributed initially iu an exponential disk with Q=1.5. ciumbedded in a live dark matter halo.," The total number of particles is $1\times10^6$, distributed initially in an exponential disk with $Q=1.5$, embedded in a live dark matter halo."
+ Αςy 1.5 Cr the bar becomes very strong and buckles. thereby. weakening.," After $\sim 1.5$ Gyr the bar becomes very strong and buckles, thereby weakening."
+ Laer the bar resumes its evolution auxL erows again. resullug in a τοΙολ! ar structure.," Later the bar resumes its evolution and grows again, resulting in a prominent structure."
+ We cousider tje simnulated galaxy at time ~1.9 Cor. after the boxy biilee has formed and the bar has regrown.," We consider the simulated galaxy at time $\sim 1.9$ Gyr, after the boxy bulge has formed and the bar has regrown."
+" The density disrbutiou for this snapshot is shown iu Figure Le. oricuted at au anele a=257?5"" with respect to the line from the Calactic center to the observer."," The density distribution for this snapshot is shown in Figure \ref{fig:snapshot1}{, oriented at an angle $\alpha =
+25^\circ$ with respect to the line from the Galactic center to the observer."
+ The boxy bulge is aparent in Figure Lb., The boxy bulge is apparent in Figure \ref{fig:snapshot1}{.
+ The model is scaled so thatthe eud of the planar bar appears just inside longitude |=3Fo as seen fromtheobserver., The model is scaled so thatthe end of the planar bar appears just inside longitude $l=30^\circ$ as seen fromtheobserver.
+The bar lenethis~L5 kpc. aud the maxima ellipticity is 0.16.,"The bar length is$\sim4.5$ kpc, and the maximum ellipticity is $0.46$ ."
+ Relative tothe pauar bar. theboxybulgeis ~20% larger in? in the scaled model than in the MW. as measured by," Relative tothe planar bar, theboxybulgeis $\sim20\%$ larger in $l$ in the scaled model than in the MW, as measured by"
+one deep exposure in the aarchive. wilh a second shorter observation (see Table 3)).,"one deep exposure in the archive, with a second shorter observation (see Table \ref{tab:prop}) )."
+ The second shorter exposure of AI83 was taken 16 months after the first. and the second exposure of MIOI taken 7 months alter the first.," The second shorter exposure of M83 was taken 16 months after the first, and the second exposure of M101 taken 7 months after the first."
+ In the long (100ks) pointing of ALLOL there are 28 soft sources (potential SNRs) detected., In the long (100ks) pointing of M101 there are 28 soft sources (potential SNRs) detected.
+ Most of these have 20-100 counts and are not detected in the short (10ks) observation., Most of these have 20-100 counts and are not detected in the short (10ks) observation.
+ This does not provide verv stringent constraints on the variabilitv of the solt sources. since to be detected in the short observation they. would have to increase in [αν by factors of 5-10.," This does not provide very stringent constraints on the variability of the soft sources, since to be detected in the short observation they would have to increase in flux by factors of 5-10."
+" There are two ""soft"" sources detected in both the long and short observation.", There are two “soft” sources detected in both the long and short observation.
+ One source with //1=—0.8 shows no evidence for variability., One source with $H1=-0.8$ shows no evidence for variability.
+ Another source is significantly variable. and has a //[1=—0.78 in the long observation.," Another source is significantly variable, and has a $H1=-0.78$ in the long observation."
+" This source has a bumninositv of ~LO"" ere lI. and is the brightest object in MIOL (source 98 [from Penceetal(2001)))."," This source has a luminosity of $\sim10^{39}$ erg $^{-1}$, and is the brightest object in M101 (source 98 from \citet{pence01}) )."
+ Il is highly variable and is clearly an accretion source., It is highly variable and is clearly an accretion source.
+ Two “hard” (possible ILMXD) sources are detected in both observations in MIOI., Two “hard” (possible HMXB) sources are detected in both observations in M101.
+ One has //2=0.31 in the long observation and shows no evidence for variabilitv., One has $H2=0.31$ in the long observation and shows no evidence for variability.
+ One source has //2=0.48. and the flux approximately doubles in the second observation.," One source has $H2=0.48$, and the flux approximately doubles in the second observation."
+ There are 28 soft (possible SNR) sources detected in M82 in the long (50ks) observation., There are 28 soft (possible SNR) sources detected in M83 in the long (50ks) observation.
+ A total of 7 sources with soft colors that were detected in the long observation were also detected in the short observation., A total of 7 sources with soft colors that were detected in the long observation were also detected in the short observation.
+ Two of these are clearly variable. while the remaining 5 have approximately constant flix.," Two of these are clearly variable, while the remaining 5 have approximately constant flux."
+ Of the 28 soft sources detected in the long observation. 5 sources should have been detected in the short (10ks) observation. assuming no variability.," Of the 28 soft sources detected in the long observation, 5 sources should have been detected in the short (10ks) observation, assuming no variability."
+ Three of these have extremely soft colors. are clearly variable and are probably supersolt sources.," Three of these have extremely soft colors, are clearly variable and are probably supersoft sources."
+ The remaining (wo sources may have declined in fInx between the two observations., The remaining two sources may have declined in flux between the two observations.
+ Three hard (potentially WAINB sources) were detected in the long observation of M53., Three hard (potentially HMXB sources) were detected in the long observation of M83.
+ One of (hese sources may be variable., One of these sources may be variable.
+ We conclude that we have insullicient data to make definitive statements about the variability properties of (he soft sources in either M83 or MIOI., We conclude that we have insufficient data to make definitive statements about the variability properties of the soft sources in either M83 or M101.
+" The spatial distribution of the soft (vellow Ns) and ""LMXD"" (red Ns) sources is shown in Fiewwe 8 lor M83.", The spatial distribution of the soft (yellow X's) and “LMXB” (red X's) sources is shown in Figure \ref{fig:opt_and_x-ray} for M83.
+ The AA-rav image is shown on the left. and the U-band optical image on the right.," The X-ray image is shown on the left, and the U-band optical image on the right."
+" Both “LAINB"" and soft sources have a tendency {ο occur in the spiral arms.", Both “LMXB” and soft sources have a tendency to occur in the spiral arms.
+ The main difference between the two distributions. however. is that the nuclear X-ray. sources are almost all sources with “LAINB™ colors.," The main difference between the two distributions, however, is that the nuclear X-ray sources are almost all sources with “LMXB” colors."
+ In MIOLI the X-ray. sources follow a similar pattern (see also (2001))) with both tvpes of sources following the spiral structure., In M101 the X-ray sources follow a similar pattern (see also \citet{pence01}) ) with both types of sources following the spiral structure.
+ There is only a single diffuse source al the center of MIOI (there is essentially no bulge in this svstem) so segregation of central sources is nol seen., There is only a single diffuse source at the center of M101 (there is essentially no bulge in this system) so segregation of central sources is not seen.
+ellipses ou the Schechter parameters are shown in Figures and 10..,ellipses on the Schechter parameters are shown in Figures \ref{fig-jenv_sty} and \ref{fig-env_sty}.
+ The Schechter parameters of the J band cluster LF are inconsistent with those of the field LF at the confidence level. as determined from the likelihood. contours.," The Schechter parameters of the $J-$ band cluster LF are inconsistent with those of the field LF at the confidence level, as determined from the likelihood contours."
+ The trend is more significant in the J) baud due o the larger sample size as a consequence of the deeper photometry., The trend is more significant in the $J-$ band due to the larger sample size as a consequence of the deeper photometry.
+ The difference is in the seuse that the cluster LF has a brighter and a steeper oj: however. giveu the degeneracy between lese paranueters if is nof clear precisely how the LF shape varies.," The difference is in the sense that the cluster LF has a brighter and a steeper $\alpha_J$; however, given the degeneracy between these parameters it is not clear precisely how the LF shape varies."
+ If we restrict the Schechter fit to a Linited uuimositv rauge (e... brighter than A;=20.5. or fainter than A;= 23.5). the best-fit parameters change along the major axis of the error cllipse.," If we restrict the Schechter fit to a limited luminosity range (e.g., brighter than $M_J=-20.5$, or fainter than $M_J=-23.5$ ), the best-fit parameters change along the major axis of the error ellipse."
+ However. voth the cluster and feld paramcters move iu the sune direction. and the difference between them is approximately maintained.," However, both the cluster and field parameters move in the same direction, and the difference between them is approximately maintained."
+ It appears. therefore. that the environmental depeudeuce of the LF is not restricted to just the bright or faint cud.," It appears, therefore, that the environmental dependence of the LF is not restricted to just the bright or faint end."
+ The differences among the LEs of different euvironnieuts are huger when only the NEL galaxies are cousidered. as shown in Figure 9..," The differences among the LFs of different environments are larger when only the NEL galaxies are considered, as shown in Figure \ref{fig-jenv_sty}."
+ In the field. the LF of NEL galaxies is much fatter than that of the full sample. but in clusters the LFs of EL aud NEL are statistically inclistineuishable.," In the field, the LF of NEL galaxies is much flatter than that of the full sample, but in clusters the LFs of EL and NEL are statistically indistinguishable."
+ We discuss this result further in retsec-spec.., We discuss this result further in \\ref{sec-spec}.
+ The Schechter paraicters of the J aud A. huuaiuosityv functions. for cach data subsample. are tabulated in Table ]," The Schechter parameters of the $J$ and $K_s$ luminosity functions, for each data subsample, are tabulated in Table \ref{tab-params}."
+ It is encouragiug that the global ήπιο]τν functions measured here. based on redshifts frou the LCRS. are in such eood agreement with recent determinations by Cole ((2)) and ISochanek ((2)).," It is encouraging that the global luminosity functions measured here, based on redshifts from the LCRS, are in such good agreement with recent determinations by Cole \shortcite{Cole-2mass}) ) and Kochanek \shortcite{Kochanek-KLF}) )."
+ As shown by those authors. all infrared LFs derived from large samples are consistent (e.g. ?:77:7T:778 PY).," As shown by those authors, all infrared LFs derived from large samples are consistent (e.g., \cite{MSE,GPMC,Gardner,SSCM,Loveday-klf}) )."
+ Thus concordance is in contrast with the ustorical situation at optical wavelengths. where there las been considerable disagreement between results in he literature.," This concordance is in contrast with the historical situation at optical wavelengths, where there has been considerable disagreement between results in the literature."
+ In particular. optical LEs derived from he LORS (?: ?7)) are eenerally shallower than most other measurements. aud it has receutlv been shown that lis difference results from an underestimation of galaxy uagenitudes due to the shallow isoploal luit of the LORS photometry aud from the exclusion o flow ceutral surface xiehtuess galaxies (Blanton 7?) ).," In particular, optical LFs derived from the LCRS \cite{L+96,Christ}) ) are generally shallower than most other measurements, and it has recently been shown that this difference results from an underestimation of galaxy magnitudes due to the shallow isophotal limit of the LCRS photometry and from the exclusion of low central surface brightness galaxies (Blanton \shortcite{Sloan_lf}) )."
+ The fact that we do uot find a shallow slope in the iifrared. LF supports he interpretation that. at least over the relatively narrow. wielt magnitude range of 2MASS. it is the R magnitudes. and uot the survey selection itself. that is responsible for he discrepancy in the optical.," The fact that we do not find a shallow slope in the infrared LF supports the interpretation that, at least over the relatively narrow, bright magnitude range of 2MASS, it is the $R$ magnitudes, and not the survey selection itself, that is responsible for the discrepancy in the optical."
+ Analysis of the optical LF dependence ou euviromneut was done bv Christlei (7)). using the LORS sample.," Analysis of the optical LF dependence on environment was done by Christlein \shortcite{Christ}) ), using the LCRS sample."
+" Although the LORS iaenitucles are systematically biased due to the right isophotal Ενα, this bias applies uuiformly to the whole sample aud is therefore uulikely to affect the trends discussed in that paper."," Although the LCRS magnitudes are systematically biased due to the bright isophotal limit, this bias applies uniformly to the whole sample and is therefore unlikely to affect the trends discussed in that paper."
+ Clhiistleiu fouud strong evolution of the LF with cuviromuecnt. with the Schechter parameters changing along the major axis of the ellipse. toward steeper à and brighter M with larger o.," Christlein found strong evolution of the LF with environment, with the Schechter parameters changing along the major axis of the ellipse, toward steeper $\alpha$ and brighter $M^\ast$ with larger $\sigma$."
+ The magnitude aud sense of this chanee are fully cousisteut with the results preseuted in refsec-env.., The magnitude and sense of this change are fully consistent with the results presented in \\ref{sec-env}.
+ The fact that Christlein finds such a continuous depeudeuce ou 6 (something we are unable to investigate due to our smaller sample size) provides encourage support that the effect is real., The fact that Christlein finds such a continuous dependence on $\sigma$ (something we are unable to investigate due to our smaller sample size) provides encouraging support that the effect is real.
+ Zabludoff ADulehaeyv (2)) found similar evidence for steep. bright A- baud LEs in X-ray selected eroups (a=1.3).," Zabludoff Mulchaey \shortcite{ZM00}) ) found similar evidence for steep, bright $R$ -band LFs in X-ray selected groups $\alpha=-1.3$ )."
+" Because these groups generally have AD>1 keV. may would be considered ""clusters"" by our criteria. aud them results for steeper optical LFs iu these svstenis are very similar to our results for the cluster infrared LPs."," Because these groups generally have $kT\gtrsim 1$ keV, many would be considered “clusters” by our criteria, and their results for steeper optical LFs in these systems are very similar to our results for the cluster infrared LFs."
+ Previously published evidence for an cuvirommental dependence of the infrared LF is not as strong as for the optical LF., Previously published evidence for an environmental dependence of the infrared LF is not as strong as for the optical LF.
+ Both De Propris ((2)) and Andreou Pelló (7)) fud a steep faint cud slope in the 77 baud LE of the Coma cluster. but ouly at maguitudes faiuter than our sample probes.," Both De Propris \shortcite{Roberto98}) ) and Andreon Pelló \shortcite{AP00}) ) find a steep faint end slope in the $H$ band LF of the Coma cluster, but only at magnitudes fainter than our sample probes."
+" Trentham Mobasher (7)) find no evidence for a difference between the A-baud LEs of cluster and field galaxies. over 21<My«339, although the wucertaintics are large,"," Trentham Mobasher \shortcite{TMo98}) ) find no evidence for a difference between the $K$ -band LFs of cluster and field galaxies, over $-24<M_K<-22$, although the uncertainties are large."
+ Our sample is the first that is large chough to detect this weak effect at these huninosities., Our sample is the first that is large enough to detect this weak effect at these luminosities.
+ We look forward to dramatically inproviug the precision of this result whe the full 2\LASS catalog is released., We look forward to dramatically improving the precision of this result when the full 2MASS catalog is released.
+ Do the observed dependencies of the infrared LF on environment and spectral type necessarily imply au analogousdifference in the stellar mass functions?, Do the observed dependencies of the infrared LF on environment and spectral type necessarily imply an analogousdifference in the stellar mass functions?
+ The infrared light is a good. but not perfect tracer of the stellar nass the mass-to-Heht ratio (M/L) can be expected to increase by about a factor of two between earlv- aud late-type galaxies (e.g.. 2)).," The infrared light is a good, but not perfect tracer of the stellar mass; the mass-to-light ratio (M/L) can be expected to increase by about a factor of two between early- and late-type galaxies (e.g., \cite{BdJ_ML}) )."
+ To estimate M/L. constraints on the mean stellar age are required.," To estimate M/L, constraints on the mean stellar age are required."
+ The optical-IR. colors of our sanple are uureliable due to the shallow isophotal limit of the LORS. aud the IR colors alone are insensitive to stellar population differences.," The optical-IR colors of our sample are unreliable due to the shallow isophotal limit of the LCRS, and the IR colors alone are insensitive to stellar population differences."
+" The best iudicator available to us is therefore the Di999 spectral iudex frou the LCRS spectra,", The best indicator available to us is therefore the $D_{4000}$ spectral index from the LCRS spectra.
+ We use the solaranetallicitv. models of Fioc Rocca-Volinerauge (2)) to generate spectra frou a giveu star formation history., We use the solar-metallicity models of Fioc Rocca-Volmerange \shortcite{PEGASE}) ) to generate spectra from a given star formation history.
+ We then compute Doy. directly from the spectra. using the same definition as Zabludoff (24).," We then compute $D_{4000}$ directly from the spectra, using the same definition as Zabludoff \shortcite{Z+96}) )."
+ The correlation between ML; and Djjj depends ouly weakly on star formation history., The correlation between $M/L_J$ and $D_{4000}$ depends only weakly on star formation history.
+ Figure 11 shows model results for a single burst. simple stellar population (SSP) aud for a ealaxy with a constant star formation rate. using a Ikeunicutt (2)) initial mass function.," Figure \ref{fig-MLR} shows model results for a single burst, simple stellar population (SSP) and for a galaxy with a constant star formation rate, using a Kennicutt \shortcite{K83}) ) initial mass function."
+" The two models can differ bv a sauuchasθέ, over a narrow rauee of Dioo. but the ransine with M/L is unchanged."," The two models can differ by as much as, over a narrow range of $D_{4000}$, but the ranking with $M/L$ is unchanged."
+ Dyoy is sensitive to mctallicitv aud reddening. which we neglect. but is less sensitive than optical colors.," $D_{4000}$ is sensitive to metallicity and reddening, which we neglect, but is less sensitive than optical colors."
+ For example. an extinction of ly=1.5 magnitudes only reddeus Digo) by about 0.2.," For example, an extinction of $A_{\rm V}=1.5$ magnitudes only reddens $D_{4000}$ by about 0.2."
+ Iu Figure 12. we show the aud error cllipses ou the Schechter parameters for the stellar wiass functious derived using the M/L corrections derived from the SSP aodel., In Figure \ref{fig-mfunc} we show the and error ellipses on the Schechter parameters for the stellar mass functions derived using the M/L corrections derived from the SSP model.
+ The differences observed in the huuinositv functions are eracerbated m the stellar mass functions., The differences observed in the luminosity functions are in the stellar mass functions.
+ Thus. we conclude that the stellar mass function is uot," Thus, we conclude that the stellar mass function is not"
+anisotropic driving results in a weakly anisotropic shock.,anisotropic driving results in a weakly anisotropic shock.
+" Soon after the shock reaches the edge of the core at to and enters the steep density gradient in the envelope, it becomes unstable and develops a collimated jet, see Figs. 2-- 3.."," Soon after the shock reaches the edge of the core at $t_0$ and enters the steep density gradient in the envelope, it becomes unstable and develops a collimated jet, see Figs. \ref{shocks-shape-1}- - \ref{shocks-shape}."
+" For w>>4, the jets develops on the time scale to."," For $\om \gg 4$, the jets develops on the time scale $t_0$."
+" Secondly, we consider an isotropic driver, f—1 but starting with a non-spherical initial shape, r(t=0)1+cos?0."," Secondly, we consider an isotropic driver, $f=1$ but starting with a non-spherical initial shape, $r(t=0)= 1 + \cos^2 \theta$."
+" In this case the GRB shock dynamics is qualitatively similar to the wind: the shock remains weakly anisotropic in the core and develops a jet soon after reaching the envelope, Figs. 2-- 3.."," In this case the GRB shock dynamics is qualitatively similar to the non-spherical wind: the shock remains weakly anisotropic in the core and develops a jet soon after reaching the envelope, Figs. \ref{shocks-shape-1}- - \ref{shocks-shape}. ."
+"to luminosity is immediate and roughly constant in time, as assumed in our model, we would conclude that a low-density region must have existed between the star and out from the star.","to luminosity is immediate and roughly constant in time, as assumed in our model, we would conclude that a low-density region must have existed between the star and out from the star."
+" For example, as a stellar system transitions from an AGB star to a nebula (PPN), it changes from emitting a denser, cooler wind, to a hotter, less dense wind (Kwok1993)."," For example, as a stellar system transitions from an AGB star to a proto-planetary nebula (PPN), it changes from emitting a denser, cooler wind, to a hotter, less dense wind \citep{kwok93}."
+. This hot wind pushes the older wind out farther and creates a sharp density gradient and possible clumping near the interface between the cool and hot winds 1992)., This hot wind pushes the older wind out farther and creates a sharp density gradient and possible clumping near the interface between the cool and hot winds \citep{young92}.
+. This overall structure is similar to that which we infer from our modeling of SN 2002ic., This overall structure is similar to that which we infer from our modeling of SN 2002ic.
+" Assuming a SN ejecta speed of v,=30,000 km s! and a progenitor star hot wind speed of v,=100 km s! (Youngetal.1992;Herpin2002),, we conclude that the hot wind must have begun just ~15 years prior to the SN explosion."," Assuming a SN ejecta speed of $v_s=30,000$ km $^{-1}$ and a progenitor star hot wind speed of $v_w=100$ km $^{-1}$ \citep{young92,herpin02}, we conclude that the hot wind must have begun just $\sim15$ years prior to the SN explosion."
+" Alternatively, there is also the possibility that the conversion from kinetic energy to optical luminosity is for some reason significantly less efficient at very early times."," Alternatively, there is also the possibility that the conversion from kinetic energy to optical luminosity is for some reason significantly less efficient at very early times."
+ It is interesting to note that the observed light curve decline rate of SN 2002ic after 40 days past maximum light is apparently constant during these observations., It is interesting to note that the observed light curve decline rate of SN 2002ic after $40$ days past maximum light is apparently constant during these observations.
+ Spectroscopic study 2004) shows the highest observed velocity of the ejecta to be around 11000 km s! at day 200 after maximum light., Spectroscopic study \citep{wang04} shows the highest observed velocity of the ejecta to be around $11000$ km $^{-1}$ at day $200$ after maximum light.
+" If we assume a constant expansion rate, these observations of continuing emission through ~320 days after maximum provide a lower limit of ~3x1016 cm for the spatial extent of the CSM."," If we assume a constant expansion rate, these observations of continuing emission through $\sim320$ days after maximum provide a lower limit of $\sim3\times10^{16}$ cm for the spatial extent of the CSM."
+" Compared to a nominal pre-explosion stellar wind speed of 10 km s, the ejecta is moving ~1000 times more rapidly and thus has overtaken the progenitor wind from the past ~800 years."," Compared to a nominal pre-explosion stellar wind speed of $10$ km $^{-1}$, the ejecta is moving $\sim1000$ times more rapidly and thus has overtaken the progenitor wind from the past $\sim800$ years."
+ The overall smoothness of the late-time light curve shows the radial density profile of the CSM to be similarly smooth and thus implies a fairly uniform mass-loss rate between 100—800 years prior to the SN explosion., The overall smoothness of the late-time light curve shows the radial density profile of the CSM to be similarly smooth and thus implies a fairly uniform mass-loss rate between $100$ $800$ years prior to the SN explosion.
+ We take the lack of enhanced flux at early times and the bump after maximum light as evidence for a gap between the SN progenitor and the dense CSM as well as a significant further change in the mass-loss of the progenitor system ~100 years prior to the SN explosion., We take the lack of enhanced flux at early times and the bump after maximum light as evidence for a gap between the SN progenitor and the dense CSM as well as a significant further change in the mass-loss of the progenitor system $\sim100$ years prior to the SN explosion.
+" These new results prompt a reexamination of supernovae previously classified as Type In, specifically SN 1988Z (Pollasetal.1988;Stathakis&Sadler1991),, SN 1997cy 1997;Turattoetal.2000;Germany 2000), and SN 1999E (Cappellaroetal.1999;Silotietal.2000;Rigon 2003).."," These new results prompt a reexamination of supernovae previously classified as Type IIn, specifically SN 1988Z \citep{iauc4691,stathakis91}, SN 1997cy \citep{iauc6706,turatto00,germany00}, and SN 1999E \citep{iauc7091,siloti00,rigon03}."
+ These supernovae bear striking similarities in their light curves and their late-time spectra to SN 2002ic., These supernovae bear striking similarities in their light curves and their late-time spectra to SN 2002ic.
+" However, SN 2002ic is the only one of these supernovae to have been observed early in its evolution."," However, SN 2002ic is the only one of these supernovae to have been observed early in its evolution."
+" If SN 2002ic had been observed at the later times typical of the observations of these Type IIn SNe, it would not have been identified as a Type Ia. It is"," If SN 2002ic had been observed at the later times typical of the observations of these Type IIn SNe, it would not have been identified as a Type Ia. It is"
+by just over half a solar mass.,by just over half a solar mass.
+ In fact. there is not one single maximum-mass limit for a secondary to undergo WR evolution at a particular metallicity — whether it does or not depends on the amount of helium accreted. the amount of helium synthesised in the secondarys core before the onset of mass transfer and the time of mass transfer (e.g. Braun Langer 1995) both of which are determined by the initial binary parameters.," In fact, there is not one single maximum-mass limit for a secondary to undergo WR evolution at a particular metallicity – whether it does or not depends on the amount of helium accreted, the amount of helium synthesised in the secondary's core before the onset of mass transfer and the time of mass transfer (e.g. Braun Langer 1995) both of which are determined by the initial binary parameters."
+ Thus there is some variation in the threshold mass — however. in practice. this variation is only around a solar mass for fully-conservative systems.," Thus there is some variation in the threshold mass – however, in practice, this variation is only around a solar mass for fully-conservative systems."
+ Whilst the formulae used for the partial-accretion model above work well in some cases. they do not work so well for all models.," Whilst the formulae used for the partial-accretion model above work well in some cases, they do not work so well for all models."
+" In particular, models for which the metallicity and time offset behaviour are most suecessfully-matehed are generally those with periods roughly in the middle of the case A range and are at the bottom or middle of the mass range."," In particular, models for which the metallicity and time offset behaviour are most successfully-matched are generally those with periods roughly in the middle of the case A range and are at the bottom or middle of the mass range."
+ Differences are partly systematic. with the predicted time offset being too large for long periods by up to 30 and occasionally too small for short periods.," Differences are partly systematic, with the predicted time offset being too large for long periods by up to 30 and occasionally too small for short periods."
+ Similarly. the metallicity of the best-fit model and the helium metallicity sometimes differ by a small amount. although this does not appear to be systematic.," Similarly, the metallicity of the best-fit model and the helium metallicity sometimes differ by a small amount, although this does not appear to be systematic."
+ This is likely to be partially an effect of our fitting against a library of discrete models. particularly with regard to the metallicity (as the single star model grid is more closely-spaced in mass that in metallicity. for which we have only ten values available).," This is likely to be partially an effect of our fitting against a library of discrete models, particularly with regard to the metallicity (as the single star model grid is more closely-spaced in mass that in metallicity, for which we have only ten values available)."
+ Of course. it should be noted once more that fitting to single stars can only ever be an approximation and that. at the higher end of the mass range the acceretor may have synthesised significant amounts of helium itself before accretion occurs. so if it is thoroughly mixed a higher effective metallicity is expected.," Of course, it should be noted once more that fitting to single stars can only ever be an approximation and that, at the higher end of the mass range the accretor may have synthesised significant amounts of helium itself before accretion occurs, so if it is thoroughly mixed a higher effective metallicity is expected."
+ Also. as noted previously. there remains a dichotomy between fitting the HR diagram position with time well and fitting the mass evolution with time well — both are possible. but frequently not with the same model (Fig.," Also, as noted previously, there remains a dichotomy between fitting the HR diagram position with time well and fitting the mass evolution with time well – both are possible, but frequently not with the same model (Fig."
+ 3)., 3).
+ With all this noted. the results of fitting using equations | and 2 with the metallicity-dependent main-sequence lifetime equation from Hurley et al. (," With all this noted, the results of fitting using equations 1 and 2 with the metallicity-dependent main-sequence lifetime equation from Hurley et al. ("
+2000) are still significantly better than assuming standard rejuvenation at the same metallicity.,2000) are still significantly better than assuming standard rejuvenation at the same metallicity.
+ The effects of conservative accretion which we have looked at in this paper can manifest observationally in a number of ways., The effects of conservative accretion which we have looked at in this paper can manifest observationally in a number of ways.
+ For instance. if at least some mass transfer is helium-enriched. we would expect a corresponding increase in the type Ibe/type II supernova rate ratio. the Crunaway) WR star population and the CNO production from massive stars.," For instance, if at least some mass transfer is helium-enriched, we would expect a corresponding increase in the type Ibc/type II supernova rate ratio, the (runaway) WR star population and the CNO production from massive stars."
+ What is not clear is whether such effects can be distinguishable against the uncertainties arising from the question of how much matter can be accreted in any given binary., What is not clear is whether such effects can be distinguishable against the uncertainties arising from the question of how much matter can be accreted in any given binary.
+ If one assumes that primaries of binaries have a Salpeter IMF. mass transfer is conservative. the initial mass ratio distribution of binary systems is and the initial period distribution is flat in log(P). the inclusion of enhanced-metallicity rejuvenation in initially solar-metallicity accretion stars increases the population of secondaries which go through a WR phase by about 30 (excluding those systems which survive common-envelope evolution).," If one assumes that primaries of binaries have a Salpeter IMF, mass transfer is conservative, the initial mass ratio distribution of binary systems is and the initial period distribution is flat in $\log(P)$, the inclusion of enhanced-metallicity rejuvenation in initially solar-metallicity accretion stars increases the population of secondaries which go through a WR phase by about 30 (excluding those systems which survive common-envelope evolution)."
+ At the metallicity of the Small Magellanic Cloud the population is nearly doubled., At the metallicity of the Small Magellanic Cloud the population is nearly doubled.
+ This suggests that this is potentially an observable effect. if it can be distinguished from the other uncertainties which beset binary evolution.," This suggests that this is potentially an observable effect, if it can be distinguished from the other uncertainties which beset binary evolution."
+ Since WR stars produce much larger amounts of carbon and other elements in their winds than stars which do not go through a WR phase. this population increase may have implications for the enrichment from WR stars (e.g. Dray Tout 2003). even though stars in binaries which avoic a contact phase and go on to become WR-like are a relatively small part of the total population.," Since WR stars produce much larger amounts of carbon and other elements in their winds than stars which do not go through a WR phase, this population increase may have implications for the enrichment from WR stars (e.g. Dray Tout 2003), even though stars in binaries which avoid a contact phase and go on to become WR-like are a relatively small part of the total population."
+ It is also worth considering wha happens to binaries which interact and do not manage to avoic contact., It is also worth considering what happens to binaries which interact and do not manage to avoid contact.
+ If the outcome is a period of common-envelope evolution during which the secondary accretes no further matter. followed by further evolution as a detached but close binary. then we would expect the evolution of the secondary to continue as discussed in section 3 and enhanced-metallicity evolution would have a negligible effect in most cases.," If the outcome is a period of common-envelope evolution during which the secondary accretes no further matter, followed by further evolution as a detached but close binary, then we would expect the evolution of the secondary to continue as discussed in section 3 and enhanced-metallicity evolution would have a negligible effect in most cases."
+ If instead the components of the binary merge as a consequence of this process — a fate which may wppen to a majority of interacting binaries — then it is likely that he structure of the resulting star will have been quite thoroughly mixed and it would behave similarly to the conservatively-accreting secondaries previously discussed., If instead the components of the binary merge as a consequence of this process – a fate which may happen to a majority of interacting binaries – then it is likely that the structure of the resulting star will have been quite thoroughly mixed and it would behave similarly to the conservatively-accreting secondaries previously discussed.
+ If this happens. the population of binary-formed WR stars may be much larger than estimated above.," If this happens, the population of binary-formed WR stars may be much larger than estimated above."
+ It is also interesting to note that a large proportion of these ormed WR stars would be single during their WR phase., It is also interesting to note that a large proportion of these binary-formed WR stars would be single during their WR phase.
+ However. whether or not it can fit the whole-population woperties. a successful model of massive-binary evolution must also be able to explain individual systems.," However, whether or not it can fit the whole-population properties, a successful model of massive-binary evolution must also be able to explain individual systems."
+ In particular. the ooperties of individual systems may be used to constrain. the model set used because. with the large number of free or poorly-constrained parameters available to vary in population synthesis. it is not hard to make most of the possible populations fit in some way or other.," In particular, the properties of individual systems may be used to constrain the model set used because, with the large number of free or poorly-constrained parameters available to vary in population synthesis, it is not hard to make most of the possible populations fit in some way or other."
+ The rejuvenatory effects discussed in this paper only come into effect if there is either conservative of nearly-conservative mass transfer if there is a late mass transfer event (not necessarily wholly conservative)or in which matter from a star with an already strongly helium-enriched surface composition is transferred across., The rejuvenatory effects discussed in this paper only come into effect if there is either conservative of nearly-conservative mass transfer or if there is a late mass transfer event (not necessarily wholly conservative) in which matter from a star with an already strongly helium-enriched surface composition is transferred across.
+Oue&OL.,$\alpha\ind{ov}\simeq0.17$.
+ This sudden increase of the core Lifetime for ov>OL is in fact caused x the start of the pplI aud pplIH reaction chains. aud later by the CNO chain. because of the increase in temperature in the center as the star evolves.," This sudden increase of the core lifetime for $\aov>0.1$ is in fact caused by the start of the ppII and ppIII reaction chains, and later by the CNO chain, because of the increase in temperature in the center as the star evolves."
+ If the ITe has kept the convective core going until the pplI chain begins to compete. the couvection preveuts the pplI reactions from achieving equilibriuu. in exactly the same wav as described above for the ppl chain.," If the $^3$ He has kept the convective core going until the ppII chain begins to compete, the convection prevents the ppII reactions from achieving equilibrium, in exactly the same way as described above for the ppI chain."
+ The star starts burning ‘Li through the reaction ‘Li(p.a) Πο out of equilibrium. with a temperature scusitivity of πα which sustains the core.," The star starts burning $^7$ Li through the reaction $^7$ $\alpha$ $^4$ He out of equilibrium, with a temperature sensitivity of $\nu\ind{Li7}=10.8$, which sustains the core."
+ This is what happoeus for modelB., This is what happens for model.
+. If. ou the coutrary. the convective core las already disappeared at that time. there is no more mixine in the center and the pplI reactions achieve equilibria without trigecring convection.," If, on the contrary, the convective core has already disappeared at that time, there is no more mixing in the center and the ppII reactions achieve equilibrium without triggering convection."
+ This is the case for modclΑ., This is the case for model.
+. We can see in Fie., We can see in Fig.
+ 5. that the pplI chain makes a uajor contibution to the nuclear production rate z()., \ref{fig_contr_eps} that the ppII chain makes a major contribution to the nuclear production rate $\varepsilon(m)$.
+ The importance of the CNO reactions is still small., The importance of the CNO reactions is still small.
+" As nentioncd above. the cucrev brought by the pplI chain is alinost cutirely due to the burning of ""Li outside equilibriun. with the other reactions contributing very ittle to μμ."," As mentioned above, the energy brought by the ppII chain is almost entirely due to the burning of $^7$ Li outside equilibrium, with the other reactions contributing very little to $\varepsilon\ind{ppII}$."
+ Tf the star were only slightly wore evolved. he CNO evele would take over. causing the convective core to grow.," If the star were only slightly more evolved, the CNO cycle would take over, causing the convective core to grow."
+ We notice that. even though the structure of the core is different for our models with aud without overshooting. their age is very comparable (sce Table 1)).," We notice that, even though the structure of the core is different for our models with and without overshooting, their age is very comparable (see Table \ref{fit}) )."
+ This suggests that the survival of the convective core does nof sieuificautlv influence the elobal nuclear cucrey produced by dauriug its evolution., This suggests that the survival of the convective core does not significantly influence the global nuclear energy produced by during its evolution.
+ This constitutes a striking difference with higher mass models. where overshooting is kuownu to have a strongC» nupact on the evolution of the star. especially ou its age. for a given Tig aud £.," This constitutes a striking difference with higher mass models, where overshooting is known to have a strong impact on the evolution of the star, especially on its age, for a given $T\ind{eff}$ and $L$."
+ For higlhi-uass stars. since the temperature dependence of the douinaut CNO evele is huge. it operates in a narrow area di the ceuter. and the wider convective core can act as a reservoir.," For high-mass stars, since the temperature dependence of the dominant CNO cycle is large, it operates in a narrow area in the center, and the wider convective core can act as a reservoir."
+ For203608.. the pp chain is douinaut. aud its temperature dependence is uel lower.," For, the pp chain is dominant, and its temperature dependence is much lower."
+ Therefore. the reactions take place in an area wider than the extent of the couvective core. and the mixing has less effect on the evolution.," Therefore, the reactions take place in an area wider than the extent of the convective core, and the mixing has less effect on the evolution."
+" This menus that it would be hopeless to try to characterize an extension of the couvective core iu this type of star by classical fundamental stellar parameters alone (Tig. L. M). as was done by Claret (2007) for higher Lasses,"," This means that it would be hopeless to try to characterize an extension of the convective core in this type of star by classical fundamental stellar parameters alone $T\ind{eff}$, $L$, $M$ ), as was done by Claret (2007) for higher masses."
+ We present here a modeling of bbased ou the analysis of ddata performed in MOS., We present here a modeling of based on the analysis of data performed in M08.
+ Our iain result is that. on this basis. we find. strong evidence that this old low-1uass star has a convective core.," Our main result is that, on this basis, we find strong evidence that this old low-mass star has a convective core."
+ Models with convective cores enabled us to solve the disagreement with observations that was pointed out in MOS for models without convective cores. bringing the 47 function from 9.1 to 0.8.," Models with convective cores enabled us to solve the disagreement with observations that was pointed out in M08 for models without convective cores, bringing the $\chi^2$ function from $9.1$ to $0.8$."
+ All the observed xuanueters for aare now fitted within 1-0 error bars., All the observed parameters for are now fitted within $\sigma$ error bars.
+ Iu the case of our modeling of203608.. the value obtained or the ayy parameter (0.17c 0.03) is stronely constrained.," In the case of our modeling of, the value obtained for the $\aov$ parameter $0.17 \pm 0.03$ ) is strongly constrained."
+ Overshooting was here used as a proxy to model the couples processes of transport at the οσο of a convective core. ax is usually done iu the preseut state of stellar inodeling.," Overshooting was here used as a proxy to model the complex processes of transport at the edge of a convective core, as is usually done in the present state of stellar modeling."
+ Rather than flucing a unique absolute value for doy. the current ain is to try to observationallv determine which values of the oy paraimcter are needed to represent stars of different masses aud evolution stages.," Rather than finding a unique absolute value for $\aov$, the current aim is to try to observationally determine which values of the $\aov$ parameter are needed to represent stars of different masses and evolution stages."
+ Iu this respect. the value obtained for cconstitutes a valuable iuput for low-1ass objects.," In this respect, the value obtained for constitutes a valuable input for low-mass objects."
+ We discussed im detail how the existence ofa convective core m such an evolved loxv-1uass star can be explained by a reasonable amount of extra mixing (modeled here as core overshooting) inducing the survival of the carly convective core., We discussed in detail how the existence of a convective core in such an evolved low-mass star can be explained by a reasonable amount of extra mixing (modeled here as core overshooting) inducing the survival of the early convective core.
+ For lowanass stars stich as203608.. an early convective. core exists. because of. the burning. of 212 C aud Πο outside equilibrimm.," For low-mass stars such as, an early convective core exists because of the burning of $^{12}$ C and $^3$ He outside equilibrium."
+ An extra mixing at the edee of the core increases its lifetime. by brineiie more ο to the center. as mentioned in Boxbureh(1985)..," An extra mixing at the edge of the core increases its lifetime, by bringing more $^3$ He to the center, as mentioned in \cite{1985SoPh..100...21R}."
+" Tere. we showed that. above a certain amount of overshooting (Aoy~ 0.15). the burning of ""Ie out of equilibrimm sustains the core until the pplII aud. pplII reactions take over."," Here, we showed that, above a certain amount of overshooting $\aov\sim0.15$ ), the burning of $^3$ He out of equilibrium sustains the core until the ppII and ppIII reactions take over."
+ Convection prevents these reactious frou achieving equilibrium. and the burning of ‘Li outside equilibrium," Convection prevents these reactions from achieving equilibrium, and the burning of $^7$ Li outside equilibrium"
+2007)).,).
+ Such a core is separated fom the hydrogen envelope bv a thin C-O shell and an even thinner Ie shell., Such a core is separated from the hydrogen envelope by a thin C-O shell and an even thinner He shell.
+ However. the density. [alls off so steeply in the region immediately above the core that we shall refer to the C-O and Ile shells as the surface lavers of the core.," However, the density falls off so steeply in the region immediately above the core that we shall refer to the C-O and He shells as the surface layers of the core."
+ These lavers result. from He and 11 burning coupled with convection during the pre-SN evolution. and (his pre-supernova burning can give rise {ο a metallicitv-independent neutron excess through reaction sequences such as ορ.)PN followed by 3 decay of P'N to PC (Nomoto 2007. personal communication).," These layers result from He and H burning coupled with convection during the pre-SN evolution, and this pre-supernova burning can give rise to a metallicity-independent neutron excess through reaction sequences such as $^{12}{\rm C}(p,\gamma)^{13}$ N followed by $\beta^+$ decay of $^{13}$ N to $^{13}$ C (Nomoto 2007, personal communication)."
+ The SN shock rapidly accelerates as il propagates through the surface lavers of the core., The SN shock rapidly accelerates as it propagates through the surface layers of the core.
+ This gives rise {ο [ast expansion of the shocked ejecta on timescales of ~10| s. Together with an entropy of οον100 (in units of Boltzmann's constant & per nucleon) and an initial electron lraction of Y;~0.495 (e.g.. for a composition of PC:PC:ο[:3:3 by mass). this fast expansion enables an r-process to occur in the shocked ejecta. producing neutron-rich nuclei with 44>130 through the actinides.," This gives rise to fast expansion of the shocked ejecta on timescales of $\sim 10^{-4}$ s. Together with an entropy of $S\sim 100$ (in units of Boltzmann's constant $k$ per nucleon) and an initial electron fraction of $Y_e\sim 0.495$ (e.g., for a composition of ${^{13}{\rm
+C}}:{^{12}{\rm C}}:{^{16}{\rm O}}\sim 1:3:3$ by mass), this fast expansion enables an $r$ -process to occur in the shocked ejecta, producing neutron-rich nuclei with $A>130$ through the actinides."
+ We discuss the conditions during (he expansion of the shocked ejecta in relsec:conditions.., We discuss the conditions during the expansion of the shocked ejecta in \\ref{sec:conditions}.
+ Nucleosvntliesis caleulations for (vo different initial compositions are presented in relsec:nucleosvuthesis with and without neutrino reactions on free nucleons aud o-particles., Nucleosynthesis calculations for two different initial compositions are presented in \\ref{sec:nucleosynthesis} with and without neutrino reactions on free nucleons and $\alpha$ -particles.
+ We discuss (he implications of this new r-process model for abundances in metal-poor stars and lor general Galactic ehemical evolution in re[sec:gce.., We discuss the implications of this new $r$ -process model for abundances in metal-poor stars and for general Galactic chemical evolution in \\ref{sec:gce}.
+ In the generic scenario of stellar core collapse. the inner core becomes a proto-neutron star (DNS) and bounces on reaching supra-nuclear density.," In the generic scenario of stellar core collapse, the inner core becomes a proto-neutron star (PNS) and bounces on reaching supra-nuclear density."
+ This launches a shock. which propagates into the still collapsing outer core and falters.," This launches a shock, which propagates into the still collapsing outer core and falters."
+ According to the neutrino-driven SN mechanism (Bethe&Wilson1985).. the shock is re-energized through heating of the material behind it by the neutrinos emitted from the PNS.," According to the neutrino-driven SN mechanism \citep{1985ApJ...295...14B}, the shock is re-energized through heating of the material behind it by the neutrinos emitted from the PNS."
+ Using Nomotos (1934. 1987) model for a 1.38Al. O-Ne-Mg core. Mavle&Wilson(1988).. ancl most recently. Nitaura(2006) (see also Jankaetal. 2007)) indeed obtained neutrino-driven explosion al(hough wilh very different final explosion energies.," Using Nomoto's (1984, 1987) model for a $1.38\,M_\odot$ O-Ne-Mg core, \citet{1988ApJ...334..909M}, and most recently, \citet{2006A&A...450..345K} (see also \citealp{2007arXiv0706.3056J}) ) indeed obtained neutrino-driven explosion although with very different final explosion energies."
+ For our purpose. the exact explosion energy does nol matter so long as the re-energized shock propagates through the surface lavers of the core with a sullicient speed.," For our purpose, the exact explosion energy does not matter so long as the re-energized shock propagates through the surface layers of the core with a sufficient speed."
+ As only stars in a very narrow range of ~8 11M. develop Q-Ne-Me cores. we consider Nomoto's (1984. 1987) 1.38.M.. core model. which was evolved from the Ile core of à zz9.V. progenitor. to be representative ancl adopt ils quantitative," As only stars in a very narrow range of $\sim 8$ $11\,M_\odot$ develop O-Ne-Mg cores, we consider Nomoto's (1984, 1987) $1.38\,M_\odot$ core model, which was evolved from the He core of a $\approx 9\,M_\odot$ progenitor, to be representative and adopt its quantitative"
+an acceptable ft (sce also T97). but such low abuudanuces are unexpected for dominant eroupg,"an acceptable fit (see also T97), but such low abundances are unexpected for dominant group."
+"alaxies! Iu a second step. two-component models cousisting of contributions frou, both. a rs plasma and a pl source. Or two rs sources were applied."," In a second step, two-component models consisting of contributions from both, a rs plasma and a pl source, or two rs sources were applied."
+ Dui he rs|pl description. the pl iudex was fixed to 1.9 (the value vpically observed im Sy: since we want to specifically tes for the presence of an ACN).," In the rs+pl description, the pl index was fixed to –1.9 (the value typically observed in Sy; since we want to specifically test for the presence of an AGN)."
+" This provides a eood fit (A?,4 1.1).", This provides a good fit $\chi{^{2}}_{\rm red}$ = 1.1).
+ Te Nyy is treated as additional free parameter. it is found to be of the order of the Galactic value.," If $N_{\rm H}$ is treated as additional free parameter, it is found to be of the order of the Galactic value."
+ For the rs coniponent. we obtain &Tz 0.6 keV (error contours are displaved iu Fie. 7)).," For the rs component, we obtain $kT \simeq$ 0.6 keV (error contours are displayed in Fig. \ref{chi2}) )."
+" The pl conmiponeut contributes with a (0.1:2.1 keV) huuinositv of £444=3.510H. cress, the total (0.12.1 keV) οςαν is Ly=L7104! ere/s (for Aq, fixed o Nea)."," The pl component contributes with a (0.1–2.4 keV) luminosity of $L_{\rm x,pl} = 3.5\,10^{41}$ erg/s, the total (0.1–2.4 keV) luminosity is $L_{\rm x} = 4.7\,10^{41}$ erg/s (for $N_{\rm H}$ fixed to $N_{\rm gal}$ )."
+ Alternatively. the spectrum cau be fit with a double rs model.," Alternatively, the spectrum can be fit with a double rs model."
+ In this case we find the hotter component to be ill-coustrained., In this case we find the hotter component to be ill-constrained.
+" If we fix AT» = 1.5 keV. we get hE, = 0.10.2 keV. OPνα 1.1: we used δι= UN) aud Duamnosities of Da,=Lilpit cre/s. Lay,=2.910"" eres. For comparison. if kT> = 5 keV is chosen. AT, = O.540.2 keV ενα = 1.0) and Ly,=2.0103. ore/s. Ly.T,=2.7104 ere/s. The results of the spectral fits are sinumniarized in Table 1.."," If we fix $kT_2$ = 1.5 keV, we get $kT_1$ = $\pm{0.2}$ keV $\chi{^{2}}_{\rm red}$ = 1.1; we used $N_{\rm H} = N_{\rm gal}$ ) and luminosities of $L_{\rm x,T_1} = 1.7\,10^{41}$ erg/s, $L_{\rm x,T_2} = 2.9\,10^{41}$ erg/s. For comparison, if $kT_2$ = 5 keV is chosen, $kT_1$ = $\pm{0.2}$ keV $\chi{^{2}}_{\rm red}$ = 1.0) and $L_{\rm x,T_1} = 2.0\,10^{41}$ erg/s, $L_{\rm x,T_2} = 2.7\,10^{41}$ erg/s. The results of the spectral fits are summarized in Table \ref{fitres}."
+ The Nav enuüsson from the direction of 3383 is strouglv peaked., The X-ray emission from the direction of 383 is strongly peaked.
+ To check how iuuch of the cussion might arise from a point source we compared the radial source profile with the instrumental poiut-spread fiction (PSF)., To check how much of the emission might arise from a point source we compared the radial source profile with the instrumental point-spread function (PSF).
+ We fiud that the emission from 3383 is not sienificautly extended bevoud the PSF of the PSPC., We find that the emission from 383 is not significantly extended beyond the PSF of the PSPC.
+ Thus. the data are consistent with the bulk of the X-ray cussion arising fom a point source.," Thus, the data are consistent with the bulk of the X-ray emission arising from a point source."
+ IxEforming a simular analysis for the IRI observation. we flu a deviation of the source profile from the ITRI PSE.," Performing a similar analysis for the HRI observation, we find a deviation of the source profile from the HRI PSF."
+ However. similar deviations are found for the (prestmably ontlike) F star which is located within the field of view (after aking iuto account the appropriate off-axis PSF for he star).," However, similar deviations are found for the (presumably pointlike) F star which is located within the field of view (after taking into account the appropriate off-axis PSF for the star)."
+ We conclude that in the present data there is 10 evidence for source exteut., We conclude that in the present data there is no evidence for source extent.
+ The N-rax lighteurve of 3383 is displaved in Fig. 8.., The X-ray lightcurve of 383 is displayed in Fig. \ref{light}.
+ Au AGN/point-source iüghnt reveal itself dy variability (but not necessarily)., An AGN/point-source might reveal itself by variability (but not necessarily).
+ We find a constant source flux within the errors., We find a constant source flux within the errors.
+ A spectral analysis was performed for the two X-ray brightest ealaxies. 3379 and 3380.," A spectral analysis was performed for the two X-ray brightest galaxies, 379 and 380."
+ A pl model does not provide au acceptable fit οντω = 3.1 aud 5.6) with residuals strongly indicative of the presence of anu rs component., A pl model does not provide an acceptable fit $\chi{^{2}}_{\rm red}$ = 3.4 and 5.6) with residuals strongly indicative of the presence of an rs component.
+" Such a model indeed gives au excellent fit (47,44 = 1.0 and 0.8).", Such a model indeed gives an excellent fit $\chi{^{2}}_{\rm red}$ = 1.0 and 0.8).
+ We fined temperatures of 0.5 keV. 3379). aud 0.9 keV 3380)., We find temperatures of 0.5 keV 379) and 0.9 keV 380).
+" The absorption-corrected fluxes for this model description iu the (0.12.1) keV baud are f,=8510+! erg/eau? /s 3379) and fe=Laslo29 en? /s 3380), and the corresponding huuinosiies L4=1.1101! ere/s 3379) and Ly=2.310"" cress 3380)."," The absorption-corrected fluxes for this model description in the (0.1–2.4) keV band are $f_{\rm x} = 8.5\,10^{-14}$ $^2$ /s 379) and $f_{\rm x} = 1.85\,10^{-13}$ $^2$ /s 380), and the corresponding luminosities $L_{\rm x} = 1.1\,10^{41}$ erg/s 379) and $L_{\rm x} = 2.3\,10^{41}$ erg/s 380)."
+ To estimate Iunuinosities also for 338 Land 385. which are too weak to allow direct spectral fits. we adopted a rs spectrum of 0.5 keV. The derived hnuuinosities are even in Table 2.," To estimate luminosities also for 384 and 385, which are too weak to allow direct spectral fits, we adopted a rs spectrum of 0.5 keV. The derived luminosities are given in Table 2."
+ To compare the derived X-rav luninositics with blue huninosities we used the observed blue maguituces of de Vaucouleur et al. (, To compare the derived X-ray luminosities with blue luminosities we used the observed blue magnitudes of de Vaucouleur et al. (
+1991 via NED: see also Suüth ο al.,1991 via NED; see also Smith et al.
+ 1997)., 1997).
+ For the extinction correction we converted the Galactic Nyy as given in Dickey Lockman (1990) into Ap assuminga standard sas/dust ratio. the relation of Bohlin et al. (," For the extinction correction we converted the Galactic $N_{\rm H}$ as given in Dickey Lockman (1990) into $A_{\rm B}$ assuminga standard gas/dust ratio, the relation of Bohlin et al. ("
+1978: sce also Predehl Schiuitt 1995). iux the extinction curve as eiven in Osterbrock et al. (,"1978; see also Predehl Schmitt 1995), and the extinction curve as given in Osterbrock et al. ("
+1989. his Tab.,"1989, his Tab."
+ 7.2)., 7.2).
+" This vields Ap = 0.38"".", This yields $A_{\rm B}$ = $^{\rm m}$ .
+ Lp was then caleulatec, $L_{\rm B}$ was then calculated
+assumptions on the distributions of masses. intervals aud Lorentz [actors and that the index and (he normalization of the variability - Iuminositw relation highlv depends ou a and Vp.,"assumptions on the distributions of masses, intervals and Lorentz factors and that the index and the normalization of the variability - luminosity relation highly depends on $\alpha$ and $\Gamma_0$."
+ Our numerical results suggest a~1 and Dy—107., Our numerical results suggest $\alpha \sim 1$ and $\Gamma_0 \sim 10^2$.
+"5 Another assumption concerns the values of e, and eg. fractions of internal energy distributed to electrons and magnetic field at collisions."," Another assumption concerns the values of $\epsilon_e$ and $\epsilon_B$, fractions of internal energy distributed to electrons and magnetic field at collisions."
+" We have done another set of simulations with e,=e,—0.3 and get verv similar results.", We have done another set of simulations with $\epsilon_e=\epsilon_B=0.3$ and get very similar results.
+ We have shown (hat there existsa correlation between (he jel opening angle 9 and the eamnma-ray light curve variability V., We have shown that there existsa correlation between the jet opening angle $\theta$ and the gamma-ray light curve variability $V$.
+ Though the correlation is based on only seven events al present and needs to be further confirmed with more events. it is naturally expected if the luminosity L is correlated with the variability (Fenimore Ramirez-Ruiz 2000: Reichart οἱ al.," Though the correlation is based on only seven events at present and needs to be further confirmed with more events, it is naturally expected if the luminosity $L$ is correlated with the variability (Fenimore Ramirez-Ruiz 2000; Reichart et al."
+ 2001). and if GRBs have a standard energy output (Frail et al.," 2001), and if GRBs have a standard energy output (Frail et al."
+ 2001: Panaitescu Ixumar 2002: Piran et al., 2001; Panaitescu Kumar 2002; Piran et al.
+ 2001)., 2001).
+ This correlation might give us a wav to measure (the opening angle for a long burst directly from the GRB light curve., This correlation might give us a way to measure the opening angle for a long burst directly from the GRB light curve.
+ We have shown (hat the opening angle - variability relation. or equivalently. due to the constancy of burst energy. (he Iuminosity - variability relation can be interpreted as the correlation between the opening anele of a fireball jet aud the Lorentz [actor.," We have shown that the opening angle - variability relation, or equivalently, due to the constancy of burst energy, the luminosity - variability relation can be interpreted as the correlation between the opening angle of a fireball jet and the Lorentz factor."
+ Larger opening angles are expected to be associated wilh ereater mass loading al the central engine and might result in lower Lorentz [actors., Larger opening angles are expected to be associated with greater mass loading at the central engine and might result in lower Lorentz factors.
+ We also show that such a correlation ean be a natural consequence of the collapsar model., We also show that such a correlation can be a natural consequence of the collapsar model.
+ Using a multiple-shell model. we numerically caleulate the temporal profiles and estimate the Iuminosity and the variability.," Using a multiple-shell model, we numerically calculate the temporal profiles and estimate the luminosity and the variability."
+ Our numerical results suggest Dx0. For equivalently M.x8., Our numerical results suggest $\Gamma \propto \theta^{-1}$ or equivalently $M \propto \theta$.
+" If the opening augle of a jet is verv wide 665. the shells are almost ordered. by increasing values of the Lorentz [actors in the photosphere. only minor collisions or. at the eXtreme, only external shocks happen."," If the opening angle of a jet is very wide $\theta \gg \theta_2$, the shells are almost ordered by increasing values of the Lorentz factors in the photosphere, only minor collisions or, at the extreme, only external shocks happen."
+ The resulting dim burst should be smooth and has a soft spectrum., The resulting dim burst should be smooth and has a soft spectrum.
+ GRBOS0425 and recently reported Fast X-ray Transients (Ileise et al., GRB980425 and recently reported Fast X-ray Transients (Heise et al.
+" 2001: ]xippen et al 2001) could be classified into (his ""Wide Jet Case”.", 2001; Kippen et al 2001) could be classified into this “Wide Jet Case”.
+ Norris. Marani Bonnell (2000) found that the huninositw £ is inversely proportional to the GRD pulse lag time.," Norris, Marani Bonnell (2000) found that the luminosity $L$ is inversely proportional to the GRB pulse lag time."
+ The detailed study on the lag Gime is bevond the scope of this paper., The detailed study on the lag time is beyond the scope of this paper.
+ However. if the time lag is also determined by the angular spreading time scale. which is the key time scale in our model. we can get the same relation.," However, if the time lag is also determined by the angular spreading time scale, which is the key time scale in our model, we can get the same relation."
+" In our model.luminosity is scaled by the opening. angle Lx@E . while. the angular spreading. time. is. /,,;xP27 67."," In our model,luminosity is scaled by the opening angle $L \propto \theta^{-2}$ , while the angular spreading time is $t_{ang}\propto \Gamma^{-2} \propto \theta^2$ ."
+stars studied by Authouv-Twarog et al. (,stars studied by Anthony-Twarog et al. (
+1988).,1988).
+ In Fie., In Fig.
+ 1 je two possible counterparts are represented by triangle sviubols., 1 the two possible counterparts are represented by triangle symbols.
+ Finally. the position of source #111 coincides with a conglomerate of at least six stars.," Finally, the position of source 11 coincides with a conglomerate of at least six stars."
+ Iu Fig., In Fig.
+ we plot) je three brightest of these. which we regard as tlhe most ikely candidates candidates. as crosses.," 1 we plot the three brightest of these, which we regard as the most likely candidates candidates, as crosses."
+" All the sources narked in Table d with ""OUT"" as optical ID are outside 1e regions studied by the above authors and probably are iot Cluster incumbers.", All the sources marked in Table 1 with `OUT' as optical ID are outside the regions studied by the above authors and probably are not cluster members.
+ We four of ideutifiedthom with stars yon the Ciuide Star Catalog. one with a SAO star. while o the other six we teutatively assigneda iuagnuitude using he Digitized Sky Survey (Postiman ct al..," We identified four of them with stars from the Guide Star Catalog, one with a SAO star, while to the other six we tentatively assigned a magnitude using the Digitized Sky Survey (Postman et al.,"
+ iu preparation)., in preparation).
+ The remaining eight A-vayv sources do not have visible optical candidates in their error boxes aud are not listed in Table 1: two of thei (113.21) le iu the ceutral part of the cluster.," The remaining eight X-ray sources do not have visible optical candidates in their error boxes and are not listed in Table 1; two of them 13,21) lie in the central part of the cluster."
+" 11651 is located at low ealactic latitude (oy,~ STO)", 4651 is located at low galactic latitude $_{II} \sim -8^\circ$ ).
+ The total galactic interstellar absorption along the Lue of sight. estimated from radio data. is Nyy~L.7«10270 7 (Dickey Lockman 1990).," The total galactic interstellar absorption along the line of sight, estimated from radio data, is $_H \sim 1.7\times 10^{21}$ $^{-2}$ (Dickey Lockman 1990)."
+ To estimate the nuniber of feld sources expected. we assumed a typical power-law spectrum with photon iudex 2.0 and the above value of Ny.," To estimate the number of field sources expected, we assumed a typical power-law spectrum with photon index 2.0 and the above value of $_H$."
+ With these parameters our liuütiug count rate of —1 ct/Esec translates iuto a flux limit of 2.0410 τοσο 7s + (in the 0.12.5 keV baud)., With these parameters our limiting count rate of $\sim$ 1 ct/ksec translates into a flux limit of $\times 10^{-14}$ erg $^{-2}$ $^{-1}$ (in the 0.4–2.5 keV band).
+ From the log Nlog S distribution derived. from the ROSAT Lockman hole deep survey (Iasiuser ct al., From the log N–log S distribution derived from the ROSAT Lockman hole deep survey (Hasinger et al.
+ 1997). we estimate that roughly 15 detections of extragalactic sources are expected iun our observation.," 1997), we estimate that roughly 15 detections of extragalactic sources are expected in our observation."
+ All the sources without a POSS optical counterpart are therefore most Likely extragalactic. as well as some of the weak identifications outside the region of the cluster.," All the sources without a POSS optical counterpart are therefore most likely extragalactic, as well as some of the weak identifications outside the region of the cluster."
+and C are dimensionless parameters.,and $C$ are dimensionless parameters.
+" For several low-charge states. the RR rate coefficients exhibit large ""bumps"" at high temperature."," For several low-charge states, the RR rate coefficients exhibit large “bumps” at high temperature."
+ These features are caused by electron capture into low principal quantum number states. in which the nucleus ts not as effectively screened as for higher states.," These features are caused by electron capture into low principal quantum number states, in which the nucleus is not as effectively screened as for higher states."
+ Equations | and 2. are not able to reproduce the size of the bumps. and become inaccurate at high temperatures.," Equations \ref{rr1} and \ref{rr2} are not able to reproduce the size of the bumps, and become inaccurate at high temperatures."
+" For these Se7-Se*- target states. we therefore fit the high-temperature RR rate coefficients with the function where 7 and £; are in temperature units (ΚΚ). the rate coefficient is in em’ ss7!. c; in em? ss7! KK??, and 5 ranges from 5 to 7 depending on the ion."," For these $^+$ $^{3+}$ target states, we therefore fit the high-temperature RR rate coefficients with the function where $T$ and $E_i$ are in temperature units (K), the rate coefficient is in $^3$ $^{-1}$, $c_i$ in $^3$ $^{-1}$ $^{3/2}$, and $n$ ranges from 5 to 7 depending on the ion."
+ This analytical function is typically used to fit DR rate coefficients (e.g...2).. as done in Sect. ?2..," This analytical function is typically used to fit DR rate coefficients \citep[e.g., ][]{zatsarinny03}, as done in Sect. \ref{dr}."
+ The fit coefficients were determined using a non-linear least-squares fit algorithm. generally leading to fit accuracies within for the low-temperature fits. and <2% at high temperatures.," The fit coefficients were determined using a non-linear least-squares fit algorithm, generally leading to fit accuracies within for the low-temperature fits, and $\leq$ at high temperatures."
+ However. the low-temperature fits are only accurate to within for Se (level 1) and Se (level 1). and the accuracies of the high-temperature fits for Se are (levels 2-4) to (level 5).," However, the low-temperature fits are only accurate to within for $^+$ (level 1) and $^{3+}$ (level 1), and the accuracies of the high-temperature fits for $^+$ are (levels 2-4) to (level 5)."
+ Moreover. Equations | and 2. are unable to reproduce the rate coefficient for RR onto level 5 of Se at the lowest temperatures (<20 K). where they deviate from the calculated rate coefficients by up to20%.," Moreover, Equations \ref{rr1} and \ref{rr2} are unable to reproduce the rate coefficient for RR onto level 5 of $^+$ at the lowest temperatures $\leq 20$ K), where they deviate from the calculated rate coefficients by up to."
+. Aside from this exception. the fits have the correct asymptotic forms outside of the temperature range (10!— 10) K (specifically. Equations | and 2. deseribe the asymptotic behavior as T-0 K. and Equation 3. correctly describes the rate coefficient at the high-temperature limit).," Aside from this exception, the fits have the correct asymptotic forms outside of the temperature range $^1-10^7$ $z^2$ K (specifically, Equations \ref{rr1} and \ref{rr2} describe the asymptotic behavior as $T\rightarrow 0$ K, and Equation \ref{dreq} correctly describes the rate coefficient at the high-temperature limit)."
+" Table 1. displays the fit coefficients from equations | and 2.. as well as the maximum temperature 7,4, for which the fits are validto within the stated accuracies."," Table \ref{rrlowtfits} displays the fit coefficients from equations \ref{rr1} and \ref{rr2}, as well as the maximum temperature $T_{\rm max}$ for which the fits are validto within the stated accuracies."
+ If Τι Is not given in Table {.. then equations | and 2 are accurate over the full temperature range (10!— 10) K and correctly describe the asymptotic behavior of the rate coefficients at the high-temperature limit.," If $T_{\rm max}$ is not given in Table \ref{rrlowtfits}, then equations \ref{rr1} and \ref{rr2} are accurate over the full temperature range $^1-10^7$ $z^2$ K and correctly describe the asymptotic behavior of the rate coefficients at the high-temperature limit."
+ Table 2. provides fit coefficients for the RR rate coefficients from Tyas to 10/2 K. In refrrfig.. we display the ground state RR rate coefficients for low-charge Se tons as a function of temperature.," Table \ref{rrhitfits} provides fit coefficients for the RR rate coefficients from $T_{\rm max}$ to $10^7 z^2$ K. In \\ref{rrfig}, we display the ground state RR rate coefficients for low-charge Se ions as a function of temperature."
+ We tested the sensitivity of the calculated RR rate coefficients to the use of different CI expansions and other internal parameters in our caleulations., We tested the sensitivity of the calculated RR rate coefficients to the use of different CI expansions and other internal parameters in our calculations.
+ The uncertainties that we discuss were estimated at temperatures near 107 K. similar to electron temperatures in photoionized nebulae.," The uncertainties that we discuss were estimated at temperatures near $^4$ K, similar to electron temperatures in photoionized nebulae."
+ As is the case for PI. the RR rate coefficients are most sensitive to the CI expansion used.," As is the case for PI, the RR rate coefficients are most sensitive to the CI expansion used."
+ The rate coefhcients calculated with the small and large configuration sets (Table 4)) differ from the presented rate coefficients (from the medium CI expansion) typically by less than near 10 Κ. The sole exception is that of Se. for which the deviation from the presented rate coefficients ts25-30%.," The rate coefficients calculated with the small and large configuration sets (Table \ref{ciexp}) ) differ from the presented rate coefficients (from the medium CI expansion) typically by less than near $^4$ K. The sole exception is that of $^+$, for which the deviation from the presented rate coefficients is."
+ Other tested parameters produced smaller changes in the RR rate coefficients., Other tested parameters produced smaller changes in the RR rate coefficients.
+ For example. to test the sensitivity of forcing velocity gauge at high energies. we computed the rate coefficients from PI cross sections cut off at the maximum velocity gauge energy (beyond which velocity gauge must be forced in AUTOSTRUCTURE).," For example, to test the sensitivity of forcing velocity gauge at high energies, we computed the rate coefficients from PI cross sections cut off at the maximum velocity gauge energy (beyond which velocity gauge must be forced in AUTOSTRUCTURE)."
+ The RR rate coefficients were unaffected except at the highest temperatures (typically a few times 10°-10° K)., The RR rate coefficients were unaffected except at the highest temperatures (typically a few times $^5$ $^6$ K).
+ This indicates that uncertainties associated to our use of velocity gauge at high energies are negligible for the computed RR rate coefficients at photoionized plasma temperatures. and become important only above 10° K for singly- and doubly-charged Se. and above 10° K for higher charge states.," This indicates that uncertainties associated to our use of velocity gauge at high energies are negligible for the computed RR rate coefficients at photoionized plasma temperatures, and become important only above $^5$ K for singly- and doubly-charged Se, and above $^6$ K for higher charge states."
+ The calculations were performed using 3 interpolatior energies per decade., The calculations were performed using 3 interpolation energies per decade.
+ We tested the effects of using a finer interpolation mesh of five energies per decade. and found differences of less than in the RR rate coefficients for all of the investigated Se tons.," We tested the effects of using a finer interpolation mesh of five energies per decade, and found differences of less than in the RR rate coefficients for all of the investigated Se ions."
+ Similarly. the use of target ion scaling parameters versus those of the N+ I-electron ion produced discrepancies in the rate coefficients larger than only i the case of Se and Se. where the differences are (depending on the state) near 10. K. Finally. we computed RR rate coefficients without Schmidt orthogonalizing the radial orbitals. and found that the rate coefficient differed by a maximum of (for Se and Se) compared to our orthogonalized calculations.," Similarly, the use of target ion scaling parameters versus those of the $N+1$ -electron ion produced discrepancies in the rate coefficients larger than only in the case of $^+$ and $^{2+}$, where the differences are (depending on the state) near $^4$ K. Finally, we computed RR rate coefficients without Schmidt orthogonalizing the radial orbitals, and found that the rate coefficient differed by a maximum of (for $^+$ and $^{2+}$ ) compared to our orthogonalized calculations."
+ We conclude that our computed RR rate coetficients have internal uncertainties of less than with the exception of Se. where the uncertainty 1s 30—406c..," We conclude that our computed RR rate coefficients have internal uncertainties of less than with the exception of $^+$ , where the uncertainty is ."
+across the sky. waxing and waning on various timescales.,"across the sky, waxing and waning on various timescales."
+ In fact. at some level. everything in the sky is variable.," In fact, at some level, everything in the sky is variable."
+ The introduction of CCD detectors to astronomy greatly enhanced the ability to conduct variability surveys., The introduction of CCD detectors to astronomy greatly enhanced the ability to conduct variability surveys.
+ The extension of CCD cameras to mosaic and wide-field formats along with the exponential progression of computing power have allowed the subsequent development of more ambitious surveys reaching to deeper magnitudes. higher cadences and larger sky areas.," The extension of CCD cameras to mosaic and wide-field formats along with the exponential progression of computing power have allowed the subsequent development of more ambitious surveys reaching to deeper magnitudes, higher cadences and larger sky areas."
+ For more details we direct the reader to Beckeretal.(2004). who present a clear summary of modern variability surveys., For more details we direct the reader to \citet{bec2004} who present a clear summary of modern variability surveys.
+" In this work we concentrate on optical photometric and astrometric variability thence “light and motion"") over a ~249 deg patch of sky.", In this work we concentrate on optical photometric and astrometric variability (hence “light and motion”) over a $\sim$ 249 $^{2}$ patch of sky.
+ Large sky surveys such aus. the Sloan Digital Sky Survey (SDSS: Yorketal. 2000)) have in many ways revolutionised our knowledge of the Universe., Large sky surveys such as the Sloan Digital Sky Survey (SDSS; \citealt{yor2000}) ) have in many ways revolutionised our knowledge of the Universe.
+ SDSS has imaged approximately a quarter of the sky in five photometric wave bands., SDSS has imaged approximately a quarter of the sky in five photometric wave bands.
+ The exploitation of this impressive dataset has resulted in hundreds of publications covering a wide range of astronomical topics. from the structure of the Milky Way to the mapping of a large fraction of the Universe.," The exploitation of this impressive dataset has resulted in hundreds of publications covering a wide range of astronomical topics, from the structure of the Milky Way to the mapping of a large fraction of the Universe."
+ The bulk of this data. however. contain only single measurements of objects from the north Galactic cap with no information on possible photometric variability or astrometric motion.," The bulk of this data, however, contain only single measurements of objects from the north Galactic cap with no information on possible photometric variability or astrometric motion."
+ Substantial efforts have been made by Munnetal.(2004) (see also Gould&Kollmeier (2004))) to measure proper motions by matching SDSS data from the north Galactic cap with the USNO-B catalogue (Monetetal. 2003))., Substantial efforts have been made by \citet{mun2004} (see also \citet{gou2004}) ) to measure proper motions by matching SDSS data from the north Galactic cap with the USNO-B catalogue \citealt{mon2003}) ).
+ The resultant proper motion catalogue is complete down to g=19.7 with the magnitude limit being set by the USNO-B catalogue faint magnitude limits., The resultant proper motion catalogue is complete down to $g~=~19.7$ with the magnitude limit being set by the USNO-B catalogue faint magnitude limits.
+ One of the primary goals of the SDSS is the study of the variable sky CAdelman-MeCarthyetal. 2007)) of which our knowledge is still very incomplete (Paczyfski 2000)., One of the primary goals of the SDSS is the study of the variable sky \citealt{ade2007}) ) of which our knowledge is still very incomplete \citealt{pac2000}) ).
+ Το this end. the SDSS has repeatedly imaged a 300 square degree area. the so called Stripe 82. during the later half of each year since 1998.," To this end, the SDSS has repeatedly imaged a 300 square degree area, the so called Stripe 82, during the later half of each year since 1998."
+ In 2005. the SDSS-II Supernova Survey (Friemanetαἱ. 2008)) started with the aim of detecting Type-I supernovae in Stripe 82. greatly improving the cadence of measurements within the stripe.," In 2005, the SDSS-II Supernova Survey \citealt{fri2008}) ) started with the aim of detecting Type-I supernovae in Stripe 82, greatly improving the cadence of measurements within the stripe."
+ By averaging a subset of the repeated observations of unresolved sources in Stripe 82. Iveziéetal.(2007) built a standard star catalogue containing 1 million nonvariable sources with + band magnitudes in the range 14-22. by far the deepest and most numerous set of photometric standards available.," By averaging a subset of the repeated observations of unresolved sources in Stripe 82, \citet{ive2007} built a standard star catalogue containing $\sim$ 1 million nonvariable sources with $r$ band magnitudes in the range 14-22, by far the deepest and most numerous set of photometric standards available."
+ Using these same multi-epoch photometric data. Sesaratal.(2007). analysed. the photometric variability for ~ 1.4 million unresolved sources in the stripe. drawing interesting conclusions on the spatial distribution of RR Lyrae stars and the variability of quasars.," Using these same multi-epoch photometric data, \citet{ses2007}
+ analysed the photometric variability for $\sim$ 1.4 million unresolved sources in the stripe, drawing interesting conclusions on the spatial distribution of RR Lyrae stars and the variability of quasars."
+ Here we present a new public archive of light-motion curves in SDSS Stripe 82., Here we present a new public archive of light-motion curves in SDSS Stripe 82.
+ The archive has been constructed from the set of high-precision multi-epoch photometric and astrometric measurements made in the stripe since the first SDSS runs in 1998 until the end of 2005., The archive has been constructed from the set of high-precision multi-epoch photometric and astrometric measurements made in the stripe since the first SDSS runs in 1998 until the end of 2005.
+ In constructing the catalogue. we only use measurements of objects that are cleanly detected in individual SDSS runs.," In constructing the catalogue, we only use measurements of objects that are cleanly detected in individual SDSS runs."
+ The catalogue contains almost 4 million objects. galaxies and stars. and is complete down to magnitude 21.5 in i. g. r and 7. and to magnitude 20.5 in z.," The catalogue contains almost 4 million objects, galaxies and stars, and is complete down to magnitude 21.5 in $u$, $g$, $r$ and $i$, and to magnitude 20.5 in $z$."
+ Each object has its proper motion calculated based only on the multi-epoch SDSS J2000 astrometric measurements., Each object has its proper motion calculated based only on the multi-epoch SDSS J2000 astrometric measurements.
+ The catalogue reaches almost two magnitudes deeper than the SDSS/USNO-B catalogue. making it the deepest large-area photometric and astrometric catalogue available.," The catalogue reaches almost two magnitudes deeper than the SDSS/USNO-B catalogue, making it the deepest large-area photometric and astrometric catalogue available."
+ The catalogue comes in two flavours. the Light-Motion Curve Catalogue (LMCC). which contains the set of individual light- curves. where measured quantities for each object are," The catalogue comes in two flavours, the Light-Motion Curve Catalogue (LMCC), which contains the set of individual light-motion curves, where measured quantities for each object are"
+dB below the peak.,dB below the peak.
+ The GBT's gregorian 21 cm receiver system overilluminates the subreflector somewhat. producing a broad forward spillover lobe which contains about of the telescope’s total response.," The GBT's gregorian 21 cm receiver system overilluminates the subreflector somewhat, producing a broad forward spillover lobe which contains about of the telescope's total response."
+ However. this sidelobe has a diameter >30° on the sky and its effects on the current observations should be negligible (Lockmandon 2005).," However, this sidelobe has a diameter $> 30\degr$ on the sky and its effects on the current observations should be negligible \citep{loc05}."
+. The GBT was used to measure spectra at ten positions along the SW major axis of M31 in August and September 2005., The GBT was used to measure spectra at ten positions along the SW major axis of M31 in August and September 2005.
+ At 21 em. the GBT has an FWHM angular resolution of 9/1.," At 21 cm, the GBT has an FWHM angular resolution of $9\farcm1$."
+ Spectra were taken with 2.5 effective velocity resolution over a range >1000 kms''centered on the M31 velocity., Spectra were taken with 2.5 effective velocity resolution over a range $> 1000$ centered on the M31 velocity.
+" Frequency switching ""in. band” gave good instrumental baselines and good sensitivity.", Frequency switching `in band' gave good instrumental baselines and good sensitivity.
+ The dual-polarization receiver had a system| temperature Ta~ 18 K. The ten positions were observed for times varying from a few minutes to more than one hour. with a typical value being 40 minutes.," The dual--polarization receiver had a system temperature $_{\rm sys} \sim$ 18 K. The ten positions were observed for times varying from a few minutes to more than one hour, with a typical value being 40 minutes."
+ emission was detected out to R=45 kpe with the Effelsberg dish. but here we will discuss only data at <35 kpe where spectra were measured with both Effelsberg and the GBT.," emission was detected out to $R =45$ kpc with the Effelsberg dish, but here we will discuss only data at $\leq35$ kpc where spectra were measured with both Effelsberg and the GBT."
+" Moreover. the disk of M31 is known to be slightly warped (e.g. Newton Emerson 1977; Henderson 1979),"," Moreover, the disk of M31 is known to be slightly warped (e.g. Newton Emerson 1977; Henderson 1979)."
+ Knowing our beamsize and estimates of the warp parameters. we have calculated that profiles obtained along a constant line of aare reliable only up to 35 kpe.," Knowing our beamsize and estimates of the warp parameters, we have calculated that profiles obtained along a constant line of are reliable only up to 35 kpc."
+ Within this range the agreement between data from the two telescopes is good., Within this range the agreement between data from the two telescopes is good.
+ For the positions observed with both telescopes (27.6. 32.7. 34.7 kpe) the mean difference in radial velocity is less than 2.5.. and at these three locations the radial velocity 15 thus given by the average of the two observations.," For the positions observed with both telescopes (27.6, 32.7, 34.7 kpc), the mean difference in radial velocity is less than 2.5, and at these three locations the radial velocity is thus given by the average of the two observations."
+ The spectra. baseline subtracted. with the best fit Gaussian profile superposed. are shown in Figure 1..," The spectra, baseline subtracted, with the best fit Gaussian profile superposed, are shown in Figure \ref{fig1}."
+ While there is some scaling differences between the profiles coming from the two telescopes. the profile shapes are remarkably similar (e.g.. the two spectra at R=27.6 kpe).," While there is some scaling differences between the profiles coming from the two telescopes, the profile shapes are remarkably similar (e.g., the two spectra at $R=27.6$ kpc)."
+ There are discrepancies between the different determinations| of the RC of M31 in the inner— 10 kpe (Sofue&Kato1981).. but there is general agreement for 10xRx:30 kpe.," There are discrepancies between the different determinations of the RC of M31 in the inner 10 kpc \citep{sof81}, but there is general agreement for $10 \leq R \leq 30$ kpc."
+ One interesting feature of the RC of M31 ts that it seems to be declining more or less regularly from the center all the way out to 30 kpe (Braun1991).., One interesting feature of the RC of M31 is that it seems to be declining more or less regularly from the center all the way out to 30 kpc \citep{bra91}. .
+ While there are à few galaxies with declining RCs over a limited radius range (Carignan&Puche1990:Casertano&vanGorkom1991:HonmaSofue 1997).. this 1s quite unlike what is seen in most spirals. and is one of the reasons we wished to obtain the velocity information at larger radit to check whether that trend continues.," While there are a few galaxies with declining RCs over a limited radius range \citep{cap90,cas91,hon97}, this is quite unlike what is seen in most spirals, and is one of the reasons we wished to obtain the velocity information at larger radii to check whether that trend continues."
+ Table |. and Figure 2 give our derived RC., Table \ref{table1} and Figure \ref{fig2} give our derived RC.
+ For R<90’ (or ~ 20.5 Κρο). we have redetermined the velocity field using the early data-cube of Unwin(1983). and recomputed the rotation velocities by fitting a tilted-ring model to the velocity map.," For $R \le 90'$ (or $\sim$ 20.5 kpc), we have redetermined the velocity field using the early data-cube of \citet{unw83} and recomputed the rotation velocities by fitting a tilted-ring model to the velocity map."
+ The task (Begeman 1989) of the GIPSY package (van der Hulst et al., The task (Begeman 1989) of the GIPSY package (van der Hulst et al.
+ 1992) was used for that purpose., 1992) was used for that purpose.
+" No velocities are given inside a radius of 25’, where the deficiency of neutral gas introduces large uncertainties. and where it is probable that there are large non-circular motions."," No velocities are given inside a radius of $25\arcmin$, where the deficiency of neutral gas introduces large uncertainties, and where it is probable that there are large non-circular motions."
+ The error-bars for the inner RC are given by the difference between the velocities of the approaching and receding sides of the galaxy. except for the very few points where the formal error given by is chosen because it exceeds this difference.," The error-bars for the inner RC are given by the difference between the velocities of the approaching and receding sides of the galaxy, except for the very few points where the formal error given by is chosen because it exceeds this difference."
+ The inner part of the new RC compares well with the composite one presented in Widrow et al. (, The inner part of the new RC compares well with the composite one presented in Widrow et al. (
+2003). which was compiled from several data sets.,"2003), which was compiled from several data sets."
+ The curves of the receding and approaching halves are shown to illustrate the symmetry of the gas motions in the galaxy: except. perhaps. near R~70! where a difference of ~ 29 iis observed. similar values are derived for the two sidesof," The curves of the receding and approaching halves are shown to illustrate the symmetry of the gas motions in the galaxy: except, perhaps, near $R \sim 70'$ where a difference of $\sim$ 29 is observed, similar values are derived for the two sidesof"
+burst started at ~06:26 UT. almost at the same time (slightly earlier) as the wave started to separate from the driver.,"burst started at $\sim$ 06:26 UT, almost at the same time (slightly earlier) as the wave started to separate from the driver."
+ The radio observation suggests Chat a shock was generated al (hat lime. indicating that the compression wave had just turned itself into a shock wave.," The radio observation suggests that a shock was generated at that time, indicating that the compression wave had just turned itself into a shock wave."
+ Recent studies by Patsourakos&Vourlidas(20092). ancl Patsourakosetal.(2010) also noticed two fronts: the wave Iront and the bubble front., Recent studies by \citet{pat09a} and \citet{pat10} also noticed two fronts: the wave front and the bubble front.
+ However. due to the relative low cadence of the STEREO-EUVI observations (5 minutes). the two distinet [ronts were seen in only a few frames.," However, due to the relative low cadence of the -EUVI observations (5 minutes), the two distinct fronts were seen in only a few frames."
+ Here. the ALA observations in everv 12 seconds not only reveal the existence of two fronts but also the detailed separation process of the diffuse front from the sharp bubble front alter (he expansion of the bubble slows down.," Here, the AIA observations in every 12 seconds not only reveal the existence of two fronts but also the detailed separation process of the diffuse front from the sharp bubble front after the expansion of the bubble slows down."
+ These observations clearly demonstrate (hat the separated diffuse front is a true MIID wave. driven by the early accelerating expansion of the CAIE bubble.," These observations clearly demonstrate that the separated diffuse front is a true MHD wave, driven by the early accelerating expansion of the CME bubble."
+ In conclusion. our observations help understand (he physical nature of usually termed EUV waves.," In conclusion, our observations help understand the physical nature of usually termed EUV waves."
+" The EUV wave is actually a composite phenomenon. consisting of two distinct fronts, a wave front and a CALE (compressed plasma) front."," The EUV wave is actually a composite phenomenon, consisting of two distinct fronts, a wave front and a CME (compressed plasma) front."
+ We further find that its evolution can be divided into (wo stages., We further find that its evolution can be divided into two stages.
+ The first stage takes place in the accelerating expansion phase of the CME bubble. which acts as the piston-driver of the MIID wave.," The first stage takes place in the accelerating expansion phase of the CME bubble, which acts as the piston-driver of the MHD wave."
+ During (his stage. the wave [ront is coupled together with the compression front of the CME bubble.," During this stage, the wave front is coupled together with the compression front of the CME bubble."
+ In the second stage. when the expanding velocity of the CME bubble slows down. the MIID wave lront decouples from the compression bubble front. Forms a distnet [ront. aud propagates across (he solar disk.," In the second stage, when the expanding velocity of the CME bubble slows down, the MHD wave front decouples from the compression bubble front, forms a distinct front, and propagates across the solar disk."
+ The observational result For the coexistence of both wave and non-wave Ironts are in general consistent with the models and numerical simulations lor EUV waves by 2005).," The observational result for the coexistence of both wave and non-wave fronts are in general consistent with the models and numerical simulations for EUV waves by \citet{chen02,chen05}."
+. We believe that the previous dispute about the nature of EUV waves resides in. al least. partly. different parts of the composite phenomenon that the authors mav have observed.," We believe that the previous dispute about the nature of EUV waves resides in, at least partly, different parts of the composite phenomenon that the authors may have observed."
+ In fact. the duration of the wave and CME compression front coupling is different. for each event. depending on the dvnamucs of the CME ancl the surrounding environment.," In fact, the duration of the wave and CME compression front coupling is different for each event, depending on the dynamics of the CME and the surrounding environment."
+ We thank P. F. Chen for many valuable comments (hat helped to improve (he manuscript significantly., We thank P. F. Chen for many valuable comments that helped to improve the manuscript significantly.
+ SDO is a mission of NASA's Living With a Star Program., SDO is a mission of NASA's Living With a Star Program.
+ X.C.. ancl M.D.D. are supported by NSFC under grants. LOGT3004. 10525000. and. 10933003. and NINBRSE under grant 2011CD811402.," X.C., and M.D.D. are supported by NSFC under grants 10673004, 10828306, and 10933003 and NKBRSF under grant 2011CB811402."
+ A.C. is also supported by the scholarship granted by the China Scholarship Council (CSC) under file No., X.C. is also supported by the scholarship granted by the China Scholarship Council (CSC) under file No.
+ 2010619071., 2010619071.
+ J.Z. is supported by NSF grant. and NASA grant NNGO5GGI9G. A.V. is supported by NASA contract S-136361-Y., J.Z. is supported by NSF grant ATM-0748003 and NASA grant NNG05GG19G. A.V. is supported by NASA contract S-136361-Y.
+the known densities.,the known densities.
+" Underestinates are anticipate because, af one extreme. nass that is not stronely clustered on the scale of the survey has little dvuamica consequence and. at the other extreme. couples. shellcrossing orbits cannot be modeled."," Underestimates are anticipated because, at one extreme, mass that is not strongly clustered on the scale of the survey has little dynamical consequence and, at the other extreme, complex shell-crossing orbits cannot be modeled."
+ Analysis of reconstructions in 12 cosinological simulations with ciffereut initial fluctuation characteristics discussed by Mohavaceetal.(2005) resulted in recovery of SOY20% (lo) of the known model densitv., Analysis of reconstructions in 12 cosmological simulations with different initial fluctuation characteristics discussed by \cite{moh05} resulted in recovery of $80\% \pm 20\%$ $1 \sigma$ ) of the known model density.
+ Iu the case of Least Action. it las been appreciated that iuass can be muderestimated on σπα. scales (Branchini&Carlbere1991:Dranchiuietal. 2002).," In the case of Least Action, it has been appreciated that mass can be underestimated on small scales \citep{bra94,bra02}."
+". A quantitative nieasure of the effect is found using our Least Actiou reconstruction algorithiis with au N-body simulation with O4=0.7. O,,=0.3. evaluated by. placing the observer at multiple locations to test for cosmic variance."," A quantitative measure of the effect is found using our Least Action reconstruction algorithms with an N-body simulation with $\Omega_\Lambda=0.7$, $\Omega_m=0.3$, evaluated by placing the observer at multiple locations to test for cosmic variance."
+ Recovered deusities were 70:4420% of the model deusity (Phelps. Desjacques. and Nusser: ongoing work).," Recovered densities were $70\% \pm 20\%$ of the model density (Phelps, Desjacques, and Nusser; ongoing work)."
+" The estimates that we quote for O,, mcelude adjustineuts for the systematics of and for MLATS and Least Action. respectively."," The estimates that we quote for $\Omega_m$ include adjustments for the systematics of and for MAK and Least Action, respectively."
+ The results of the two methods are sunmuuidzed in Fie., The results of the two methods are summarized in Fig.
+" 1 where contours of the \? parameter are plotted in the domain (t9.0,,]..", 1 where contours of the $\chi^2$ parameter are plotted in the domain $t_0$ $\Omega_{\rm m}$.
+" In the case of the Least Action analysis. niass density @,, is derived from AL/£ values bv accepting the mean B band huninosity deusitv of the Universe deteriumed from Blautonctal.(2005) with reddening corrections."," In the case of the Least Action analysis, mass density $\Omega_m$ is derived from $M/L$ values by accepting the mean $B$ band luminosity density of the Universe determined from \cite{bla03} with reddening corrections."
+ We use οE)/(15 10h). a transform with ~20% uncertainty.," We use $\Omega_m = (M/L) / (1540 h)$ , a transform with $\sim 20\%$ uncertainty."
+ The error cllipses frou the and studies are elougated., The error ellipses from the and studies are elongated.
+ With a linear analysis the elliptical V? troughs would open to infinity. but non-linear effects create a specific minimum along cach 4? trough.," With a linear analysis the elliptical $\chi^2$ troughs would open to infinity, but non-linear effects create a specific minimum along each $\chi^2$ trough."
+" At a fixed age. fy. there are relatively tight constraints on O,, requirements."," At a fixed age, $t_0$, there are relatively tight constraints on $\Omega_m$ requirements."
+ If shorter ages are cutertained. then higher deusities are required to arrive at the observed dvianiical state iu the specified tine.," If shorter ages are entertained, then higher densities are required to arrive at the observed dynamical state in the specified time."
+ The coutrary dependence between deusity and time are seen in the results from the microwave backeround (Sperecl2003) and galaxy redshift (Teemarketal.2001). experiments that are also superimposed ou Fie., The contrary dependence between density and time are seen in the results from the microwave background \citep{spe03} and galaxy redshift \citep{teg04} experiments that are also superimposed on Fig.
+ d., 1.
+ Oulv a small domain around fy=13.5 Cove lies within or near the 20 contours of all the methodologies., Only a small domain around $t_0=13.5$ Gyr lies within or near the $2 \sigma$ contours of all the methodologies.
+" With the constraint 5=0.8 consistent with the zero poiut of the distance estimates. there is eood aegrecinent between theWALAP.SDSS.ALA. and iieasures of the deusity paraiucter: O,,=0.2240.02."," With the constraint $h=0.8$ consistent with the zero point of the distance estimates, there is good agreement between the, and measures of the density parameter: $\Omega_m=0.22\pm0.02$."
+" The low value of Q,, obtained here is consistent with that found from divers recent studies (Bahcalletal.2003:Ox-tiker2003:vaudenBoschctal.Coleet 2005)."," The low value of $\Omega_{\rm m}$ obtained here is consistent with that found from divers recent studies \citep{ostriker20031, ostriker20032, vdb03, col05}."
+. The velocity field analyses discussed. above make the sinplistic assuniption of coustant AL/L., The velocity field analyses discussed above make the simplistic assumption of constant $M/L$.
+ There has becu evidence in the literature. though. for huge variations in ML with cuviromment (Marinoni&IIudson:2002:vandeuBoschetal.2005:Parker 2005).," There has been evidence in the literature, though, for large variations in $M/L$ with environment \citep{mar02, vdb03, eke04,
+tul05b, vdb05, par05}."
+. Here we note that galaxy flows near the Virgo Cluster contradict the hypothesis of constant AZ/L., Here we note that galaxy flows near the Virgo Cluster contradict the hypothesis of constant $M/L$.
+ The high velocities of infall suggest a auch larger ML for the cluster than the value consistent with a good fit over the full region within 3.000 kins 7.," The high velocities of infall suggest a much larger $M/L$ for the cluster than the value consistent with a good fit over the full region within 3,000 km $^{-1}$."
+ The critical galaxies for this discussion lie outside the cluster ou the plane of the sky so only a modest component of the iufall iiotiou is projected ite the line of sight but there is no confusion with cluster membership., The critical galaxies for this discussion lie outside the cluster on the plane of the sky so only a modest component of the infall motion is projected into the line of sight but there is no confusion with cluster membership.
+ Galaxies iufalliug from the foreground. of the cluster are redshifted with respect to the cluster aud ealaxies infalling from the backeround are blueshifted., Galaxies infalling from the foreground of the cluster are redshifted with respect to the cluster and galaxies infalling from the background are blueshifted.
+ If oue could imagine many test particles distributed along a liue of sight. they would create a wave in velocity as a function of distance with the same velocity arising at three distances (Tourv&Davis1981) two located within the iufall region aud one located at the cosmic expansion position.," If one could imagine many test particles distributed along a line of sight, they would create a wave in velocity as a function of distance with the same velocity arising at three distances \citep{ton81}
+ – two located within the infall region and one located at the cosmic expansion position."
+ The amplitude of the peaks of the ‘triplewaluc’ waves depend on the mass. AM. iuterior to the position associated with the peaks. r.," The amplitude of the peaks of the `triple-value' waves depend on the mass, $M$, interior to the position associated with the peaks, $r$."
+ The euvelope ofobserved infall velocities with + provides a description (Tully&Shava1981). of the run of AL(<r), The envelope of observed infall velocities with $r$ provides a description \citep{tul84} of the run of $M(<r)$.
+ The outermost caustic of the Virgo Cluster proper (the radius of second turnaround) is at ~2 Mpe from the center of the cluster., The outermost caustic of the Virgo Cluster proper (the radius of second turnaround) is at $\sim 2$ Mpc from the center of the cluster.
+ Galaxies on first iufall acquire very hiel radial velocities at 2.1 Alpe from the Virgo core., Galaxies on first infall acquire very high radial velocities at 2–4 Mpc from the Virgo core.
+ Ax exiuuple is given in Fie., An example is given in Fig.
+ 2., 2.
+ The run of velocities with distance are what is expected along the specific line of seht of the 113 eroup iu the Nearby Galaxies Catalog with three different assuuptious reearding the mass of the Virgo Cluster., The run of velocities with distance are what is expected along the specific line of sight of the 11–3 group in the Nearby Galaxies Catalog with three different assumptions regarding the mass of the Virgo Cluster.
+ The observed. distance and velocity of the 113 group is indicated in the figure., The observed distance and velocity of the 11–3 group is indicated in the figure.
+ It is clear that the 113 eroup is at oue of the two infall locations (or thelocation where the two iufall positionsbecome degenerate at the tip of the triple-value curve) rather than at the cosmic expansion position around 30 Mpc., It is clear that the 11–3 group is at one of the two infall locations (or thelocation where the two infall positionsbecome degenerate at the tip of the triple-value curve) rather than at the cosmic expansion position around 30 Mpc.
+ The total mass within 3.000 kins + is the same in all 3," The total mass within 3,000 km $^{-1}$ is the same in all 3"
+We have performed the first survey of electric. field data in the solar wind to measure the perpendicular spectrum of inertial range fluctuations.,We have performed the first survey of electric field data in the solar wind to measure the perpendicular spectrum of inertial range fluctuations.
+ It was found that there is a spread of spectral index values but that the average spectral index depends on the frame of measurement., It was found that there is a spread of spectral index values but that the average spectral index depends on the frame of measurement.
+ In the spacecraft frame. the y-component of the electric field is primarily due to the magnetic field fluctuations being convected past the spacecraft at the average solar wind speed.," In the spacecraft frame, the $y$ -component of the electric field is primarily due to the magnetic field fluctuations being convected past the spacecraft at the average solar wind speed."
+ It. therefore. has the same average spectral index (to within errors) as the magnetic field of —1.66+0.01.," It, therefore, has the same average spectral index (to within errors) as the magnetic field of $-1.66\pm0.01$."
+ This is consistent with previous single interval electric field measurements in. the spacecraft frame (Baleetal.2005:Sahraoui2009).," This is consistent with previous single interval electric field measurements in the spacecraft frame \citep{bale05,sahraoui09}."
+. In the mean solar wind frame. the electric field is primarily due to velocity fluctuations in à mean magnetic field and has a spectral index of —1.404:0.01. which ts closer to the velocity spectral index than the magnetic field spectral index. although not the same to within errors.," In the mean solar wind frame, the electric field is primarily due to velocity fluctuations in a mean magnetic field and has a spectral index of $-1.40\pm0.01$, which is closer to the velocity spectral index than the magnetic field spectral index, although not the same to within errors."
+ The compressive fluctuations (|B| and 7) were found to have the same spectral index as the magnetic field and not the velocity., The compressive fluctuations $|\mathbf{B}|$ and $n$ ) were found to have the same spectral index as the magnetic field and not the velocity.
+ The difference between the scaling of the electric field in the spacecraft frame and the mean solar wind frame provides independent confirmation of the difference in scaling between the velocity and magnetic field., The difference between the scaling of the electric field in the spacecraft frame and the mean solar wind frame provides independent confirmation of the difference in scaling between the velocity and magnetic field.
+ This difference is not expected for ffluctuations. since àv is proportional to 6B in an wwave. and is not predicted by any of the current theories of tturbulence (although see recent work by (2011) and Wangetal.(2011))).," This difference is not expected for fluctuations, since $\delta\mathbf{v}$ is proportional to $\delta\mathbf{B}$ in an wave, and is not predicted by any of the current theories of turbulence (although see recent work by \citet{boldyrev11} and \citet{wang11}) )."
+ Recently. Roberts(2010) found that further out into the heliosphere. past 5 AU. the velocity spectral index evolves toward —5/3 to match the magnetic field.," Recently, \citet{roberts10} found that further out into the heliosphere, past 5 AU, the velocity spectral index evolves toward $-5/3$ to match the magnetic field."
+ Although an important result. this does not explain the difference at | AU.," Although an important result, this does not explain the difference at 1 AU."
+ Possible reasons for the difference include the effects of scale-dependent alignment. imbalance and residual energy and these will be investigated in à future paper.," Possible reasons for the difference include the effects of scale-dependent alignment, imbalance and residual energy and these will be investigated in a future paper."
+ This work was supported by NASA grant NNXO9AEAITG. We acknowledge the team and NASA contract AS5-02099., This work was supported by NASA grant NNX09AE41G. We acknowledge the team and NASA contract NAS5-02099.
+ We thank J. Bonnell. J. McFadden. A. Schekochihin. and D. Sundkvist for useful conversations and an anonymous referee for helpful comments.," We thank J. Bonnell, J. McFadden, A. Schekochihin, and D. Sundkvist for useful conversations and an anonymous referee for helpful comments."
+where we assume the adiabatic index ο=5/3 for non-relativistic gas. ancl for m we assume primordial abundances of HE and He.,"where we assume the adiabatic index $\gamma = 5/3$ for non-relativistic gas, and for $m$ we assume primordial abundances of H and He."
+ By equating these speeds. we have a formula for a critical mass: Ata ogiven ogas temperature.1 halos below this critical mass will have gas thermal speeds greater than their escape velocities.," By equating these speeds, we have a formula for a critical mass: At a given gas temperature, halos below this critical mass will have gas thermal speeds greater than their escape velocities."
+ Hence they will not be able to hold on to their gas. and therefore be unable to form stars.," Hence they will not be able to hold on to their gas, and therefore be unable to form stars."
+ For a temperature of 2.101 Ix. corresponding to the peak of the LE cooling curve. Moi=2510M..," For a temperature of $2\times 10^4$ K, corresponding to the peak of the H cooling curve, $M_{\text{crit}} = 2 \times 10^8 M_{\odot}$."
+ Vherefore we expect halos with less than this total mass to have very little star formation. and at [ate times. very little barvonie mass of any kind. assuming there is a mechanism to heat the gas to 2107 IX bv the epoch where star formation in these halos would. be significant.," Therefore we expect halos with less than this total mass to have very little star formation, and at late times, very little baryonic mass of any kind, assuming there is a mechanism to heat the gas to $2\times
+10^4$ K by the epoch where star formation in these halos would be significant."
+ lleating above this temperature will further increase the cutoll mass., Heating above this temperature will further increase the cutoff mass.
+ Moreover. since larger svstenis are formed. by the hierarchical assembly of smaller. ones. we expect. this barvon deficit to creep to larger-mass halos as time passes. ellectively increasing the critical mass at the present by some [actor f1.," Moreover, since larger systems are formed by the hierarchical assembly of smaller ones, we expect this baryon deficit to creep to larger-mass halos as time passes, effectively increasing the critical mass at the present by some factor $f \ga 1$."
+ ‘To verily the physical argument presented above. we analyze several high-resolution simulations that are. performed with various ionizing (UV. and. N-rav) backgrounds.," To verify the physical argument presented above, we analyze several high-resolution simulations that are performed with various ionizing (UV and X-ray) backgrounds."
+ These simulations were performed with a slightlv updated: version of the GADGET-2-basec code. used. in.E: refer to that work for additional details., These simulations were performed with a slightly updated version of the GADGET-2-based code used in; refer to that work for additional details.
+ In brief the code is based on GADGET-2. and the halos are selected. from. a. (50 Alpe)? box with cosmological (ACDAL) initial conditions. and resimulated in a (1 Alpe)? box centered on cach halo. with the central (0.5 Mpc) using high-resolution DM and eas/star particles2009).," In brief, the code is based on GADGET-2, and the halos are selected from a (50 $^3$ box with cosmological $\Lambda$ CDM) initial conditions, and resimulated in a (1 $^3$ box centered on each halo, with the central (0.5 $^3$ using high-resolution DM and gas/star particles."
+. Since we specifically. select high-density regions containing massive galaxies. our results are not necessarily applicable to other environments.," Since we specifically select high-density regions containing massive galaxies, our results are not necessarily applicable to other environments."
+ To ameliorate the cllects of cosmic variance. we choose seven galaxies (that is. seven initial-conditions (10) boxes). with final virial masses in the range 12107A4..," To ameliorate the effects of cosmic variance, we choose seven galaxies (that is, seven initial-conditions (IC) boxes), with final virial masses in the range $1-2\times 10^{12} M_{\odot}$."
+ The simulations do not include optical depth effects. in particular the self-shielding of dense star-forming regions from the UV background.," The simulations do not include optical depth effects, in particular the self-shielding of dense star-forming regions from the UV background."
+ They do (cillerenthy from 13) include. feedback. from. supernovac though not ACN. Following the two-phase model. and a simple prescription for metal-line cooling using cooling rates caleulatecd by Cloudy.1998). which presumes photoionization ancl collisional equilibrium for the gas and. metal atoms. and assumes 0.1 solar metallicity.," They do (differently from ) include feedback from supernovae though not AGN, following the two-phase model, and a simple prescription for metal-line cooling using cooling rates calculated by Cloudy, which presumes photoionization and collisional equilibrium for the gas and metal atoms, and assumes $0.1$ solar metallicity."
+ Star formation is performed at a fixed density threshold. calculated by the code such that 90% of the gas above the threshold will be in cold phase: his threshold is poi=1.6.107 e  or nguLOT em7.," Star formation is performed at a fixed density threshold calculated by the code such that $\sim 90\%$ of the gas above the threshold will be in cold phase: this threshold is $\rho_{\text{crit}}=1.6\times
+ 10^{-25}$ g $^{-3}$ or $n_{H,\text{crit}}=0.07$ $^{-3}$."
+ Star particles have half the mass of original gas xwticles: thus each gas particle can ultimately turn into two star particles., Star particles have half the mass of original gas particles; thus each gas particle can ultimately turn into two star particles.
+ As inL. all simulations were performed with initially 100 each of SPLL (i.e. barvon) particles and DAL particles. with a gravitational softening length of 0.25 kpe for the gas and star particles ancl twice that for the dark matter particles.," As in, all simulations were performed with initially $100^3$ each of SPH (i.e. baryon) particles and DM particles, with a gravitational softening length of 0.25 kpc for the gas and star particles and twice that for the dark matter particles."
+" Gas and star particles have. masses in the range 47.107AL. depending on the size of the ον: the assumed cosmoloev is (O3;04.0,/3.0sh)=(0.3.0.7.0.2.0.86.0.65)."," Gas and star particles have masses in the range $4 - 7 \times 10^5 M_{\odot}$, depending on the size of the box; the assumed cosmology is $(\Omega_M,\Omega_{\Lambda},\Omega_b/\Omega_M,\sigma_8,h)=(0.3,0.7,0.2,0.86,0.65)$ ."
+ We run these simulations using a variety of ionizing xickerounds: the same set which was used inL. with wo additions.," We run these simulations using a variety of ionizing backgrounds: the same set which was used in, with two additions."
+ The first. which we call No UV. has as the name suggests no ionizing radiation at all.," The first, which we call No UV, has as the name suggests no ionizing radiation at all."
+ Phe next three xickerounds are Old. UV. an updated version of(1996): New UV. which has the same spectrum as Old UV but an intensity (às measured by c.g. the 11 shotolonization rate) which declines as (1|2)+ above the oealkk at zo2.2: and FO UV. the background. constructed in(2009).," The next three backgrounds are Old UV, an updated version of; New UV, which has the same spectrum as Old UV but an intensity (as measured by e.g. the H photoionization rate) which declines as $(1+z)^{-1}$ above the peak at $z\simeq 2.2$; and FG UV, the background constructed in."
+. Phe final two backgrounds acid an X-ray component. which has à spectrum taken from and a normalization from(20073.," The final two backgrounds add an X-ray component, which has a spectrum taken from and a normalization from."
+.. The New UV|X model takes the New UV model and adds this Naray. component. with the same redshilt dependence. while FG: UV|X adds this component to the FG UV background but with the redshift dependence of Old UV. to more closely replicate the quasar background that the X-rav component is meant to model.," The New UV+X model takes the New UV model and adds this X-ray component with the same redshift dependence, while FG UV+X adds this component to the FG UV background but with the redshift dependence of Old UV to more closely replicate the quasar background that the X-ray component is meant to model."
+ The quasar background peaks at 2~ and becomes negligible for z32007)., The quasar background peaks at $z\sim 2$ and becomes negligible for $z\ga 3$.
+. Lavine performed. the simulations. we identify virialized systems in both the star and total (stars plus gas plus DM) xwticle populations using ALL. the Amiga Lalo Finder2009).," Having performed the simulations, we identify virialized systems in both the star and total (stars plus gas plus DM) particle populations using AHF, the Amiga Halo Finder."
+. ALIE uses a density grid- method followed by the removal of unbound. particles o find halos and subhalos. and returns the virial mass of each.," AHF uses a density grid-tree method followed by the removal of unbound particles to find halos and subhalos, and returns the virial mass of each."
+ The virial mass is defined as the total mass of all (relevant) particles inside the virial radius. which in urn is defined as the radius where the calculated density xolfile exceeds a threshold. multiple of the critical density determined by cosmology and redshift (~ 200): for subhalos he virial raclius is instead the largest radius of any particle which is gravitationally bound. to the halo.," The virial mass is defined as the total mass of all (relevant) particles inside the virial radius, which in turn is defined as the radius where the calculated density profile exceeds a threshold multiple of the critical density determined by cosmology and redshift $\sim 200$ ); for subhalos the virial radius is instead the largest radius of any particle which is gravitationally bound to the halo."
+ See the ALIE documentation For details., See the AHF documentation for details.
+ Figure 1. shows the mass distribution of stellar systems (identified using star particles only)., Figure \ref{fig:fracts} shows the mass distribution of stellar systems (identified using star particles only).
+ ALIE can. identify systems as small as 1.5«10A4. or about 40 star particles. but we adopt 10747. as a lower mass limit to ensure completeness.," AHF can identify systems as small as $1.5\times 10^7 M_{\odot}$, or about 40 star particles, but we adopt $10^8 M_{\odot}$ as a lower mass limit to ensure completeness."
+ Phe mean number of svstems (averaged over the seven IC's) in the lowest mass decade (107—10AZ.) declines from 131 with No UV to —43 with the three UV. models and ~20 for the two UV plus X-ray cases. differences," The mean number of systems (averaged over the seven ICs) in the lowest mass decade $(10^8 -
+10^9 M_{\odot})$ declines from $131$ with No UV to $\sim 43$ with the three UV models and $\sim 20$ for the two UV plus X-ray cases, differences"
+suggests that wwas weakly detected during both 2011. observations (see Sect. 22)).,suggests that was weakly detected during both 2011 observations (see Sect. \ref{subsec:2011}) ).
+ We extracted count rates and spectra using the tools DATENPRACT and SPECENTRACT. respectively.," We extracted count rates and spectra using the tools DMEXTRACT and SPECEXTRACT, respectively."
+ We combined the 2011 spectral data using the CLAO--task COMBINE SPECTRA to improve the statistics., We combined the 2011 spectral data using the -task $\_$ SPECTRA to improve the statistics.
+ Phe 2003. oobservation of Terzan 5 was used by. Wijnandsetal.(2005) to study the quicseent properties of248., The 2003 observation of Terzan 5 was used by \citet{wijnands2005} to study the quiescent properties of.
+".. A that time the source was detected at a mean count rate of (0.71+0.05)102countss."" anc displaving. significan.. variability in the light curve (0.510 keV).", At that time the source was detected at a mean count rate of $(0.71\pm0.05)\times10^{-2}~\cnts$ and displaying significant variability in the light curve (0.5–10 keV).
+ The quiescen spectrum was found to be hard. ancl absent of a therma emission component., The quiescent spectrum was found to be hard and absent of a thermal emission component.
+ The data could be best described by a single absorbed powerlaw model with a hydrogen column densitv Nig(1.4x0.5)1077em and a spectra index of P=LScx0.5., The data could be best described by a single absorbed powerlaw model with a hydrogen column density $N_H=(1.4\pm0.5)\times10^{22}~\nh$ and a spectral index of $\Gamma=1.8\pm0.5$.
+ Wijnandsetal.(2005). conclude that any possible thermal emission component contributec at most LO% to the unabsorbed 0.510 keV. [lux of Py0m23Qolo0ergem7s|, \citet{wijnands2005} concluded that any possible thermal emission component contributed at most $\sim10\%$ to the unabsorbed 0.5–10 keV flux of $F_X=2.2^{+0.7}_{-0.3}\times10^{-13}~\flux$.
+ An upper limit on the neutron star temperature as seen by an observer at infinity of ATNO eV. was inferred., An upper limit on the neutron star temperature as seen by an observer at infinity of $kT^{\infty}\lesssim80$ eV was inferred.
+ We reprocessed ancl refitte the 2003 data in this work and obtained similar results as found. previously (see Table 3))., We reprocessed and refitted the 2003 data in this work and obtained similar results as found previously (see Table \ref{tab:spec}) ).
+ During the 2009 observation wwas detected at an average count rate of (1.00+0.06)10countss (0.5LO keV)., During the 2009 observation was detected at an average count rate of $(1.09\pm0.06)\times10^{-2}~\cnts$ (0.5–10 keV).
+ The count rate lighteurve. shown in Fig. 3..," The count rate lightcurve, shown in Fig. \ref{fig:lc},"
+ reveals considerable variation by a factor of ~3 on a time scale of hours., reveals considerable variation by a factor of $\sim3$ on a time scale of hours.
+ We fitted the 2009 spectral cata using a variety of models., We fitted the 2009 spectral data using a variety of models.
+ Similar to the results for the 2003 observation. we find that a thermal emission moclel does not provide a good description of the data.," Similar to the results for the 2003 observation, we find that a thermal emission model does not provide a good description of the data."
+ We first tried a neutron star atmosphere mocel. which is often used to deseribe the N-rav. spectra of quiescent. neutron. star LMXDs.," We first tried a neutron star atmosphere model, which is often used to describe the X-ray spectra of quiescent neutron star LMXBs."
+ Using the model NSATAIOS (Lleinkeetal.2006) we explored fits by either fixing several parameters or leaving these [ree to float., Using the model NSATMOS \citep[][]{heinke2006} we explored fits by either fixing several parameters or leaving these free to float.
+ TE trials resulted in a reduced chi-squared value A26. henee no aeceptable fit could be obtained.," All trials resulted in a reduced chi-squared value $\chi_{\nu}^2>6$, hence no acceptable fit could be obtained."
+ A simple blackbody can describe the data. but returns unrealistic model parameters: Ny«0.31077em7. a temperature AY=L4d0.2 keV and an emitting radius R=0.04ed km (for D=5.5 kpe). vielding M—l.l for 21 degrees of [reedom (ο).," A simple blackbody can describe the data, but returns unrealistic model parameters: $N_H<0.3\times10^{22}~\nh$, a temperature $kT=1.4\pm0.2$ keV and an emitting radius $R=0.04^{+0.6}_{-0.04}$ km (for $D=5.5$ kpc), yielding $\chi_{\nu}^2=1.1$ for 21 degrees of freedom (dof)."
+ Phe obtained hydrogen column density is well below that inferred. for the 2000 outburst ancl 2003 quiescent data of citepNy~1.5.1077 wijnands2005.aswellastheavera: ragsourcesinerzanb(Ng ~1910-7y Furthermore. the temperature and emitting radius would be highly unusual for a (non-pulsating) thermallv-emitting neutron star in quiescence (e.g.Itutledgectal.1999)...," The obtained hydrogen column density is well below that inferred for the 2000 outburst and 2003 quiescent data of \\citep[$N_H \sim 1.5 \times10^{22}~ , as well as the average value found for the X-ray sources in Terzan 5 \citep[$N_H \sim 1.9 \times10^{22}~. Furthermore, the temperature and emitting radius would be highly unusual for a (non-pulsating) thermally-emitting neutron star in quiescence \citep[e.g.,][]{rutledge1999}."
+ We thus conclude that this model does not correctly describe the data., We thus conclude that this model does not correctly describe the data.
+ Fitting the 2009 spectral data with a simple absorbed powerlaw vields Ny=(1.1+0.3)J1077em? and E.—1.5250.3 (AL=1.3 for 22 dof)., Fitting the 2009 spectral data with a simple absorbed powerlaw yields $N_H=(1.1\pm0.3)\times10^{22}~\nh$ and $\Gamma=1.3\pm0.3$ $\chi_{\nu}^2=1.3$ for 22 dof).
+ The resulting 0.510 keV unabsorbed [ux is £y=(2.7+0.2).10Pergem7s1 (sec Table 3))., The resulting 0.5–10 keV unabsorbed flux is $F_X=(2.7\pm0.2) \times10^{-13}~\flux$ (see Table \ref{tab:spec}) ).
+ To constrain any thermal emission. component. we add an NS.VEMOS model with a canonical neutron star mass and radius of Αν=1.4M. and Rays=10 km. a source distance of 2=5.5 kpe and a normalisation. of unity.," To constrain any thermal emission component, we add an NSATMOS model with a canonical neutron star mass and radius of $M_{NS}=1.4~\Msun$ and $R_{NS}=10$ km, a source distance of $D=5.5$ kpc and a normalisation of unity."
+ The only free fit parameter for this model component is the neutron star temperature., The only free fit parameter for this model component is the neutron star temperature.
+ Addition of this mocel component does not improve the lit (A5.=193 for 21 dof)., Addition of this model component does not improve the fit $\chi_{\nu}^2=1.3$ for 21 dof).
+ Since the spectral data do not require the inclusion ol a thermal component. we consider the obtained neutron star temperature and thermal (lux as upper limits to their true values.," Since the spectral data do not require the inclusion of a thermal component, we consider the obtained neutron star temperature and thermal flux as upper limits to their true values."
+ As such we find that the thermal component contributes 42% to the total 0.510 keV unabsorbed Lux and we obtain a neutron star temperature of &TS5 eV. ὃν extrapolating the NS.NEMOS model fit to an energy range of 0.01.100. keV. we estimate a thermal bolometric luminosity of L1qol1075ergs.1.," As such we find that the thermal component contributes $\lesssim42\%$ to the total 0.5–10 keV unabsorbed flux and we obtain a neutron star temperature of $kT^{\infty}\lesssim85$ eV. By extrapolating the NSATMOS model fit to an energy range of 0.01–100 keV, we estimate a thermal bolometric luminosity of $L^{\mathrm{th}}_{\mathrm{q,bol}}\lesssim 1 \times10^{33}~\lum$."
+" ""his is comparable to the results for the 2003 observation (see Table 3)).", This is comparable to the results for the 2003 observation (see Table \ref{tab:spec}) ).
+ The 0.510 keV. unabsorbect [lux inferred. from the 2009 spectral data is higher than observed. in. 2003., The 0.5–10 keV unabsorbed flux inferred from the 2009 spectral data is higher than observed in 2003.
+ This sugeests that the source intensity may be variable on a timescale of vears. although the two measurements are consistent within the 90% confidence errors (see Table 3)).," This suggests that the source intensity may be variable on a timescale of years, although the two measurements are consistent within the $90\%$ confidence errors (see Table \ref{tab:spec}) )."
+ ‘Lo further investigate whether there are intensity or spectral variations between 2003. anc 2009. we fitted the two data sets simultaneously to an absorbed powerlaw model.," To further investigate whether there are intensity or spectral variations between 2003 and 2009, we fitted the two data sets simultaneously to an absorbed powerlaw model."
+ When all spectral parameters are forced. to be the same (i.e. assuming there are no spectral dillerences between the| two data sets)ὃν we obtain Na=(1320.3)«1077em7 and |—L6:0.3 for x5=1.3 (38 dof).," When all spectral parameters are forced to be the same (i.e., assuming there are no spectral differences between the two data sets), we obtain $N_H=(1.3\pm0.3)\times10^{22}~\nh$ and $\Gamma=1.6\pm0.3$ for $\chi_{\nu}^2=1.3$ (38 dof)."
+ Enabling the powerlaw normalisation to vary between the two epochs (i.e. allowing only the intensity to change) provides a better fit that vields Ng=(12x03).107cm and P—15£03for VS=1.0 (36 dof).," Enabling the powerlaw normalisation to vary between the two epochs (i.e., allowing only the intensity to change) provides a better fit that yields $N_H=(1.2\pm0.3)\times10^{22}~\nh$ and $\Gamma=1.5\pm0.3$for $\chi_{\nu}^2=1.0$ (36 dof)."
+ This mocel fit is shown in Fig. 4.., This model fit is shown in Fig. \ref{fig:spec}.
+ An [-tests suggests that there is a TX probability of achieving this level of improvement by chance., An f-tests suggests that there is a $1\%$ probability of achieving this level of improvement by chance.
+ The inferred. 0.5.10 keV unabsorbed Uuxes are £x=(2.020:2)1027 and (2.7+0.2).10Pergem7s+ for the 2003 and 2009 data. respectively.," The inferred 0.5–10 keV unabsorbed fluxes are $F_X=(2.0\pm0.2) \times10^{-13}$ and $(2.7\pm0.2) \times10^{-13}~\flux$ for the 2003 and 2009 data, respectively."
+ This suggests that the source intensity changed by ~30% between the two epochs., This suggests that the source intensity changed by $\sim30\%$ between the two epochs.
+ We next allowed. the spectral shape to vary between 2003 ancl 2009., We next allowed the spectral shape to vary between 2003 and 2009.
+ Phe 2000 outburst lighteurve of deisplaved X-ray dips (Markwardt:et.al. 2000)..., The 2000 outburst lightcurve of displayed X-ray dips \citep[][]{markwardt2000_2}. .
+ Such behavior has been seen in some LAINBs and is thought to result from obscuration of the central X-ray source. e.g.. by," Such behavior has been seen in some LMXBs and is thought to result from obscuration of the central X-ray source, e.g., by"
+ (Klypin ," \citep{Klypin99, Moore99b}."
+10° (Madauetal.2008:Springel ," $10^9$ \citep{Madau08a,
+ Springel08, Stadel08}."
+My=—18.5 Vinay~60 (vandenBergh2000:," $M_V =
+-18.5$ $\vmax \sim 60$ \citep{vandenBergh00, VanDerMarel02, Stanimirovic04}."
+ satellites compared to the MCs in the MW (see.e.g..Fig-ure5inMadauetal. 2008)..," satellites compared to the MCs in the MW \citep[see, e.g.,
+Figure 5 in][]{Madau08a}."
+ Indeed. most of these high resolution halos contain no objects of similar mass to either of the two MCs.," Indeed, most of these high resolution halos contain no objects of similar mass to either of the two MCs."
+" While this should hardly be taken as a serious problem since the handful of high resolution simulated halos does not constitute a statistical sample. it is important to understand how typical or atypical the Milky Way ts. and whether or not not there is an “extra satellites"" problem at high satellite masses within the modern Lambda Cold Dark Matter (ACDM) paradigm."," While this should hardly be taken as a serious problem since the handful of high resolution simulated halos does not constitute a statistical sample, it is important to understand how typical or atypical the Milky Way is, and whether or not not there is an “extra satellites” problem at high satellite masses within the modern Lambda Cold Dark Matter $\Lambda$ CDM) paradigm."
+ Because the MCs are relatively nearby and have been studied with high-precision instruments such as HST for some time. detailed 3-dimensional measurements of their velocities have been made (Kallivayaliletal.2006b.a).," Because the MCs are relatively nearby and have been studied with high-precision instruments such as HST for some time, detailed 3-dimensional measurements of their velocities have been made \citep{Kallivayalil06a, Kallivayalil06b}."
+. Some of these studies have indicated that the Magellanic Clouds may be on their first orbit around the Milky Way (Beslaetal.2007:Bushaetal. 20100: this may make such an extra satellite problem less worrisome tf the presence of the Magellanic Clouds is a transient event.," Some of these studies have indicated that the Magellanic Clouds may be on their first orbit around the Milky Way \citep{Besla07, Busha10c}; this may make such an extra satellite problem less worrisome if the presence of the Magellanic Clouds is a transient event."
+ Regardless of the dynamical state of the Clouds. it is an important test of galaxy formation to understand the likelihood for MW-like systems to host massive satellite galaxies.," Regardless of the dynamical state of the Clouds, it is an important test of galaxy formation to understand the likelihood for MW-like systems to host massive satellite galaxies."
+ Recent developments have afforded us the possibility to address this problem in more detail with both observations and theoretical models., Recent developments have afforded us the possibility to address this problem in more detail with both observations and theoretical models.
+ From the observational side. wide area surveys such as the Sloan Digital Sky Survey (SDSSYorketal.2000:Abazajian2009) have given us the ability to probe galaxy content for a large number of Milky Way-magnitude objects (Chenetal.2006:James&Ivory2010;Liuetal. 2010).," From the observational side, wide area surveys such as the Sloan Digital Sky Survey \citep[SDSS][]{York00,Abazajian09} have given us the ability to probe galaxy content for a large number of Milky Way-magnitude objects \citep{Chen06, James10, Liu10}."
+. The main theoretical effort has been in populating dark matter halos with galaxies using semi-analytic methods and hydrodynamic simulations., The main theoretical effort has been in populating dark matter halos with galaxies using semi-analytic methods and hydrodynamic simulations.
+ Koposovetal.(2009) used a number of toy models to add galaxy properties to dark matter halos generated using à combination of Press-Schechter theory with semi-analytic models for tracking subhalo orbits (Zentneretal.2005).., \citet{Koposov09} used a number of toy models to add galaxy properties to dark matter halos generated using a combination of Press-Schechter theory with semi-analytic models for tracking subhalo orbits \citep{Zentner05}.
+ They found that it was very difficult to model objects as bright as Magellanic clouds without allowing foran extremely high star formation efficiency., They found that it was very difficult to model objects as bright as Magellanic clouds without allowing foran extremely high star formation efficiency.
+ This, This
+the model PSF for each image using typically from 60 to 120 stars.,the model PSF for each image using typically from 60 to 120 stars.
+ The absolute calibration of the observations to the V-Johnson and I-Cousins systems 1s based on a set of standard stars from the catalog of Landolt (1992))., The absolute calibration of the observations to the V-Johnson and I-Cousins systems is based on a set of standard stars from the catalog of Landolt \cite{landolt92}) ).
+" Specifically. the observed standard stars were in the fields: PGO231. SA95 (41. 3. 96. 97, 98. 100. 101. 102. 112. 113). SA98 (556. 557. 563. 580. 581. LI. 614. 618. 626. 627. 634. 642). RUBIN 149. RUBIN 152. PG0918. PG0942. PGIO47. PGI323. PG1525. PGI530. PG1633. and Mark A. At least 3 exposures were taken for each standard field. with a total of ~100 standard star measurements per night and per filter."," Specifically, the observed standard stars were in the fields: PG0231, SA95 (41, 43, 96, 97, 98, 100, 101, 102, 112, 115), SA98 (556, 557, 563, 580, 581, L1, 614, 618, 626, 627, 634, 642), RUBIN 149, RUBIN 152, PG0918, PG0942, PG1047, PG1323, PG1525, PG1530, PG1633, and Mark A. At least 3 exposures were taken for each standard field, with a total of $\sim100$ standard star measurements per night and per filter."
+ The reduction απά aperture photometry of standard star fields were performed in the same way as for the cluster images., The reduction and aperture photometry of standard star fields were performed in the same way as for the cluster images.
+" The aperture magnitudes were corrected for atmospheric extinction. assuming ντ=0.11 and 4,=0.08 as extinction coefficients for the V. and / filters. respectively."," The aperture magnitudes were corrected for atmospheric extinction, assuming $A_V=0.14$ and $A_I=0.08$ as extinction coefficients for the $V$ and $i$ filters, respectively."
+ As shown in Fig. 2..," As shown in Fig. \ref{calib},"
+ a straight line well reproduces the calibration equations., a straight line well reproduces the calibration equations.
+ As the seeing and the overall observing conditions were stable during the run. the slopes of the calibration equations for each observing run and for each filter have been computed using the data from all the nights.," As the seeing and the overall observing conditions were stable during the run, the slopes of the calibration equations for each observing run and for each filter have been computed using the data from all the nights."
+ As it can be seen in Table 2.. the standard deviations of the calibration constants for each run and filter is 0.015mag. corroborating our assumption that all nights were photometric. and that we can assume a constant slope for each filter and run.," As it can be seen in Table \ref{calib_par}, the standard deviations of the calibration constants for each run and filter is $0.015$ mag, corroborating our assumption that all nights were photometric, and that we can assume a constant slope for each filter and run."
+ Standard stars for which previous problems were reported (PG 1047C. RUI49A. RUI49G. PGI323A: see Johnson Bolte 1998)) were excluded. as well as saturated stars. those close to a cosmic ray. ete... After this cleaning. the mean slope was computed. and finally the different night constants were found using this slope to fit the individual data. night by night.," Standard stars for which previous problems were reported (PG 1047C, RU149A, RU149G, PG1323A; see Johnson Bolte \cite{johnsonbolte98}) ) were excluded, as well as saturated stars, those close to a cosmic ray, etc... After this cleaning, the mean slope was computed, and finally the different night constants were found using this slope to fit the individual data, night by night."
+ The adopted values are presented in Table 2.., The adopted values are presented in Table \ref{calib_par}.
+ The typical errors (ris) are also given., The typical errors $rms$ ) are also given.
+ The calibration curves are shown in Fig., The calibration curves are shown in Fig.
+ 2. for both runs., \ref{calib} for both runs.
+ In this figure. the represents the best fitting equation. while the is obtained by best fitting the data imposing the adopted mean slope.," In this figure, the represents the best fitting equation, while the is obtained by best fitting the data imposing the adopted mean slope."
+ The two lines are almost overlapping., The two lines are almost overlapping.
+ The mean number of standard star measures used, The mean number of standard star measures used
+Over the last decade. it has become clear that a sizeable. perhaps dominant. part of the stellar halo of the Milky Way is composed of the accumulated debris from the disruptior of distinet satellite galaxies.,"Over the last decade, it has become clear that a sizeable, perhaps dominant, part of the stellar halo of the Milky Way is composed of the accumulated debris from the disruption of distinct satellite galaxies."
+ Starting with the discovery and characterization. of the tidal debris from the disruption of the Sagittarius dwarf galaxy (bataetal.1995;Majewski 2003).. a variety of other stellar halo overdensities have beer discovered: the low-latitude stream (Yannyetal.2003:Ibataetal. 2003).. the Virgo and Hercules-Aquila overdensities (Duffauetal.2006:Jurié2008;Belokurov 2007a).. and a number of smaller tidal streams (Odenkirchenetal.2001:Yanny&Willett 2009).," Starting with the discovery and characterization of the tidal debris from the disruption of the Sagittarius dwarf galaxy \citep{ibata95,majewski03}, a variety of other stellar halo overdensities have been discovered: the low-latitude stream \citep{yanny03,ibata03}, the Virgo and Hercules-Aquila overdensities \citep{duffau06,juric08,her_aq}, and a number of smaller tidal streams \citep{pal5,grillmair_orphan,grillmair_stream,grillmair_n5466,orphan,grillmair08,newberg09}."
+. Models of stellar halos that formed in a cosmological context through the disruption of dwarf galaxies have been developed and refined (Bullocketal.2001:Bul-lock&Johnston2005;Abadietal.2006:Cooper 2010).. providing quantitative predictions for the structure and stellar content of stellar halos formed in such a fashion.," Models of stellar halos that formed in a cosmological context through the disruption of dwarf galaxies have been developed and refined \citep{bkw,bullock05,abadi06,cooper10}, providing quantitative predictions for the structure and stellar content of stellar halos formed in such a fashion."
+ Remarkably. the structure and degree of substructure of the model stellar halos appear to be in quantitative agreement with observations of the Milky Way's stellar halo (Bullock&Johnston2005.. Belletal. 2008): stellar halos with ~IO?L... a roughly επ density profile. and showing rich substructure on a variety of spatial scales.," Remarkably, the structure and degree of substructure of the model stellar halos appear to be in quantitative agreement with observations of the Milky Way's stellar halo \citealp{bullock05}, \citealp{bell08}) ): stellar halos with $\sim 10^9 L_{\sun}$, a roughly $r^{-3}$ density profile, and showing rich substructure on a variety of spatial scales."
+ A clear prediction of such a picture is that any population differences in the stellar halo. arising from different progenitor populations. should have a similar morphology to the stellar density inhomogeneities.," A clear prediction of such a picture is that any population differences in the stellar halo, arising from different progenitor populations, should have a similar morphology to the stellar density inhomogeneities."
+ The metallicity—mass correlation of satellite galaxies will translate into variations in stellar halo metallicity. as debris from larger satellites will have higher metallicity than those of lower-luminosity satellites or globular clusters.," The metallicity–mass correlation of satellite galaxies will translate into variations in stellar halo metallicity, as debris from larger satellites will have higher metallicity than those of lower-luminosity satellites or globular clusters."
+ Furthermore. one expects to see distinctive signatures in age and detailed element abundance patterns. reflecting when satellites were acereted (Robertsonetal.2005.. Fontetal. 2006:; see also Tumlinson 20100).," Furthermore, one expects to see distinctive signatures in age and detailed element abundance patterns, reflecting when satellites were accreted \citealp{robertson05}, \citealp{font06}; see also \citealp{tumlinson10}) )."
+ Signatures of population inhomogeneity have been detected: hints of chemical signatures of kinematically-detected local streams (e.g..Helmietal.2006).. small variations in metallicity within the nearby parts of the stellar halo. especially associated with the low-latitude stream (e.g.Iveziéetal.2008)ue eene tendency SUPETIRDeSECTor The Omnarod fo REYnave meku-poorer slsreds ar arger radii (e.g..Carolloetal.2007).. and a recent color-magnitude diagram (CMD) fitting analysis of Sloan Digital Sky Survey (SDSS) stripes showing both a radial metallicity gradient and large density and metallicity variations around that gradient (deJongetal.2010).," Signatures of population inhomogeneity have been detected: hints of chemical signatures of kinematically-detected local streams \citep[e.g.,][]{helmi06}, small variations in metallicity within the nearby parts of the stellar halo, especially associated with the low-latitude stream \citep[e.g.][]{ivezic08}, the tendency for the halo to have metal-poorer stars at larger radii \citep[e.g.,][]{carollo07}, and a recent color–magnitude diagram (CMD) fitting analysis of Sloan Digital Sky Survey (SDSS) stripes showing both a radial metallicity gradient and large density and metallicity variations around that gradient \citep{dejong10}."
+. Yet. for the most part. the observational signatures of such population variations are accessible for relatively bright stars only (high. S/N ugriz imaging. or moderate S/N spectra).," Yet, for the most part, the observational signatures of such population variations are accessible for relatively bright stars only (high S/N $ugriz$ imaging, or moderate S/N spectra)."
+ Accordingly. those radii in our own stellar halo where substructure is rich CIO<<40 kkpe)are not probed particularly well.," Accordingly, those radii in our own stellar halo where substructure is rich $10<r<40$ kpc) are not probed particularly well."
+ Some of the most impressive evidence for stellar halo population variations are from M31. where significant variations in stellar age and metallicity have been documented (Brownetal.2006:Richardson 2008).," Some of the most impressive evidence for stellar halo population variations are from M31, where significant variations in stellar age and metallicity have been documented \citep{brown06,richardson08}."
+. In particular. [bataetal.(2007) and MeConnachieetal.(2009) have used the color of the red giant branch stars (practically inaccessible in the Milky Way's halo because distances are uncertain) to provide panoramic views of the stellar populations in the M31 halo. showing clearly that stellar population variation can be associated," In particular, \citet{ibata07} and \citet{mcconnachie09} have used the color of the red giant branch stars (practically inaccessible in the Milky Way's halo because distances are uncertain) to provide panoramic views of the stellar populations in the M31 halo, showing clearly that stellar population variation can be associated"
+This bright object is listed in the Luvteu Double Star catalogue aud has a separation of 127 arcsec from its colupanion star.,This bright object is listed in the Luyten Double Star catalogue and has a separation of 127 arcsec from its companion star.
+ The hot subdwart has a Iipparcos parallax of 1541.75 niniülliarcssecouds (ias). however the reliability of this value is uncertain.," The hot subdwarf has a Hipparcos parallax of $\pm$ milliarcsseconds (mas), however the reliability of this value is uncertain."
+ The associated nucertaity is too high to meaniugfullv coustrain the stars distance and mass m any case., The associated uncertainty is too high to meaningfully constrain the star's distance and mass in any case.
+" The correspondiug Ilipparcos proper iiotions of jt,cosd=83 2nunuas/yvr iib gs=δΙ lunnas/vr.", The corresponding Hipparcos proper motions of $\mu_\alpha \cos\delta = 83\pm 2$ mas/yr and $\mu_\delta = -84\pm 1$ mas/yr.
+" The companion star is faint bv comparison it has SDSS photometry. with — (21.07I. 18.512. 17.097. 16.138. 15.635) id has proper niotkis of fi,cosÓ=SOc Duuuas/vr Had pg=8st nuuas/vr. almost an exact match."," The companion star is faint by comparison – it has SDSS photometry, with = (21.074, 18.512, 17.097, 16.138, 15.635) – and has proper motions of $\mu_\alpha \cos\delta = 80\pm 3$ mas/yr and $\mu_\delta = -83\pm
+3$ mas/yr, almost an exact match."
+ The proper motions from SuperCosmos of stars with 150 ‘aresecouds of 11777 (prdare shown in Figure2 2.., The proper motions from SuperCosmos of stars with 150 arcseconds of 1777 are shown in Figure \ref{fig:BDp48PM}.
+ The two stars (at jp~~100 nnuas/yr). are clearly separated from the main group around the origin., The two stars (at $\mu_\delta\sim -100$ mas/yr) are clearly separated from the main group around the origin.
+ Note that the SuperCosiios. proper motious shown in the figure are somewhat hieher than those from the more reliable and precise UCAC?2 catalogue quoted above., Note that the SuperCosmos proper motions shown in the figure are somewhat higher than those from the more reliable and precise UCAC2 catalogue quoted above.
+ 11777 is a reasonably well-studied object aud we are fortunate to have access to a spectrum. taken using CAFOS| on the Calar Alto 2.2nuu with resolution aud covering waveleugth raugeAA.," 1777 is a reasonably well-studied object and we are fortunate to have access to a spectrum, taken using CAFOS on the Calar Alto m with resolution and covering wavelength range."
+. A model fit of this spectrmu kindly carried out by Uli Heber (aud shown in Figure 3)) usine the NLTE model evid of Stróeretal. (2007).., A model fit of this spectrum kindly carried out by Uli Heber (and shown in Figure \ref{fig:BDp48fit}) ) using the NLTE model grid of \citet{Stroeer2007}. .
+ The best-fit stellar given iu Table ::, The best-fit stellar parameters are given in Table \ref{tab:BDp48}.
+ The star is a “garden- parametershelimm-richare subdwarf O star aud sits with the majority ofthese objects in the TR diagram at around 15 ΤΙΝ and logg756., The star is a ``garden-variety'' helium-rich subdwarf O star and sits with the majority of these objects in the HR diagram at around $\sim$ K and $\log g\sim 5-6$.
+" The extinction iu the direction of the svstem is low (0,01 mag}.", The extinction in the direction of the system is low $\sim$ 0.04 mag).
+ The SDSS colours of the companion star are consistent with it beige an MO dwart., The SDSS colours of the companion star are consistent with it being an M0 dwarf.
+ If we use he empirical relations of ILwwvlewetal.(2002) we Bud a photometric parallax of (/ AL;)z7.9 uae.," If we use the empirical relations of \citet{Hawley2002} we find a photometric parallax of $(i- M_i)
+\approx 7.9$ mag."
+ The; distance. to the star is. therefore. z: ppc., The distance to the star is therefore $\approx$ pc.
+" Thisqi in turus leads AEto the absoluteA visual""ue imaguitude of tsli777 of Myc 2.5nuuag. using nennenV-—10.1of Or sdO "," This in turns leads to the absolute visual magnitude of 1777 of $_V \approx$ mag, using $V =
+10.7$ for the sdO star."
+We uote that these values are very xelininary., We note that these values are very preliminary.
+the star.Unfortunatelv the spectruni we have of he Te-sdO is not flux calibrated. which makes mass determination forthe star currently difficult.," Unfortunately the spectrum we have of the He-sdO is not flux calibrated, which makes mass determination forthe star currently difficult."
+function (?).. among CDÀM-halo catalogues (2).. and among semianalvtical galaxy catalogues (?)..,"function \citep{Schmidt01}, among CDM-halo catalogues \citep{Arbabi02}, and among semianalytical galaxy catalogues \citep{Benson03}."
+ Here we could provide an analytical description of the void size distribution given bv Eq. (, Here we could provide an analytical description of the void size distribution given by Eq. (
+3).,3).
+ The similarity. relation Eq. (, The similarity relation Eq. (
+4) is à cirect consequence of such a void. size distribution.,4) is a direct consequence of such a void size distribution.
+ The scaling of the median sizes is comparable to our earlier findings taking into account the transition to the 3-cimensional survey geometry., The scaling of the median sizes is comparable to our earlier findings taking into account the transition to the 3-dimensional survey geometry.
+ The sealing relation. (i.e. το constant of proportionality) of void sizes 2oxA ds Latter than [ound by 22., The scaling relation (i.e. the constant of proportionality) of void sizes $R \propto \lambda$ is flatter than found by \citet{Benson03}.
+ We suspect this is due to our procedure of identifving both Large and small voids with the same algorithm. not just looking inside of wall galaxies of fixed local overdensity as other void finclers do.," We suspect this is due to our procedure of identifying both large and small voids with the same algorithm, not just looking inside of wall galaxies of fixed local overdensity as other void finders do."
+ The scaling relation is a representation of the self-similarity of the void phenomenon., The scaling relation is a representation of the self-similarity of the void phenomenon.
+ In distinction to earlier. analyses. we have a large sample of voids and therefore a good statistics in determining the void size distribution.," In distinction to earlier analyses, we have a large sample of voids and therefore a good statistics in determining the void size distribution."
+ “Phe flat) scaling relation is a consequence of relatively small. voids among the bright. galaxy samples with large A and conversely. comparatively large. voids among faint galaxies.," The flat scaling relation is a consequence of relatively small voids among the bright galaxy samples with large $\lambda$ and conversely, comparatively large voids among faint galaxies."
+" ""This is a new property well established in the densely sampled 2dPCGIS and in the large. Millenium: simulation.", This is a new property well established in the densely sampled 2dFGRS and in the large Millenium simulation.
+ Lt is the same void phenomenon discussed in the stimulating paper by ον, It is the same void phenomenon discussed in the stimulating paper by \citet{Peebles01}.
+ However. he claims a significant cliscrepaney of the void phenomenon and the standard. hierarchical clustering phenomenon.," However, he claims a significant discrepancy of the void phenomenon and the standard hierarchical clustering phenomenon."
+ Contrary to this claim we established. void statistics that are well reproduced by the Millenium galaxy catalogues., Contrary to this claim we established void statistics that are well reproduced by the Millenium galaxy catalogues.
+ As shown in Fig., As shown in Fig.
+ 7 the size distribution of voids in simulated. galaxies is well described. by a self-similar distribution for small anc intermediate sized. voids. and a more extended. tail of large relative void sizes with respect to the mean galaxy separation.," \ref{cummill_norm} the size distribution of voids in simulated galaxies is well described by a self-similar distribution for small and intermediate sized voids, and a more extended tail of large relative void sizes with respect to the mean galaxy separation."
+ This means that the void phenomenon discussed by ? is a genuine consequence of the underlving LCDM model., This means that the void phenomenon discussed by \citet{Peebles01} is a genuine consequence of the underlying LCDM model.
+ An earlier demonstration of the void phenomenon within the hierarchical CDAL clustering model by quite dilferent methods was provided by 2.., An earlier demonstration of the void phenomenon within the hierarchical CDM clustering model by quite different methods was provided by \citet{Mathis02}.
+ The new results concern the precise verification of the void scaling relations for the 2dkGRS survey ancl the detailed fitting of the void size distribution by a two-branch exponential distribution., The new results concern the precise verification of the void scaling relations for the 2dFGRS survey and the detailed fitting of the void size distribution by a two-branch exponential distribution.
+ Both properties of the 2dECGINS are surprisingly well reproduced in the semianalytical galaxy formation scheme of the Millenium: simulation (?).., Both properties of the 2dFGRS are surprisingly well reproduced in the semianalytical galaxy formation scheme of the Millenium simulation \citep{Croton05}.
+ The comparison is done in the 24hCRS survey mask and in recshift space., The comparison is done in the 2dFGRS survey mask and in redshift space.
+ This is essential for the precision comparison since we found coinciding void size distributions both for small and median void sizes. but an increase of the largest voids in redshift’ space by about10'4.. and an overall reduction of the void size distribution in the 0CARS survey mask as conipared to the complete 500 simulation box.," This is essential for the precision comparison since we found coinciding void size distributions both for small and median void sizes, but an increase of the largest voids in redshift space by about, and an overall reduction of the void size distribution in the 2dFGRS survey mask as compared to the complete 500 simulation box."
+ We established a clear representation of the color bimodality of void galaxies and an lareer fraction of blue galaxies in the 2PdkFORS voids., We established a clear representation of the color bimodality of void galaxies and an larger fraction of blue galaxies in the 2dFGRS voids.
+ An increase of the fraction of blue ealaxies in. voids has been previously established by ο for the 2dbORS and by ο [for the SDSS.," An increase of the fraction of blue galaxies in voids has been previously established by \citet{Hoyle04} for the 2dFGRS and by \citet{Rojas04, Rojas05} for the SDSS."
+ Also. ?. found a dependenceof the blue fraction on the Iarge-scale density environment of a scale of ShMpc.," Also, \citet{Croton05} found a dependenceof the blue fraction on the large-scale density environment of a scale of $8 \hMpc$."
+ 2 investigated the SDSS for the environment. depencence of star formation activity of galaxies on a more local scale., \citet{Kauffmann04} investigated the SDSS for the environment dependence of star formation activity of galaxies on a more local scale.
+ Vhey found that the most sensitive environmental dependence of galaxies is the SER. which is strongest lor galaxies with stellar masses «8LAL. but no much environmental dependence on the large scale density was found when the small scale density was spoecilied.," They found that the most sensitive environmental dependence of galaxies is the SFR, which is strongest for galaxies with stellar masses $< 3 \times 10^{10} M_{\sun}$, but no much environmental dependence on the large scale density was found when the small scale density was specified."
+ A bimodality of the color distribution of void. galaxies in the SDSS has been established by 2.., A bimodality of the color distribution of void galaxies in the SDSS has been established by \citet{Patiri06a}.
+ Our results for the 3011.1 are new and extend these earlier findings., Our results for the 2dFGRS are new and extend these earlier findings.
+ As in that paper. we confirm the stability of the red. sequence in void regions. but we find a slight blue shift of the blue cloud.," As in that paper, we confirm the stability of the red sequence in void regions, but we find a slight blue shift of the blue cloud."
+ Furthermore. there is a weak tendency of void galaxies to have a higher SER. and for more early type (12) than late type (9) galaxies.," Furthermore, there is a weak tendency of void galaxies to have a higher SFR, and for more early type (E) than late type (S) galaxies."
+ A reason for the blue slight. shift of the blue cloud. established bere but not by ο in the SDSS analysis may be that we analyzed the central void regions. while they considered the whole void. volume.," A reason for the blue slight shift of the blue cloud established here but not by \citet{Patiri06a} in the SDSS analysis may be that we analyzed the central void regions, while they considered the whole void volume."
+ Also the 2dECGIUS is denser sampled. and. therefore we have a better statistics of void galaxies.," Also the 2dFGRS is denser sampled, and therefore we have a better statistics of void galaxies."
+ While the SDSS colors have a better photometric accuracy. it was shown by ? that there is only a small non-linearity between the APAL and SDSS colors.," While the SDSS colors have a better photometric accuracy, it was shown by \citet{Norberg02} that there is only a small non-linearity between the APM and SDSS colors."
+ Some dillerence might be produced by dilferences in the adopted &- ancl e-corrections., Some difference might be produced by differences in the adopted $k$ - and $e$ -corrections.
+ Color distributions at. higher redshifts show a similar bimocal behavior as our results., Color distributions at higher redshifts show a similar bimodal behavior as our results.
+ By comparing the SDSS color distribution with DEEP? results. ? found that the blue cloud is shifted bluewards by 0.6 mag at redshift +=1.," By comparing the SDSS color distribution with DEEP2 results, \cite{Blanton06} found that the blue cloud is shifted bluewards by 0.6 mag at redshift $z=1$."
+ Of course. this shift is stronger than what we find here. but it indicates a similar trend.," Of course, this shift is stronger than what we find here, but it indicates a similar trend."
+ Therefore a natural interpretation would be a slower evolution of structures inside voids and a vounger galaxy population as compared to the field., Therefore a natural interpretation would be a slower evolution of structures inside voids and a younger galaxy population as compared to the field.
+ It the color shift might be due to à vounger population in underdense regions. we would expect the shift in the blue peak to be stronger for the lower magnitude bin.," If the color shift might be due to a younger population in underdense regions, we would expect the shift in the blue peak to be stronger for the lower magnitude bin."
+ However due to the low number of void galaxies in the fainter magnitude bin. we cannot establish such a tendency.," However due to the low number of void galaxies in the fainter magnitude bin, we cannot establish such a tendency."
+ The color shift. is of similar amount than the color uncertaintv of Om= 0.09., The color shift is of similar amount than the color uncertainty of $\delta m=0.09$ .
+ But an uncertainty would, But an uncertainty would
+is asvinhiolic recurrent nova (IN) which had previously undergone recorded outbursts in 1898. 1933. 1958. 1967 and 1985 (seeRosino1937:&lima1987).. with possible additional outbursts in 1907. (Schaefer2004) and 1945 1993).,"is a symbiotic recurrent nova (RN) which had previously undergone recorded outbursts in 1898, 1933, 1958, 1967 and 1985 \citep[see][]{ros87,rosi87}, with possible additional outbursts in 1907 \citep{sch04} and 1945 \citep{opp93}."
+. On 2006 February 12.89UT it was observed (ο be undergoing a further eruption, On 2006 February 12.83UT it was observed to be undergoing a further eruption
+[actor: &—0.411.——0.063ρα under the assumption↽− ofom s-s symmetry.,"factor $\kappa = 0.477 \:
+^{+\:0.063}_{-\:0.053}$ under the assumption of $s$ $\bar s$ symmetry."
+" Hf: allowing""E s-5 asymmetry. then ⊔∐↲⋯≸↽↔↴↕∖↽≼↲⋝∖⊽∣↘⇁∶∩⋅↱≻≡↽⊰↻⋮⋮∣↓∣∎−∎⋅⋝⊺↥∐⊳∖⇁∖∖⊽≀↧↪∖⇁≼↲≺∢∐"," If allowing $s$ $\bar s$ asymmetry, then the fit gives $\kappa = 0.536 \: ^{+\:0.109}_{-\:0.079}$."
+"⋯↲≼⇂∣↽≻∡∖⇁∐∐↲⋟∖⇁∪∐≺∢↥∪⋯⊔⊔⋯⇂≼↲↥⋝∖⊽⋅⊳∖⊽∏≺∢∐≀↧↪∖⊽∐≼↲↓⋅⋓∥∖∖↽↥∐≺∢∐ , . ⋅⋅T . ⋅"," This was echoed by meson cloud models, such as \cite{car00}, , which gives $\kappa=0.55$ ."
+" ↴∏∐↲≼⇂≼↲↓∎↓∐∐↕∪∐∪↓⋟∣↘⇁↕⋝∖⊽ Since both experiment [0]. and our model caleulation found that the strange sea .r-cependence is similar to that of the non-strange sea. the & reflects roughly the ratio of strange quark content relative to (he non-strange sea quarks. ie. Then the constraint &20.5. together with {ντDP,;=12% and Pyg+Pugt+Pos<1. limits the strange quark content in the proton to be in the range of P;=(3~20)%."," The definition of $\kappa$ is Since both experiment \cite{baz95} and our model calculation found that the strange sea $x$ -dependence is similar to that of the non-strange sea, the $\kappa$ reflects roughly the ratio of strange quark content relative to the non-strange sea quarks, i.e., Then the constraint $\kappa\approx 0.5$, together with $P_{d\bar
+d}-P_{u\bar u} =12\%$ and $P_{d\bar d}+P_{u\bar u}+P_{s\bar s} < 1$ , limits the strange quark content in the proton to be in the range of $P_{s\bar s} = (3\sim 20)\%$."
+ The lower limit D=3% corresponds to Py= 0., The lower limit $P_{s\bar s} = 3\%$ corresponds to $P_{u\bar u}=0$ .
+ In our model caleulation we assume that the wuddd system stays in the configuration of the lowest energy [22] where the dis in its ground state and the wadd subsystem has mixed orbital symmetry. [31]y.," In our model calculation we assume that the $uudd\bar{d}$ system stays in the configuration of the lowest energy \cite{Helminen}
+ where the $\bar{d}$ is in its ground state and the $uudd$ subsystem has mixed orbital symmetry $[31]_{X}$ ."
+ This configuration gives the possibility of 1/4 for d to be in P-wave hence leads naturally (ο spin flip., This configuration gives the possibility of 1/4 for $d$ to be in P-wave hence leads naturally to spin flip.
+ Considering the electroproduction process οΔΑ)—aN in the five-quark component model of nucleon. the basic eleetroproduction reaction is +q(p)—x4qp). where q' represents the final quark which joins in the remaining part of the nucleon to form the hadronic state X. as illustrated in Fig.," Considering the electroproduction process $\gamma N^{\uparrow}\to
+\pi X$ in the five-quark component model of nucleon, the basic elcetroproduction reaction is $\gamma + q(p) \to \pi + q'(p')$, where $q'$ represents the final quark which joins in the remaining part of the nucleon to form the hadronic state X, as illustrated in Fig."
+ 1., 1.
+ The vw and d can be taken as the probed quarks in thenucleon. while (he rest partons canbe taken as the spectators collectively in thescattering process.," The $u$ and $d$ can be taken as the probed quarks in thenucleon, while the rest partons canbe taken as the spectators collectively in thescattering process."
+ The model assumes thatthe, The model assumes thatthe
+In this work we propose to combine Fisher matrix orecasts with Box-Cox transformations of parameter space o obtain accurate expectations of posterior. cüstributions.,In this work we propose to combine Fisher matrix forecasts with Box-Cox transformations of parameter space to obtain accurate expectations of posterior distributions.
+ Box&Cox(1964) introduced. a parametrisecl set of »ower transformations with which a wide range of data can be transformed to follow a Gaussian. distribution to good approximation., \citet{box64} introduced a parametrised set of power transformations with which a wide range of data can be transformed to follow a Gaussian distribution to good approximation.
+ We will apply these transformations o modcl parameters. instead. of data. in order to naioclily a given posterior into a multivariate Gaussian distribution or which a Fisher matrix analysis is exact.," We will apply these transformations to model parameters, instead of data, in order to modify a given posterior into a multivariate Gaussian distribution for which a Fisher matrix analysis is exact."
+ After an inverse Box-Cox transformation the Fisher matrix results will then accurately describe the original posterior., After an inverse Box-Cox transformation the Fisher matrix results will then accurately describe the original posterior.
+ To determine the ree parameters of the Box-C'ox transformation. the original »osterior needs to be sampled. and hence an initial mock ikelihood analysis to be run.," To determine the free parameters of the Box-Cox transformation, the original posterior needs to be sampled, and hence an initial mock likelihood analysis to be run."
+ We will demonstrate this method with an example from cosmology., We will demonstrate this method with an example from cosmology.
+ Several ambitious surveys are currently planned or that are going to measure the parameters of the cosmological standard: model. particularly those of dark matter and dark energy. with high precision.," Several ambitious surveys are currently planned or that are going to measure the parameters of the cosmological standard model, particularly those of dark matter and dark energy, with high precision."
+ These experiments will investigate several cosmological probes of the [large-scale structure of the Universe. the potentially most powerful one being weak gravitational lensing of clistant galaxies (Albrechtctal.2006:Peacock2006).," These experiments will investigate several cosmological probes of the large-scale structure of the Universe, the potentially most powerful one being weak gravitational lensing of distant galaxies \citep{albrecht06,peacock06}."
+". Weak lensing features a characteristic non-linear degeneracy between the twobest-constrained parameters. Qu, (mean matter density) and os (normalisation of matter density Huetuations) as they both govern the overall aniplitucle of the signal (seee.g.Lockstraetal.2006:Schrabbacketal. 2010)."," Weak lensing features a characteristic non-linear degeneracy between the twobest-constrained parameters $\Omega_{\rm m}$ (mean matter density) and $\sigma_8$ (normalisation of matter density fluctuations) as they both govern the overall amplitude of the signal \citep[see e.g.][]{hoekstra06,schrabback09}."
+. Hence a mock weak lensing survey provides an excellent test case. but we emphasise that the method outlined is applicable to any prediction for model parameter constraints.," Hence a mock weak lensing survey provides an excellent test case, but we emphasise that the method outlined is applicable to any prediction for model parameter constraints."
+ Thepaper is organised as follows., Thepaper is organised as follows.
+ Section 2. details he principles of Box-Cox transformations. our strategies to determine optimal Box-Cox parameters.and the combined Dox-Cox-Fisher. formalism.," Section \ref{sec:theory} details the principles of Box-Cox transformations, our strategies to determine optimal Box-Cox parameters,and the combined Box-Cox-Fisher formalism."
+ 1n. Section 3 we investigate he performance of the proposed. method for a mock weak ensing experiment. comparing different variants in the implementation and quantifving the universality of the Box-Cox-Fisher formalism.," In Section \ref{sec:performance} we investigate the performance of the proposed method for a mock weak lensing experiment, comparing different variants in the implementation and quantifying the universality of the Box-Cox-Fisher formalism."
+ We apply this formalism to a test of he cegencracy-breakingὃνo capabilities of weak lensinge higher->order statistics in Section +.. before we summarise and conclude on our findings in Section 5..," We apply this formalism to a test of the degeneracy-breaking capabilities of weak lensing higher-order statistics in Section \ref{sec:application}, before we summarise and conclude on our findings in Section \ref{sec:conclusions}."
+ Power transformations such as the inverse ancl square-root ransformation. or logarithmic transformations are popular choices to render the distribution of data more Gaussian.," Power transformations such as the inverse and square-root transformation, or logarithmic transformations are popular choices to render the distribution of data more Gaussian."
+ ‘The Box-Cox transformation unites these cases with a single ree parameter per dimension and are hence widely. used in various areas of science., The Box-Cox transformation unites these cases with a single free parameter per dimension and are hence widely used in various areas of science.
+ Astrophysical applications are rare: one example is the work by Dineen&Coles(2005). who ested cosmic microwave background data for Gaussianity., Astrophysical applications are rare; one example is the work by \citet{dineen05} who tested cosmic microwave background data for Gaussianity.
+" For ao N,-climensional variable p. the Dox-Cox transformation in each dimension (p;—1....;N, reads (Box (1) μίαdy) =.anherclbenormalisation αἱ A,=0."," For a $N_p$ -dimensional variable $\vek{p}$ the Box-Cox transformation in each dimension $\mu=1,\,..\,,N_p$ reads \citep{box64}
+ \eq{ } _μ(λ_μ,a_μ) = where the normalisation has been chosen such that the transformation is continuous in the parameter $\lambda_\mu$ at $\lambda_\mu=0$."
+ We allow for a shift αμ as a second free parameter in each cimension., We allow for a shift $a_\mu$ as a second free parameter in each dimension.
+ Note that we denote. transformed quantities by a. bar. ancl drop the dependence on. the Box-Cox parameters (A.a) unless it needs to be mace explicit.," Note that we denote transformed quantities by a bar and drop the dependence on the Box-Cox parameters $\br{\vek{\lambda},\vek{a}}$ unless it needs to be made explicit."
+" Usually. equation (1) is applied to the elements of a datavector. but we will henceforth understand. p, as the parameters of an IN,dimensional parameter space."," Usually, equation ) is applied to the elements of a datavector, but we will henceforth understand $p_\mu$ as the parameters of an $N_p$ -dimensional parameter space."
+ Then the transform of a given posterior distribution P(p) is given bv with the Jacobian The second equality follows cirectly from equation (1)., Then the transform of a given posterior distribution ${\cal P}(\vek{p})$ is given by with the Jacobian The second equality follows directly from equation ).
+" “Phe first goal is to determine the set of 2.N,, parameters. (A.a). such that the transformed posterior. (p). is a multivariate Gaussian to good approximation."," The first goal is to determine the set of $2 N_p$ parameters, $\br{\vek{\lambda},\vek{a}}$, such that the transformed posterior, $\bar{\cal P}(\bar{\vek{p}})$, is a multivariate Gaussian to good approximation."
+" Suppose a rancom sample p with η elements. Le. LasτνPho] For everv p=1.25.N,. ds drawn from the posterior (p). lor instance via Monte-Carlo sampling techniques."," Suppose a random sample $\vek{\hat{p}}$ with $n$ elements, i.e. $\bc{\hat{p}_{\mu,1},\,..\,,\hat{p}_{\mu,n}}$ for every $\mu=1,\,..\,,N_p$, is drawn from the posterior ${\cal P}(\vek{p})$, for instance via Monte-Carlo sampling techniques."
+ Lf the. Box-Cox transformed. posterior is indeed Gaussian. the distribution is given by and has only the Box-Cox parameters. and the mean py... απο Covariance of the Gaussian as [ree parameters.," If the Box-Cox transformed posterior is indeed Gaussian, the distribution is given by and has only the Box-Cox parameters, and the mean $\bar{\vek{p}}_{\rm max}$ and covariance of the Gaussian as free parameters."
+ Since (p) is assumed Gaussian. one can employ the stancard maximum: likelihood estimators for the covariance and mean.," Since $\bar{\cal P}(\bar{\vek{p}})$ is assumed Gaussian, one can employ the standard maximum likelihood estimators for the covariance and mean."
+" The latter simply implies p= p,,,,- ��ο that the exponential in (4) is unity."," The latter simply implies $\bar{\vek{p}}=\bar{\vek{p}}_{\rm max}$ , so that the exponential in ) is unity."
+ Consequently. one obtains the following concentrated: for the Box-Cox parameters (for details 1993). up to an irrelevant constant.," Consequently one obtains the following concentrated log-likelihood for the Box-Cox parameters <cit.>[for details , up to an irrelevant constant."
+ We have added the subscript AIL to emphasise that the maximum likelihood estimate for, We have added the subscript ML to emphasise that the maximum likelihood estimate for
+Within a cluster. tidal interactions aud the ICAL wind can strip lass from a ealaxws disk.,"Within a cluster, tidal interactions and the ICM wind can strip mass from a galaxy's disk."
+ We estimate the οπου of both on IC 23118. noting that the lack of an appropriateτς neiehbor excludes ealaxv-ealaxy tidal interactions.," We estimate the effects of both on IC 3418, noting that the lack of an appropriate neighbor excludes galaxy-galaxy tidal interactions."
+ The ID tin IC 3118's disk was extremely susceptible to RPS. as expected for a faint. low surface brightness galaxy orbiting in a iassive cluster.," The H in IC 3418's disk was extremely susceptible to RPS, as expected for a faint, low surface brightness galaxy orbiting in a massive cluster."
+ Iu contrast. IC 3118s probable tidal radius is bevond its stellar disk. ruling out a tidal origin for the tail.," In contrast, IC 3418's probable tidal radius is beyond its stellar disk, ruling out a tidal origin for the tail."
+ The dense molecular clouds were unlikely to be RP stripped as clouds., The dense molecular clouds were unlikely to be RP stripped as clouds.
+ We first estimate the impact of RPS using the stripping condition of Cama&Cott(1972).. We use the usual 9 profile for the ICAL density aud adopt values of the profile slope. .. aud the ceutral density. py. for Virgo from Sandersonetal.(2003).," We first estimate the impact of RPS using the stripping condition of \citet{Gunn72}, We use the usual $\beta$ profile for the ICM density and adopt values of the profile slope, $\beta$, and the central density, $\rho_0$ , for Virgo from \citet{Sanderson03}."
+. Asstuing a three-dimensional distance from the clusters core of Y2rprojecrod ANC an orbital velocity of 1000luis +. the rau IS Peau _=2PICACrypitalxessureeeI0MpePkus 7.," Assuming a three-dimensional distance from the cluster's core of $\sqrt{2}r_{\rm projected}$ and an orbital velocity of $v_{\rm orbital}=1000~\rm{km~s^{-1}}$ , the ram pressure is $P_{\rm ram}=\rho_{\rm ICM} v_{\rm orbital}^2\approx40~\rm{M_{\odot}~pc^{-3}~km^2~s^{-2}}$ ."
+ To estimate the resoring pressure on the eas disk. 2vost2xGo.oj. we first determine the stellar surface cdeusity. o.. from (--banud observations.," To estimate the restoring pressure on the gas disk, $P_{\rm rest}=2\pi G\sigma_{\ast}\sigma_{\rm H}$, we first determine the stellar surface density, $\sigma_{\ast}$ , from -band observations."
+ We uext estimate the initial gas surface density. ay. by απλο ai pre-strippine TD mass aud disk size based on IC 115 morphology. optical Iuniuosity. aud optical radius.," We next estimate the initial gas surface density, $\sigma_{\rm H}$, by assuming a pre-stripping H mass and disk size based on IC 3418's morphology, optical luminosity, and optical radius."
+ The Z[-—band magnitude aud effective surface brightness of IC 3118 are mq12.72 aud He=22.2]magarcsec2 (CGavazziotal.2000.2003).," The -band magnitude and effective surface brightness of IC 3418 are $m_{\rm H} = 12.72$ and $\mu_e=22.24~{\rm mag~arcsec}^{-2}$ \citep{Gavazzi00, Gavazzi03}."
+". Assundug a dass to light ratio of 0.78 ον 2008).. 0.=12M.pe7 and AL,=105? ALL. etal.(2005). report an II deficiency (see Cüovanelli&Ilavues (1983))) aud au upper limit ou the IT mass that correspond to an expected LOSAL. of IE in IC 3118."," Assuming a mass to light ratio of 0.78 \citep{Kirby08}, $\sigma_{\ast} = 12~\rm{M_{\odot}~pc^{-2}}$ and $M_{\ast} = 10^{8.9}~{\rm M}_{\odot}$ \citet{Gavazzi05} report an H deficiency (see \citet{Giovanelli83}) ) and an upper limit on the H mass that correspond to an expected $10^{8.8}~{\rm M}_{\odot}$ of H in IC 3418."
+ We estimate that IC. 3118. contained of its ID within 1.85 by couverting its UCC optical radius to a Wolubere radius (Dickel&Rood1978) aud taking au average ratio of the IT to optical disk size from Cüiovauclli&Ilavnues (1983).., We estimate that IC 3418 contained of its H within $1\arcmin.85$ by converting its UGC optical radius to a Holmberg radius \citep {Dickel78} and taking an average ratio of the H to optical disk size from \citet{Giovanelli83}. .
+ This gives σι~--2M.pe? and 2πα...ο.GAL.pedan?s 7., This gives $\sigma_{\rm H}\approx2~\rm{M_{\odot}~pc^{-2}}$ and $P_{\rm rest} \approx 6~\rm{M_{\odot}~pc^{-3}~km^2~s^{-2}}$ .
+ A lavee initial gas fraction is consistent with IC 3118s blue optical-IT color and low surface brightuess., A large initial gas fraction is consistent with IC 3418's blue optical-H color and low surface brightness.
+ We also considered the contribution of a Navarro-White (NFW) dark matter halo το P: rootDM=Oy(Ooxgwον).," We also considered the contribution of a Navarro-Frenk-White (NFW) dark matter halo to $P_{\rm rest}$; $P_{\rm rest,DM}=\sigma_{\rm H}\left(\partial\phi_{\rm NFW}/\partial r\right)$."
+" We found that the dark halo makes a significant contribution to the restoring pressure. but that Pa,»P4 uuless the halo mass exceccds zoQUPDALL."," We found that the dark halo makes a significant contribution to the restoring pressure, but that $P_{\rm ram}>P_{\rm rest}$ unless the halo mass exceeds $\approx10^{13}~{\rm M}_{\odot}$."
+ This is uurealistieallv high. iaud IC 3118s IL disk was extremely susceptible to RPS.," This is unrealistically high, and IC 3418's H disk was extremely susceptible to RPS."
+ For the case of molecular cloud RPS. we assuue spherically svuunetrie clouds with conservative values for a molecular cloud size aud deusitv of rc:=50 pe and pyc=107cm7.," For the case of molecular cloud RPS, we assume spherically symmetric clouds with conservative values for a molecular cloud size and density of $r_{\rm MC} = 50$ pc and $\rho_{\rm MC} = 10^3~{\rm cm}^{-3}$."
+" This eives a surface density. TH,προνο. of =2000AL.pe3,"," This gives a surface density, $\sigma_{\rm H_2}=(4/3)\rho_{\rm MC}r_{\rm MC}$, of $\approx 2000~\rm{M_{\odot}~pc^{-2}}$."
+" The restoring force from the stellar disk is P220.0], 7.", The restoring force from the stellar disk is $P_{\rm rest}=2\pi G\sigma_{\ast}\sigma_{\rm H_2}\approx6000~\rm{M_{\odot}~ pc^{-3}~ km^2~s^{-2}}$ .
+ For molecular clouds. Dau ," For molecular clouds, $P_{\rm rest}\gg P_{\rm ram}$ ."
+If IC 3118 is on a purely radial orbit aud ou its first pass through the cluster potential. then its tidal radius can be estimated by setting and assuming NEW poteutials.," If IC 3418 is on a purely radial orbit and on its first pass through the cluster potential, then its tidal radius can be estimated by setting and assuming NFW potentials."
+" Were in(e) denotes the mass enclosed within r. ry the galaxys tidal radius. aud r, the orbital radius within Vireo."," Here $m(r)$ denotes the mass enclosed within $r$, $r_t$ the galaxy's tidal radius, and $r_o$ the orbital radius within Virgo."
+ This gives au upper ut on rg: amore circular orbit will have a smaller tidal radius at IC 3118s cureut position., This gives an upper limit on $r_t$; a more circular orbit will have a smaller tidal radius at IC 3418's current position.
+ A lower limit cau be estimated by using (Maton2000).. which assumes that the galaxy is at pericenter and that it has an orbital eccentricity typical of subhalos ini N-bodysimulations.," A lower limit can be estimated by using \citep{Mamon00}, which assumes that the galaxy is at pericenter and that it has an orbital eccentricity typical of subhalos in -bodysimulations."
+ Hereνου dsVirgos virial radius.," Here$r_{v,{\rm Virgo}}$ isVirgo's virial radius."
+For both estimates we assumea dark matter to barvon mass ratio of papiMaryzm15(Mp210197 ALL).,For both estimates we assumea dark matter to baryon mass ratio of $M_{\rm DM}/M_{\rm bary}\approx15$$M_{\rm DM}\approx10^{10.3}~{\rm M}_{\odot}$ ).
+" We adopt revive, from Mebaugliliu (1999).."," We adopt $r_{v,\rm{Virgo}}$ from \citet{Mclaughlin99}. ."
+ IC 3118 currenttidal racius is between 9 aud 30 Ipc., IC 3418's currenttidal radius is between 9 and 30 kpc.
+ If the ealaxys is not at poricenter. then its final tidal," If the galaxy is not at pericenter, then its final tidal"
+where g. is (he vertical gravity.,where $g_z$ is the vertical gravity.
+" When the Ri drops below a critical value. πέος, (vpically zz|/4. we have Kelvin-Helmholtz instability. (XID) and the onset of midplane turbulence."," When the $Ri$ drops below a critical value, $Ri_{\rm c}$, typically $\approx 1/4$, we have Kelvin-Helmholtz instability (KHI) and the onset of midplane turbulence."
+ At low particulate densities py<pe. we may approximate the vertical gravity as coming solely from the central star. g.&O72. while differentiation of equation (9)) vields for equation (10)): Inereasing stratification is then destabilizing. since the increase in shear outweighs the increasing huovaney. decreasing /27 until it reaches A.," At low particulate densities $\rho_{\rm p}\ll \rho_{\rm g}$, we may approximate the vertical gravity as coming solely from the central star, $g_z \approx \Omega^2 z$, while differentiation of equation \ref{azimuthal}) ) yields for equation \ref{rich}) ): Increasing stratification is then destabilizing, since the increase in shear outweighs the increasing buoyancy, decreasing $Ri$ until it reaches $Ri_{\rm c}$."
+ In our detailed calculations. we include the sell-gravitv of the combined [ια in the plane-parallel approximation: When py becomes large because of the settling of particulates to the midplane. the of sub-disk can provide a significant stabilizing inflivence against INKILI.," In our detailed calculations, we include the self-gravity of the combined fluid in the plane-parallel approximation: When $\rho_{\rm p}$ becomes large because of the settling of particulates to the midplane, the self-gravity of sub-disk can provide a significant stabilizing influence against KHI."
+ Indeed. we shall lind (hat this additional elfect is the cause of a sudden “precipitation phenomenon” when the surface densitv of solids crosses a critical value.," Indeed, we shall find that this additional effect is the cause of a sudden “precipitation phenomenon” when the surface density of solids crosses a critical value."
+ Since the sell-eravity of the subdisk of solids is also of crucial importance for the Goldreich-Ward mechanism. we find (hat “precipitation” is virtually svnonvinous with (he satisfaction of the criterion for GI. equation (2)).," Since the self-gravity of the subdisk of solids is also of crucial importance for the Goldreich-Ward mechanism, we find that “precipitation” is virtually synonymous with the satisfaction of the criterion for GI, equation \ref{Roche}) )."
+ To see this. note that for the second term on the right-hand side of equation (13)) to become of comparable importance as the first. we require which differs [rom equation (2)) only by a factor of 7.8.," To see this, note that for the second term on the right-hand side of equation \ref{grav}) ) to become of comparable importance as the first, we require which differs from equation \ref{Roche}) ) only by a factor of 7.8."
+ Precipitation is formally reached belore GI Gl we think of a process of slowly increasing the local solid/gas ratio in the disk). but since precipitation resul(s. as we shall see. in the settling out of a very dense of solids. GI follows almost immediately alter conditions in the disk become ripe for the beginning of precipitation.," Precipitation is formally reached before GI (if we think of a process of slowly increasing the local solid/gas ratio in the disk), but since precipitation results, as we shall see, in the settling out of a very dense midplane-layer of solids, GI follows almost immediately after conditions in the disk become ripe for the beginning of precipitation."
+ Belore these dramatic events occur. the particulate laver should reach an equilibrium which is marginally unstable to the KIL. because even weak turbulence is sufficient to halt the settling of well-coupled particles.," Before these dramatic events occur, the particulate layer should reach an equilibrium which is marginally unstable to the KHI, because even weak turbulence is sufficient to halt the settling of well-coupled particles."
+" We hencelorth assume with Sekiva that the subclisk of solids satisfies at each value of z. the marginally stable condition. ἐπ)=Ri,1/4. consistent with (he necessary criterion for instability in plane-parallel flows."," We henceforth assume with Sekiya that the subdisk of solids satisfies at each value of $z$, the marginally stable condition, $Ri = Ri_{\rm c} =
+1/4$, consistent with the necessary criterion for instability in plane-parallel flows."
+ The result is probably mocdified by the disk geometry and Coriolis forces., The result is probably modified by the disk geometry and Coriolis forces.
+ ILowever. preliminary investigations of an," However, preliminary investigations of an"
+"1) For models with g=1, the simulations indicate that the system evolves toward a quasi equilibrium state with both planets trapped in a mean motion resonance and evolving on non migrating orbits.","i) For models with $q=1$, the simulations indicate that the system evolves toward a quasi equilibrium state with both planets trapped in a mean motion resonance and evolving on non migrating orbits."
+" This occurs because the positive corotation torques exerted on the innermost planet counterbalance the negative corotation torques exerted on the outermost ii) For models with q>1, most of the runs resulted in a similar mode of evolution."," This occurs because the positive corotation torques exerted on the innermost planet counterbalance the negative corotation torques exerted on the outermost ii) For models with $q \ge 1$, most of the runs resulted in a similar mode of evolution."
+" Occasionally however, the final fate of the sytem was such that the two bodies are in close vicinity to one another, have stopped migrating, but are not in resonance."," Occasionally however, the final fate of the sytem was such that the two bodies are in close vicinity to one another, have stopped migrating, but are not in resonance."
+" This arises when the mass of the innermost planet is high enough to open a partial gap in the disc, such that the migration of the outer planet is stopped at the edge of the large cavity formed by both the binary and the innermost iii) For most of the models with 4«1, the planets involved in the simulations underwent different dynamical processes such as scattering or orbital exchange."," This arises when the mass of the innermost planet is high enough to open a partial gap in the disc, such that the migration of the outer planet is stopped at the edge of the large cavity formed by both the binary and the innermost iii) For most of the models with $q < 1$, the planets involved in the simulations underwent different dynamical processes such as scattering or orbital exchange."
+" When orbital exchange occurs, the final state of the system is a stable mean motion resonance with the more massive planet now being the innermost one."," When orbital exchange occurs, the final state of the system is a stable mean motion resonance with the more massive planet now being the innermost one."
+ Scattering and orbital exchange occurs once a first order resonance is established between the planets., Scattering and orbital exchange occurs once a first order resonance is established between the planets.
+ This drives up the eccentricity of the inner planet leading to a close encounter with either the outer planet or the binary., This drives up the eccentricity of the inner planet leading to a close encounter with either the outer planet or the binary.
+ A close encounter with the binary leads to ejection., A close encounter with the binary leads to ejection.
+" Having examined the two-planet problem, we then focused on more ""realistic"" systems composed of five planets with masses of 5, 7.5, 10, 12.5 and 15 Mg embedded in a circumbinary disc."," Having examined the two-planet problem, we then focused on more “realistic” systems composed of five planets with masses of 5, 7.5, 10, 12.5 and 15 $\mearth$ embedded in a circumbinary disc."
+" We performed three simulations, with the planets being placed in different initial orbital configurations."," We performed three simulations, with the planets being placed in different initial orbital configurations."
+" In general terms, the evolution of such a system proceeds as follows."," In general terms, the evolution of such a system proceeds as follows."
+ Each planet migrates inward until it is captured into resonance with an adjacent body., Each planet migrates inward until it is captured into resonance with an adjacent body.
+ This occurs either, This occurs either
+aabsorbers are weaker versions of strong aabsorbers with similar kinematics. as if they arise in the outskirts or in sparse regions of luminous galaxies.,"absorbers are weaker versions of strong absorbers with similar kinematics, as if they arise in the outskirts or in sparse regions of luminous galaxies."
+ Others. which tend to have more kinematically compact profiles. have been hvpothesized to arise in dwarf galaxies al.2005).," Others, which tend to have more kinematically compact profiles, have been hypothesized to arise in dwarf galaxies \citep{Zonak04,Ding05,Mas05}."
+ When exploring the nature of weak aabsorbers. (wo methods are often usedstatistical survevs and photoionization modeling ol individual svstems. Naravananetal. (2005)... Churchille," When exploring the nature of weak absorbers, two methods are often used–statistical surveys and photoionization modeling of individual systems. \citet{Nar05}, , \citet{Church99b},"
+tal.(1999b).. and (2006) conducted surveys of svstems with equivalent widths in the range 0.02<VV.0.3A. which when combined. span the redshift range 0<22.4.," and \citet{Lynch06} conducted surveys of systems with equivalent widths in the range $0.02 \le W_r \le 0.3~\Ang$, which when combined, span the redshift range $0 < z < 2.4$."
+ These survevsobtained number densities of absorbers per unit redshift (d.N/dz): dN/d2=1.0020.20(0<z«0.3) (Naravananetal.2005).. dN/dz=1.04dE0.10(0.4<2«1.4) (Churchilletal.1999b).. and ανα:=1.0630.12(1.4<z24) (Lynchetal.2006).," These surveysobtained number densities of absorbers per unit redshift $dN/dz$ ): $dN/dz = 1.00 \pm 0.20~(0 < z < 0.3)$ \citep{Nar05}, $dN/dz = 1.74 \pm 0.10~(0.4 < z < 1.4)$ \citep{Church99b}, and $dN/dz = 1.06 \pm 0.12~(1.4 < z < 2.4)$ \citep{Lynch06}."
+. Apparently. the population of weak aabsorbers peaks al z71 over the range 0<z«2.4.," Apparently, the population of weak absorbers peaks at $z \sim 1$ over the range $0 < z < 2.4$."
+ Furthermore. the density of absorbers al zc2 is sienilicantlv lower than expected if the change in &N/dz was due onlv {ο cosmological effects. and to the changing extragalactic background radiation 2006).," Furthermore, the density of absorbers at $z \sim 2$ is significantly lower than expected if the change in $dN/dz$ was due only to cosmological effects and to the changing extragalactic background radiation \citep{Lynch06}."
+. This indicates that either (he same process creates weak aabsorbers across (his redshift range but was less active al z~2 than ad z~1. or that the physical mechanisms responsiblefor creating the absorbers change across reclshilt aud are more ellicient al z~1 than al z~2.," This indicates that either the same process creates weak absorbers across this redshift range but was less active at $z \sim 2$ than at $z \sim 1$, or that the physical mechanisms responsiblefor creating the absorbers change across redshift and are more efficient at $z \sim 1$ than at $z \sim 2$."
+ It should be noted (that (his trend in αλας is consistent with the star formation history in dwar! galaxies etal. 2004)... which may suggest that weak aabsorbers are related to this activity.," It should be noted that this trend in $dN/dz$ is consistent with the star formation history in dwarf galaxies \citep{Gab04,Kauffmann04}, which may suggest that weak absorbers are related to this activity."
+ Photoionization models facilitate the exploration of the physical properties of the absorption svslenm. such as ionization parameter. density. temperature. abundance pattern. ancl size.," Photoionization models facilitate the exploration of the physical properties of the absorption system, such as ionization parameter, density, temperature, abundance pattern, and size."
+ By understanding these properties for absorbers at different redshifts. we can gain insight into what processes create them.," By understanding these properties for absorbers at different redshifts, we can gain insight into what processes create them."
+ If different processes are responsible for the formation of these absorbers at different. epochs. then the plvsical properties of the absorbers are likely to evolve.," If different processes are responsible for the formation of these absorbers at different epochs, then the physical properties of the absorbers are likely to evolve."
+ If no statistical difference in the plivsical properties is observed. (this would suggest that the same mechanism is responsible lor creating weak aabsorbers at dillerent epochs.," If no statistical difference in the physical properties is observed, this would suggest that the same mechanism is responsible for creating weak absorbers at different epochs."
+ To this end. we have produced photoionization models of the nine 1.4«&z«2.4 svstems found by Lynchetal.(2006) using the code Cloudy. οιal.1998). and compared our results to those of svstems al lower redshift (σονetal. 2005)..," To this end, we have produced photoionization models of the nine $1.4 < z < 2.4$ systems found by \citet{Lynch06} using the code Cloudy \citep{Fer98} and compared our results to those of systems at lower redshift \citep{Rig02,Char03,Zonak04,Ding05,Mas05}. ."
+ Our methodology for modeling the systems is the subject o£ 2.., Our methodology for modeling the systems is the subject of \ref{sec:method}. .
+ In 3 we report (the results of our models for each individual svstem., In \ref{sec:props} we report the results of our models for each individual system.
+In +4 we discuss our results and their implications,In \ref{sec:disc} we discuss our results and their implications
+rrather than a high abundance ofPF.,rather than a high abundance of.
+ Finally. the observational data regarding SC stars should be updated using more recent atmospheric models.," Finally, the observational data regarding SC stars should be updated using more recent atmospheric models."
+ AL L. deeply appreciates (he hospitality and support extended to her by Michael Wiescher. Joachim Gorrres and. Claudio Ugalde during a visit to Notre Dame University. as well as the hospitality received by John Lattanzio and Amanca Ixarakas during a visit to Monash University.," M. L. deeply appreciates the hospitality and support extended to her by Michael Wiescher, Joachim Görrres and Claudio Ugalde during a visit to Notre Dame University, as well as the hospitality received by John Lattanzio and Amanda Karakas during a visit to Monash University."
+ We thank Roberto Gallino. Enrico Arnone. Stefano Masera and Richard Stanclilfe for discussion aud help.," We thank Roberto Gallino, Enrico Arnone, Stefano Masera and Richard Stancliffe for discussion and help."
+ The manuscript was much extended and improved following strong criticisms by the anonvimous referee., The manuscript was much extended and improved following strong criticisms by the anonymous referee.
+ This work was supported by the Australian Research Council and the Australian. Partnership for Advances Computing., This work was supported by the Australian Research Council and the Australian Partnership for Advances Computing.
+ Computational resources used for this study were also partly funded by the Canada Foundation for Innovation (CET) and the Nova Seolia Research and Innovation Trust fund. (NSRIT)., Computational resources used for this study were also partly funded by the Canada Foundation for Innovation (CFI) and the Nova Scotia Research and Innovation Trust fund (NSRIT).
+Visible spectra. of TNOs 1995 SMss. 2003 ΟΡ... 2003 211). 2004 SBoo and 2005 UQs;;. were done with the 3.58m Telescopio Nazionale Galileo (TNG. El Roque de los Muchachos Observatory. Canary Islands. Spain) on three consecutive nights from September 15 to 17. 2007.,"Visible spectra of TNOs 1995 $_{55}$, 2003 $_{32}$ , 2003 $_{117}$, 2004 $_{60}$ and 2005 $_{513}$, were done with the 3.58m Telescopio Nazionale Galileo (TNG, El Roque de los Muchachos Observatory, Canary Islands, Spain) on three consecutive nights from September 15 to 17, 2007."
+ The spectrograph DOLORES with the LR-R grism and the 2.07 slit width was used.," The spectrograph DOLORES with the LR-R grism and the 2.0"" slit width was used."
+ Several spectra of each object with expossure times from 1600 to 2400 sec. covering the 0.52 to 0.86pm spectral range. were obtained (see Table 1)).," Several spectra of each object with expossure times from 1600 to 2400 sec, covering the 0.52 to $\mu $ m spectral range, were obtained (see Table \ref{Table1}) )."
+ Each object was shifted in the slit by 57 between consecutive spectra to better correct the fringing., Each object was shifted in the slit by 5” between consecutive spectra to better correct the fringing.
+ Images were bias corrected using the over-scan region and flat-field corrected using lamp flats., Images were bias corrected using the over-scan region and flat-field corrected using lamp flats.
+ The two-dimensional spectra were extracted. sky background subtracted. and collapsed to one dimension.," The two-dimensional spectra were extracted, sky background subtracted, and collapsed to one dimension."
+ The wavelength calibration was done using the Neon and Argon lamps., The wavelength calibration was done using the Neon and Argon lamps.
+ All the spectra of each TNO. obtained at different positions 1n the slit. were averaged and rebined to obtain a higher S/N. To correct for tellurie absorption and to obtain the relative reflectance. the G stars Landolt (SA) 110-36. Landolt (SA) 115-271. Landolt (SA) 93-101. Landolt (SA) 112-13. Landolt (SA) 98-978. Landolt (SA) 110-113 and Landolt (SA)115-27 (Landolt. 1992)) were observed at different airmasses (similar to those of the TNOs) during each night. before and after the TNO observations. and used as solar analogue stars.," All the spectra of each TNO, obtained at different positions in the slit, were averaged and rebined to obtain a higher S/N. To correct for telluric absorption and to obtain the relative reflectance, the G stars Landolt (SA) 110-36, Landolt (SA) 115-271, Landolt (SA) 93-101, Landolt (SA) 112-13, Landolt (SA) 98-978, Landolt (SA) 110-113 and Landolt (SA)115-27 (Landolt, \cite{Landolt}) ) were observed at different airmasses (similar to those of the TNOs) during each night, before and after the TNO observations, and used as solar analogue stars."
+ The spectrum of each object was divided by those of the solar analogue stars observed the same night and at similar airmasses. and then normalized to unity around 0.554 thus obtaining the normalized reflectance.," The spectrum of each object was divided by those of the solar analogue stars observed the same night and at similar airmasses, and then normalized to unity around $\mu $ m thus obtaining the normalized reflectance."
+ The obtained spectra are shown in Fig. |., The obtained spectra are shown in Fig. \ref{Fig1}.
+ The spectra of the five TNOs are featureless within. the uncertainties., The spectra of the five TNOs are featureless within the uncertainties.
+ No signatures of any absorption band centered at 0.7 um. typically observed in low albedo main belt asteroids and attributed to silicate aqueous alteration. are found.," No signatures of any absorption band centered at 0.7 $\mu$ m, typically observed in low albedo main belt asteroids and attributed to silicate aqueous alteration, are found."
+ This very weak absorption has been reported for other TNOs (e.g Lazzarin et al. 2004:;, This very weak absorption has been reported for other TNOs (e.g Lazzarin et al. \cite{LazzPhyll};
+ Fornasier et al. 2004))., Fornasier et al. \cite{FornaPhyll}) ).
+ However. the spectra show variety in color. from bluish to red.," However, the spectra show variety in color, from bluish to red."
+ To perform a quantitative analysis of the color distribution of the ELa;-group and to diserimminate. if the candidates studied belong to it. we computed the spectral gradient ΤΠ | as detined by Jewitt. (2002):," To perform a quantitative analysis of the color distribution of the $_{61}$ -group and to discriminate if the candidates studied belong to it, we computed the spectral gradient $S'[\% (0.1\mu$ $^{-1}]$ as defined by Jewitt. \cite{Jewittslope}) ):"
+ where $ is the relative reflectance of the object and S is the mean value of this reflectance in the wavelength range over which 4Sdl is computed., where $S$ is the relative reflectance of the object and $\overline{S}$ is the mean value of this reflectance in the wavelength range over which ${dS}$ ${d\lambda}$ is computed.
+ S is à measurement of the variation of SCU over an interval Al., $S'$ is a measurement of the variation of $S(\lambda)$ over an interval $\Delta\lambda$.
+ Assuming that a featureless spectrum has an overall linear shape (as a first order approach). we made a linear fitting of the visible spectra normalized at 0.55 um. obtaining $.," Assuming that a featureless spectrum has an overall linear shape (as a first order approach), we made a linear fitting of the visible spectra normalized at 0.55 $\mu$ m, obtaining $S'$."
+ Results are shown in Table 2.. together with the dynamical parameters of each TNO.," Results are shown in Table \ref{Table2}, , together with the dynamical parameters of each TNO."
+" The analysis of the solar analogue stars observed the three nights demonstrate that uncertainties up to 2656/0.μι in S"" are possible due to systematic errors.", The analysis of the solar analogue stars observed the three nights demonstrate that uncertainties up to $/0.1\mu$ m in $S'$ are possible due to systematic errors.
+ This is smaller than the uncertainties due to spectrophotometrie determinations that can easily be larger than 5956/0.1jum. TNOs 1995 ος. 2003 OP: are already considered as members of the family.," This is smaller than the uncertainties due to spectrophotometric determinations that can easily be larger than $/0.1\mu$ m. TNOs 1995 $_{55}$, 2003 $_{32}$ are already considered as members of the family."
+ Spectral gradients computed from their spectra (see table 2.), Spectral gradients computed from their spectra (see table 2.)
+ are consistent with Ragozzine Brown (2007)) values obtained from photometry 1.79 + 2.60 and -1.09 + 2.20 96/0. 1m respectively., are consistent with Ragozzine Brown \cite{Ragozzine}) ) values obtained from photometry 1.79 $\pm$ 2.60 and -1.09 $\pm$ 2.20 $/0.1\mu$ m respectively.
+ The spectral gradients computed from the spectra of the other members of the group 2003 ΕΙ. 2002 ΤΝ. and 2005 RR: CPinilla-Alonso et al. 2008..," The spectral gradients computed from the spectra of the other members of the group 2003 $_{61}$, 2002 $_{300}$, and 2005 $_{43}$ (Pinilla-Alonso et al. \cite{Pinilla-AlonsoEL61},"
+ Licandro et al. 2006..," Licandro et al. \cite{LicTX},"
+ Pinilla-Alonso et al.. 2007))," Pinilla-Alonso et al., \cite{Pinilla-AlonsoRR43}) )"
+ are 0.0. 1.0 and 0.4 respectively.," are 0.0, 1.0 and 0.4 respectively."
+" These give a mean S'z0.4 + 1.090/0. Lum for the ELG,-group.", These give a mean $S'$ =0.4 $\pm$ $/0.1\mu$ m for the $_{61}$ -group.
+ The visible spectra of the three observed TNOs candidates to belong to the ELa;-group according to Ragozzine Brown (2007)). are all featureless but show variety in colors.," The visible spectra of the three observed TNOs candidates to belong to the $_{61}$ -group according to Ragozzine Brown \cite{Ragozzine}) ), are all featureless but show variety in colors."
+" Two of them (2004 SBeu and 2005 UQz,:) are too red ($= 12.6 and 18.1 + 2.0 9c/0. Lum respectively) to be members of the group.", Two of them (2004 $_{60}$ and 2005 $_{513}$ ) are too red $S'=$ 12.6 and 18.1 $\pm$ 2.0 $/0.1\mu$ m respectively) to be members of the group.
+ The surfaces of these objects probably contain a significant fraction of processed materials composed of complex organics., The surfaces of these objects probably contain a significant fraction of processed materials composed of complex organics.
+ On the other hand. 2003 UZj;5 should be considered as a strong candidate to be part of the EL¢)-group.," On the other hand, 2003 $_{117}$ should be considered as a strong candidate to be part of the $_{61}$ -group."
+ Its spectrum is featureless in the visible and bluish with an spectral gradient computed from the spectrum. S’=-2 + 2 %/O.1pm. close tothe mean value for the whole group and compatible with the spectral gradient computed from the photometry (Ragozzine etal. 2007).," Its spectrum is featureless in the visible and bluish with an spectral gradient computed from the spectrum, $S'$ =-2 $\pm$ 2 $/0.1\mu$ m, close tothe mean value for the whole group and compatible with the spectral gradient computed from the photometry (Ragozzine etal. \cite{Ragozzine}) )."
+ This fact is consistent. but not conclusive as we need near-infrared photometry and/or spectroscopy to confirm the presence of water ice.," This fact is consistent, but not conclusive as we need near-infrared photometry and/or spectroscopy to confirm the presence of water ice."
+ It is well known that solar irradiation during Gy., It is well known that solar irradiation during Gy.
+ sublimates volatiles from the surface of, sublimates volatiles from the surface of
+"galaxy group was found to lead to the formation of a single elliptical galaxy (hence a ""fossil. group”) in a few billion years (Barnes1989)5.",galaxy group was found to lead to the formation of a single elliptical galaxy (hence a “fossil group”) in a few billion years (Barnes.
+. The final product retains an X-ray emitting halo of hot gas (Ponman Bertram 1993)., The final product retains an X-ray emitting halo of hot gas (Ponman Bertram 1993).
+ The whole scenario was proved to be plausible by the discovery of the fossil group | 4018 (Ponman et al., The whole scenario was proved to be plausible by the discovery of the fossil group $+$ 4018 (Ponman et al.
+ 1994)., 1994).
+ The stanclare observational definition of a fossil svsteni is based on the existence of extended: X-ray emission with a bolometric luminosity Lx101haergs and of an R-band magnitude eap Amy.ον2 between the brightest and. second brightest. members within 0.572500 (Jones. Ponman Forbes 2000).," The standard observational definition of a fossil system is based on the existence of extended X-ray emission with a bolometric luminosity $L_{\mathrm{X}} > 10^{42}~h_{50}^{-2}~\mathrm{erg~s}^{-1}$ and of an R-band magnitude gap $\mathrm{\Delta} m_{12} > 2$ between the brightest and second brightest members within $0.5 R_{200}$ (Jones, Ponman Forbes 2000)."
+ In theoretical studies. the analogous definition of a fossil svstem is less straightforward (sce c.e.. D'Onghia et al.," In theoretical studies, the analogous definition of a fossil system is less straightforward (see e.g., D'Onghia et al."
+ 2005 for a cliscussion)., 2005 for a discussion).
+ On the basis of the existing observations. mostly limited to the local Universe. fossil groups are expected. to be a considerable population among the gravitationally bound systems of galaxies (Vikhlinin et al.," On the basis of the existing observations, mostly limited to the local Universe, fossil groups are expected to be a considerable population among the gravitationally bound systems of galaxies (Vikhlinin et al."
+ L999: Romer ct al., 1999; Romer et al.
+ 2000: Jones et al., 2000; Jones et al.
+ 2003: La Barbera et al., 2003; La Barbera et al.
+ 2009: Voevodkin et al., 2009; Voevodkin et al.
+ 2010)., 2010).
+ For instance. their estimated spatial density is equal to 2.8310μάςMpe for an X-ray bolometric luminosity above O.S91072hoo7Cres (La Barbera οἱ al.," For instance, their estimated spatial density is equal to $2.83 \times 10^{-6}~h_{75}^{3}~\mathrm{Mpc}^{-3}$ for an X-ray bolometric luminosity above $0.89 \times 10^{42}~h_{75}^{-2}~\mathrm{erg~s}^{-1}$ (La Barbera et al."
+ 2009)., 2009).
+ X-rav. properties and dark-matter (DM) content of fossil. eroups are comparable to those of bright groups ancl poor clusters of galaxies. with total masses of ~1077 103NL;. whereas their brightest members are typically as luminous as a cluster ¢D galaxy (Ay</—22.5|5log hz).," X-ray properties and dark-matter (DM) content of fossil groups are comparable to those of bright groups and poor clusters of galaxies, with total masses of $\sim 10^{13}$ $10^{14}~\mathrm{M}_{\sun}$, whereas their brightest members are typically as luminous as a cluster cD galaxy $M_{\mathrm{R}} < -22.5 + 5~\mathrm{log}~h_{50}$ )."
+ In recent. high-resolution hydrodynamical cosmological simulations in a A cold. dark. matter CACDAL) universe. fossil groups have already assembled. half of their. final mass at redshifts zz1. and subsequently they typically. erow bv minor mergers only. whereas non-fossil svstems of similar total mass on average form later (D'Onghia et al.," In recent high-resolution hydrodynamical cosmological simulations in a $\mathrm{\Lambda}$ cold dark matter $\mathrm{\Lambda}$ CDM) universe, fossil groups have already assembled half of their final mass at redshifts $z \ge 1$, and subsequently they typically grow by minor mergers only, whereas non-fossil systems of similar total mass on average form later (D'Onghia et al."
+ 2005: Dariush et. al., 2005; Dariush et al.
+ 2007)., 2007).
+ Phe carly assembly leaves sullicicnt time for galaxies with a luminosity LL' (assuming a Sehechter 1976 luminosity function) to merge into the central galaxy. by dynamical friction. which vields the magnitude gap defining fossil groups.," The early assembly leaves sufficient time for galaxies with a luminosity $L \ge L^{\star}$ (assuming a Schechter 1976 luminosity function) to merge into the central galaxy by dynamical friction, which yields the magnitude gap defining fossil groups."
+" This ""early formation"" scenario is alternative to the compact group mereine scenario (Barnes 1989).", This “early formation” scenario is alternative to the compact group merging scenario (Barnes 1989).
+ Llowever. no dillerence seems to exist in the distribution of neighboring faint galaxy density excess. distance from the redsequence. structural parameters. or stellar population properties between the brightest central galaxies (BCGs) in 25 fossil groups at. z0.1 and a sample of bright field ellipticals with a similar distribution in optical Iuminosity and X-ray properties. (La Barbera ct al.," However, no difference seems to exist in the distribution of neighboring faint galaxy density excess, distance from the red, structural parameters, or stellar population properties between the brightest central galaxies (BCGs) in 25 fossil groups at $z \le 0.1$ and a sample of bright field ellipticals with a similar distribution in optical luminosity and X-ray properties (La Barbera et al."
+ 2009)., 2009).
+" “hese authors argue that this evidence undermines the idea that fossil svstems represent the send product"" of a physical mechanism that occurred at high redshift: it rather suggests that fossil groups are the final stage of mass assembly in the Universe.", These authors argue that this evidence undermines the idea that fossil systems represent the “end product” of a physical mechanism that occurred at high redshift; it rather suggests that fossil groups are the final stage of mass assembly in the Universe.
+ Whether either interpretation is true or fossil groups correspond to a transient phase in the life of a galaxy group (von Benda-Deckmann et al., Whether either interpretation is true or fossil groups correspond to a transient phase in the life of a galaxy group (von Benda-Beckmann et al.
+ 2008: Dariush et al., 2008; Dariush et al.
+ 2010: Cui et al., 2010; Cui et al.
+ 2011). whether they form. a clistinet family of groups (D'Onshia Lake 2004: Whosroshahi. Ponman Jones 2007) or not (Sales et al.," 2011), whether they form a distinct family of groups (D'Onghia Lake 2004; Khosroshahi, Ponman Jones 2007) or not (Sales et al."
+ 2007: Zibetti. Pierini Pratt 2009: Sun et al.," 2007; Zibetti, Pierini Pratt 2009; Sun et al."
+ 2009; Démmocloss ct al., 2009; Démmoclèss et al.
+ 2010: Voevodkin ct al., 2010; Voevodkin et al.
+ 2010: Cui et al., 2010; Cui et al.
+ 2011). for several reasons it is important to extend the present. knowledge of fossil systems bevond the local Universe and towards lower total masses.," 2011), for several reasons it is important to extend the present knowledge of fossil systems beyond the local Universe and towards lower total masses."
+ Lf unclisturbed for a very lone time. fossil groups may represent the ultimate example of hydrostatic equilibrium in viriliazed svstems (Vikhlinin et al.," If undisturbed for a very long time, fossil groups may represent the ultimate example of hydrostatic equilibrium in viriliazed systems (Vikhlinin et al."
+ 1999)., 1999).
+ This conclusion is supported by a recent detailed study of two such systems based. on high-quality spectro-imaging data (Dénmmoclóss οἱ al., This conclusion is supported by a recent detailed study of two such systems based on high-quality spectro-imaging data (Démmoclèss et al.
+ 2010)., 2010).
+ In addition. fossil groups could. represent the. least massive. barvonicallv closed svstems (Mathews et al.," In addition, fossil groups could represent the least massive, baryonically closed systems (Mathews et al."
+ 2005). their barvon mass fraction being equal to the cosmic meanο," 2005), their baryon mass fraction being equal to the cosmic mean."
+"τι, Lf so. they could. formi a reference. sample with a recent. quiescent merging activity with which one can investigate the origin of the average barvon deficiency within {σου (%OF fous) Found in X-ray emitting eroups at 1. (Gioclini et al."," If so, they could form a reference sample with a recent, quiescent merging activity with which one can investigate the origin of the average baryon deficiency within $R_{500}$ $\approx 0.7 R_{200}$ ) found in X-ray emitting groups at $z \le 1$ (Giodini et al."
+ 2009. hereafter. referred to as (100: see also Tevssier et al.," 2009, hereafter referred to as G09; see also Teyssier et al."
+ 2011)., 2011).
+" In. particular. two pair-matched samples of fossil groups with a range in total mass ancl redshift. one with racdio-Ioud BOCs and the other with radio-«quiet BCGs. could be used to estimate the ""instantaneous? boost of the entropy in the central regions of the ICM that is due to radio-galaxy feedback (c£."," In particular, two pair-matched samples of fossil groups with a range in total mass and redshift, one with radio-loud BCGs and the other with radio-quiet BCGs, could be used to estimate the “instantaneous” boost of the entropy in the central regions of the ICM that is due to radio-galaxy feedback (cf."
+ Giodini et al., Giodini et al.
+ 2010)., 2010).
+ From simple considerations on the hierachical nature of structure growth via merging (Milosavljevié ct al., From simple considerations on the hierachical nature of structure growth via merging (Milosavljević et al.
+ 2006). fossil systems are expected to be more. frequent (hy a [actor of ten) at the group mass-scale than at the eluster mass-scale (~lot 101”AL. ο.," 2006), fossil systems are expected to be more frequent (by a factor of ten) at the group mass-scale than at the cluster mass-scale $\sim 10^{14}$ $10^{15}~\mathrm{M}_{\sun}$ )."
+ This is confirmed. by the most recent study of fossil groups based on hyclroclynamical resimulations of five regions in the Millennium Simulation and on galaxy formation models (Cui et al., This is confirmed by the most recent study of fossil groups based on hydrodynamical resimulations of five regions in the Millennium Simulation and on galaxy formation models (Cui et al.
+ 2011 and references therein)., 2011 and references therein).
+ “Pherelore fossil groups may also help in understanding the properties. of more massive bound systems ab 2=0. especially if it turns out that they foster the formation of the most massive progenitor of the BCG of a cluster. as suggested by 00454 0309 at 2=0.26 (Schirmer et al.," Therefore fossil groups may also help in understanding the properties of more massive bound systems at $z = 0$, especially if it turns out that they foster the formation of the most massive progenitor of the BCG of a cluster, as suggested by $-$ 0309 at $z = 0.26$ (Schirmer et al."
+ 2010: οἱ., 2010; cf.
+ Barnes 1989. and references therein)., Barnes 1989 and references therein).
+ During our ongoing studies of the X-ray selected groups in the 2deg? area of the Cosmic Evolution SurveyALOS: Seoville et al., During our ongoing studies of the X-ray selected groups in the $2~\mathrm{deg}^2$ area of the Cosmic Evolution Survey; Scoville et al.
+ 2007). we have discovered two fossil eroups ats=0.372 and z=0.425.," 2007), we have discovered two fossil groups at $z = 0.372$ and $z = 0.425$."
+ As demonstrated. by G09 and Ciodini et al. (, As demonstrated by G09 and Giodini et al. (
+2010). the multi-wavelength data set of the ollers a thus far unique opportunity,"2010), the multi-wavelength data set of the offers a thus far unique opportunity"
+60e V). that do not ionize the elements effectively. heat the medium. ahead. of the ionization front.,"$> 60 \, \rm eV$ ), that do not ionize the elements effectively, heat the medium ahead of the ionization front."
+ 1 the source had a harder spectrum. e.g. 107IX black body. than the gas would be heated even at larger radii.," If the source had a harder spectrum, e.g. $10^5 \, \rm K$ black body, than the gas would be heated even at larger radii."
+ We stress that at a distance of =Skpe (namely. around the Strómmgren radius) teniperaturessteeply drop from 10!I& down to ~107Ix and recombination processes take place (see next)," We stress that at a distance of $\gtrsim 5\,\rm kpc$ (namely, around the Strömmgren radius) temperaturessteeply drop from $\sim 10^4\,\rm K$ down to $\sim 10^3\,\rm K$ and recombination processes take place (see next)."
+ Correspondingly to the temperature profile. in Fig. 3..," Correspondingly to the temperature profile, in Fig. \ref{fig:SST_um},"
+ we displav the radial profiles of the dilferent chemical abundances at 500Myr after the radiative source has been switched on.," we display the radial profiles of the different chemical abundances at $500\,\rm Myr$ after the radiative source has been switched on."
+ We assumed initial cosmicabundances?.. which means that hydrogen species account for ~93% of the total number densities and. helium species for ~7%.," We assumed initial cosmic, which means that hydrogen species account for $\sim 93\% $ of the total number densities and helium species for $\sim 7\%$."
+ As expected. ionized fractions usually have larger values closer to the sources (within a few kpe). while neutral or molecular fractions increase at larger distances (above 3Gkpc).," As expected, ionized fractions usually have larger values closer to the sources (within a few kpc), while neutral or molecular fractions increase at larger distances (above $3-6\,\rm
+kpc$ )."
+" In the following we discuss these trends. by referring to the reaction network of Table 2.. in a more precise and cdetailed way,"," In the following we discuss these trends, by referring to the reaction network of Table \ref{tab:reactions}, , in a more precise and detailed way."
+Another instrument of similar capabilities is the Ohio State Infrared. Lager/Spectrometer (OSIRIS). currently installed at the Southern Astrophysics Research Observatory (SOAR). attached to the Lin telescope.,"Another instrument of similar capabilities is the Ohio State Infrared Imager/Spectrometer (OSIRIS), currently installed at the Southern Astrophysics Research Observatory (SOAR), attached to the 4.1m telescope."
+ OSIRIS provides spectral coverage from 1.01un to [tun in cross-cdispersed moce. with a resolving power of ~ 1200.," OSIRIS provides spectral coverage from $\upmu$ m to $\upmu$ m in cross-dispersed mode, with a resolving power of $\sim$ 1200."
+ High resolution (Il ~ 3000) long-slit 1uodes are also available. but inulti-baud spectroscopy of this kind sullers Grom differences in aperture and seeiug.," High resolution (R $\sim$ 3000) long-slit modes are also available, but multi-band spectroscopy of this kind suffers from differences in aperture and seeing."
+ However. reductiou of NIR spectra has a complexity of its own. mostly related to telluric spectral features. both in absorption aud emission. and black body raciation due to the telescope itself.," However, reduction of NIR spectra has a complexity of its own, mostly related to telluric spectral features, both in absorption and emission, and black body radiation due to the telescope itself."
+ A rich literatire has been developed ou the subject (e.g.Maiolit2003:Cushingοἱal. 2001).," A rich literature has been developed on the subject \citep[e.g.][]{maiolino1996,vacca2003,cushing2004}."
+. There are currently uo specilic software packages available for the recductiou of cross-dispersec spectra takei with OSIRIS., There are currently no specific software packages available for the reduction of cross-dispersed spectra taken with OSIRIS.
+ Aimine at providing a fast ancl highly automated task. we developec the (acrouyiu for cross-dispersed spectra reduction script) package.," Aiming at providing a fast and highly automated task, we developed the (acronym for cross-dispersed spectra reduction script) package."
+ The CL language was chosen due to the availability of almost all of the basic tasks needed to perform the reduction i the lmage Reduction aud Analysis Facility (RAF) software (Tody1986.1993)..," The CL language was chosen due to the availability of almost all of the basic tasks needed to perform the reduction in the Image Reduction and Analysis Facility ) software \citep{IRAF1,IRAF2}."
+ In 82 we describe main aspects of the instrument. focusing on its ellects on the reductioi process.," In \ref{sec:osiris} we describe main aspects of the instrument, focusing on its effects on the reduction process."
+ In &)j we describe the general steps towards a fully reduced. spectrum. as well as the approach adopted by the package to each of these steps. aud finally in 81. we give a brief SUnuinuary.," In \ref{sec:reduction} we describe the general steps towards a fully reduced spectrum, as well as the approach adopted by the package to each of these steps, and finally in \ref{sec:summary} we give a brief summary."
+ Iu this section we discuss the main aspects of the cross-dispersed mode of OSIRIS. with special attention to those characteristics that are relevant to the reduction process.," In this section we discuss the main aspects of the cross-dispersed mode of OSIRIS, with special attention to those characteristics that are relevant to the reduction process."
+ A complete description of the instrument cau be found in its on line User's, A complete description of the instrument can be found in its on line User's.
+" The detector is a 102Ix102E ELAWATI array (Hodappetal...1996). sensitive to waveleneths of up to 2.5,un. Equation 1 models the uou-luear beliavior of the array. which ouly becomes critical above 28.000 counts."," The detector is a 1024x1024 HAWAII array \citep{hodapp1996}, sensitive to wavelengths of up to $\upmu$ m. Equation \ref{eq:linearity} models the non-linear behavior of the array, which only becomes critical above 28,000 counts."
+ Usually the detector is read only at the eid of the integration. but since it cali be reac uon-destructively cffe‘elt salipling methods «‘ould be impleijentect.," Usually the detector is read only at the end of the integration, but since it can be read non-destructively different sampling methods could be implemented."
+ A reslcual image is sonetimes seen. syecially when brieht sources are obseved iu acquisition uode.," A residual image is sometimes seen, specially when bright sources are observed in acquisition mode."
+ This means that eventually some of he first spectra taken after the arget acequisitiou images have to be ¢liscarded., This means that eventually some of the first spectra taken after the target acquisition images have to be discarded.
+ Resictals have approximately of he intensity of the original source. aud it should not be a problem to science exposures that have ypleal couuts »elow one thousand.," Residuals have approximately of the intensity of the original source, and it should not be a problem to science exposures that have typical counts below one thousand."
+the black hole rotates faster than the disk. then a steady state with no accretion is built.,"the black hole rotates faster than the disk, then a steady state with no accretion is built."
+ In such a case. all (he power of the disk comes from (he rotational energy of the black hole. the disk has an infinite efficiency.," In such a case, all the power of the disk comes from the rotational energy of the black hole, the disk has an infinite efficiency."
+ If the disk has so much internal torque (hat cannot be balanced by the external torque. the excess internal torque will produceaccretion®.," If the disk has so much internal torque that cannot be balanced by the external torque, the excess internal torque will produce."
+. In such a case. the power of the disk comes from both the rotational energy of the black hole and ihe exeravitational οποιονe. of the disk.," In such a case, the power of the disk comes from both the rotational energy of the black hole and the gravitational energy of the disk."
+ In the case that there is accretion and the magnetic field touches the disk at a circle of radius ry. the quasi-steady solutions are given bv equation (39)) and equation (40)).," In the case that there is accretion and the magnetic field touches the disk at a circle of radius $r_0$, the quasi-steady solutions are given by equation \ref{fdel}) ) and equation \ref{gdel}) )."
+ The radiation [hax is plotted in Fig., The radiation flux is plotted in Fig.
+ 8. ancl Fig. 9..," \ref{fig7} and Fig. \ref{fig8},"
+" respectively for the case that the black hole rotates faster than the disk (0,2 Qo) and for the case that the black hole rotates slower than the disk (0;,< O04).", respectively for the case that the black hole rotates faster than the disk $\Omega_H>\Omega_0$ ) and for the case that the black hole rotates slower than the disk $\Omega_H<\Omega_0$ ).
+" In both cases. for r<rg the solutions are the same as those predicted by the standard theory of a (hin Ixeplerian disk (NovikovandThorneandThorne1974: 1974).. in particular F ancl g ave zero al r=""μι."," In both cases, for $r<r_0$ the solutions are the same as those predicted by the standard theory of a thin Keplerian disk \citep{nov73,pag74,tho74}, in particular $F$ and $g$ are zero at $r=r_{ms}$."
+ The extensions of the standard solutions to 7>ry are shown wilh dashed curves. the positions of rp are shown with vertical dotted lines.," The extensions of the standard solutions to $r>r_0$ are shown with dashed curves, the positions of $r_0$ are shown with vertical dotted lines."
+ Due to the magnetic coupling to the black hole. F ancl yg ave mocilied [or r2ry by superposing the contribution of the magnetic coupling to the standard solutions.," Due to the magnetic coupling to the black hole, $F$ and $g$ are modified for $r>r_0$ by superposing the contribution of the magnetic coupling to the standard solutions."
+" For r>>ry. the radiation fIux given bv the standard theory decreases as oE r7, while the radiation flux contributed by the magnetic coupling decreases as r.7."," For $r\gg r_0$, the radiation flux given by the standard theory decreases as $r^{-3}$ , while the radiation flux contributed by the magnetic coupling decreases as $r^{-3.5}$."
+ Thus. ab large radii the radiation flux. Fis dominated by the contribution of accretion.," Thus, at large radii the radiation flux $F$ is dominated by the contribution of accretion."
+ This is also (ave for the viscous torque in (be disk. ie.. al large radii the torque in the disk is dominated by the contribution of accretion.," This is also true for the viscous torque in the disk, i.e., at large radii the torque in the disk is dominated by the contribution of accretion."
+" When the black hole rotates faster than the disk as the case shown in Fie. ὃ,,"," When the black hole rotates faster than the disk – as the case shown in Fig. \ref{fig7},"
+ the black hole pumps enerev and angular momentum into the disk. the energv is locally dissipated in the disk and eventually radiated away. a bright annular “bump is produced at à=rg.," the black hole pumps energy and angular momentum into the disk, the energy is locally dissipated in the disk and eventually radiated away, a bright annular “bump” is produced at $r=r_0$."
+ When the black hole rotates slower than (he disk as the case shown in Fig. 9..," When the black hole rotates slower than the disk – as the case shown in Fig. \ref{fig8},"
+" the black hole extracts energy. and angular momentum from the disk. and a dark annular ""vallev is produced at r=ry."," the black hole extracts energy and angular momentum from the disk, and a dark annular “valley” is produced at $r=r_0$."
+ When the black hole rotates slower than the disk (i.e.. Oj;« Oo). Ay is negative thus enerev and angular momentum are (ransferecl Irom (he disk tothe black hole.," When the black hole rotates slower than the disk (i.e., $\Omega_H<\Omega_0$ ), $A_0$ is negative thus energy and angular momentum are transfered from the disk tothe black hole."
+ From equation (39)). the minimum radiation [hix at r=rg ds," From equation \ref{fdel}) ), the minimum radiation flux at $r=r_0$ is"
+short duration state was continued until April 17.,short duration state was continued until April 17.
+ Then as the Keplerian rate became comparable to the sub-Keplerian rate. the spectral index rapidly became soft (T— 2.5) and the object remained in the soft-intermediate state for ~26 days until 2010 May l4.," Then as the Keplerian rate became comparable to the sub-Keplerian rate, the spectral index rapidly became soft $\Gamma \sim 2.5$ ) and the object remained in the soft-intermediate state for $\sim 26$ days until 2010 May 14."
+ Here sporadic QPOs were observed and the ΟΡΟ frequency remained at around 6 Hz., Here sporadic QPOs were observed and the QPO frequency remained at around $6$ Hz.
+ After that the Keplerian rate dominates and the spectrum becomes softer with a high spectral index (T> 3.0)., After that the Keplerian rate dominates and the spectrum becomes softer with a high spectral index $\Gamma > 3.0$ ).
+ We believe that when the viscosity is reduced in future. the Keplerian rate will be so much reduced that the Keplerian disk may itself evaporate and the inner edge of the Keplerian could be thought to have receded from the black hole.," We believe that when the viscosity is reduced in future, the Keplerian rate will be so much reduced that the Keplerian disk may itself evaporate and the inner edge of the Keplerian could be thought to have receded from the black hole."
+ In this case. the object would return to the hard state (Mandal Chakrabarti. 2010) and monotonically decreasing nature of the QPO frequencies will be observed as predicted by the POS solution (CDNPO8. CDP09).," In this case, the object would return to the hard state (Mandal Chakrabarti, 2010) and monotonically decreasing nature of the QPO frequencies will be observed as predicted by the POS solution (CDNP08, CDP09)."
+" In Chakrabarti Manickam (2000). the authors showed that the true ""soft photons! r.e.. the photons emitted from the pre-shock flow. do not exhibit QPOs."," In Chakrabarti Manickam (2000), the authors showed that the true `soft photons' i.e., the photons emitted from the pre-shock flow, do not exhibit QPOs."
+ A cross-check that the two-component model is valid and that the QPO is caused by the oscillation of the Comptonized cloud (namely. the shock region) can be made.," A cross-check that the two-component model is valid and that the QPO is caused by the oscillation of the Comptonized cloud (namely, the post-shock region) can be made."
+ As the shock moved in. the post- flow was heated up roughly as T~1/r while the pre-shock Keplerian flow was heated up as 77 (Shakura Sunyaev. 1973).," As the shock moved in, the post-shock flow was heated up roughly as $T\sim 1/r$ while the pre-shock Keplerian flow was heated up as $ r^{-3/4}$ (Shakura Sunyaev, 1973)."
+" This means that so-called ""soft-photons' in 2—4 keV range could be actually Comptonized photons at a large distance and should be treated as hard photons.", This means that so-called `soft-photons' in $2-4$ keV range could be actually Comptonized photons at a large distance and should be treated as hard photons.
+ Thus. 2—4 keV photons should show QPOs at a large distance only.," Thus, $2-4$ keV photons should show QPOs at a large distance only."
+ In Fig., In Fig.
+" 5(a-c) we plot the power density spectra of photons in the 2—4 keV. 4—15 keV and 15—30 keV bands on the April 5. 13 and 16 respectively when the shock (Compton cloud boundary) was located at ~780r,.3677, and 244r, respectively with compression ratios of 2.2. 1.27 and 1.08 respectively."," 5(a-c) we plot the power density spectra of photons in the $2-4$ keV, $4-15$ keV and $15-30$ keV bands on the April 5, 13 and 16 respectively when the shock (Compton cloud boundary) was located at $\sim 780 r_g, 367 r_g$ and $244 r_g$ respectively with compression ratios of $2.2$, $1.27$ and $1.08$ respectively."
+ The spectral and timing properties of the object in and around these days are presented in Table | and Table 2., The spectral and timing properties of the object in and around these days are presented in Table 1 and Table 2.
+ In Table | we show how the photon index and fluxes vary in these days., In Table 1 we show how the photon index and fluxes vary in these days.
+ y and the number of degrees of freedom (dof) in the spectral fits are also included., $\chi^2$ and the number of degrees of freedom (dof) in the spectral fits are also included.
+ In Table 2 we note that the rms amplitude of the fundamental QPO decreased with time in 2—4 keV range., In Table 2 we note that the rms amplitude of the fundamental QPO decreased with time in $2-4$ keV range.
+ This shows that the number of Comptonized photon is decreasing in this energy bin and increasing in the 4—15 keV bin as the shock moves in., This shows that the number of Comptonized photon is decreasing in this energy bin and increasing in the $4-15$ keV bin as the shock moves in.
+ We note that in the 4—15 keV range. the rms initially decreased and then increased.," We note that in the $4-15$ keV range, the rms initially decreased and then increased."
+ The latter increase is because the entire Comptonized photons now belong to this energy bin., The latter increase is because the entire Comptonized photons now belong to this energy bin.
+ In high-energy channels (15—30 keV) the rms amplitude decreased because of the paucity of photons., In high-energy channels $15-30$ keV) the rms amplitude decreased because of the paucity of photons.
+ Clearly Table 2 shows that the rms amplitude was highest in all energy ranges when the shock was farther out., Clearly Table 2 shows that the rms amplitude was highest in all energy ranges when the shock was farther out.
+ This shows that the oscillation of the Compton cloud. which was caused by a stronger shock at a large distance. was high enough to significantly modulate the outgoir£g flux.," This shows that the oscillation of the Compton cloud, which was caused by a stronger shock at a large distance, was high enough to significantly modulate the outgoing flux."
+ This ts particularly surprising in the 2—4 keV range. because it was supposed to be the energy of the ‘soft-photons’.," This is particularly surprising in the $2-4$ keV range, because it was supposed to be the energy of the `soft-photons'."
+ This shows that far out. these photons are actually Comptonized photons.," This shows that far out, these photons are actually Comptonized photons."
+ Note that the QPO is absent in this range on April 16., Note that the QPO is absent in this range on April 16.
+ This is expected. because the Compton cloud is so close to the black hole that these photons are contributed mostly from the Keplerian flow and are not modulated.," This is expected, because the Compton cloud is so close to the black hole that these photons are contributed mostly from the Keplerian flow and are not modulated."
+ Note that the appearance of harmonics (Yu. 2010) depends on the energy of the photons.," Note that the appearance of harmonics (Yu, 2010) depends on the energy of the photons."
+ On the April 16. the shock is very weak and a faint ΟΡΟ is seen only in 4—]5 keV. At lower energy bins QPO was not seen because they were emitted from the pre-shock flow. and at higher energy bins. the number of photons was statistically insignificant.," On the April 16, the shock is very weak and a faint QPO is seen only in $4-15$ keV. At lower energy bins QPO was not seen because they were emitted from the pre-shock flow, and at higher energy bins, the number of photons was statistically insignificant."
+ This analysis indicates that the general two-component solution of the outburst sources as developed in CDNPO8. CDPO9. and Mandal Chakrabarti (2010) can explain most of what is seen in GX 339-4 so far.," This analysis indicates that the general two-component solution of the outburst sources as developed in CDNP08, CDP09, and Mandal Chakrabarti (2010) can explain most of what is seen in GX 339-4 so far."
+ Assuming that this outburst is similar, Assuming that this outburst is similar
+light at Ty= MJD 53915.7+0.2.,light at $T_0 =$ MJD $53945.7 \pm 0.2$.
+" In addition. we find in our data another period rauging [rom 31.2 to 31.6 days. which we attribute to a ""superlhump"" resonance in the accretion dise (Haswelletal.2001).."," In addition, we find in our data another period ranging from $31.2$ to $31.6$ days, which we attribute to a “superhump"" resonance in the accretion disc \citep{Haswell:2001uq}."
+ The superliuunp period is longer than the orbital period. iudicatiug that the precession is prograde with respect to the rotation of the dise (Whitehurst&Ixiug1991)..," The superhump period is longer than the orbital period, indicating that the precession is prograde with respect to the rotation of the disc \citep{Whitehurst:1991fk}."
+ The fractional dillerence e between the orbital aud πρό periods is determiued to lie between 0.010 aud 0.026., The fractional difference $\epsilon$ between the orbital and superhump periods is determined to lie between $0.010$ and $0.026$.
+ Our measurement of e allows us to make a prediction of the mass ratio q of the system. based on au empirical formula 2005)..," Our measurement of $\epsilon$ allows us to make a prediction of the mass ratio $q$ of the system, based on an empirical formula \citep{Patterson:2005qy}."
+ We find that 0.05<q0.12. providing au independent verification of the result q=0.09£0.03 due to Greineretal.(2001a)..," We find that $0.05 < q < 0.12$, providing an independent verification of the result $q = 0.09 \pm 0.03$ due to \cite{Greiner:2001kx}."
+ The verification of the mass ratio q via the superliump period is of particular importance., The verification of the mass ratio $q$ via the superhump period is of particular importance.
+ The previous determinatiou of the mass ratio was based on spectroscopic identification of the secondary star (Greineretal.2001b).. which requires a iunber of assumptions: our inetliod is based on au uurelated set of assumptious.," The previous determination of the mass ratio was based on spectroscopic identification of the secondary star \citep{Greiner:2001vn}, which requires a number of assumptions; our method is based on an unrelated set of assumptions."
+ Couversely. our result cau be seen as a verification of the empirical relation between e aud gy. iu a system with au extremely long orbital period. compared to other known superliumpers (Pattersonetal.2005)..," Conversely, our result can be seen as a verification of the empirical relation between $\epsilon$ and $q$, in a system with an extremely long orbital period compared to other known superhumpers \citep{Patterson:2005qy}."
+ We thank the teams at MIT aud the Science Operatious Facility aud Guest Observer Facility at NASA’s Goddard Space Flight Center for providing the X-ray results., We thank the teams at MIT and the Science Operations Facility and Guest Observer Facility at NASA's Goddard Space Flight Center for providing the X-ray results.
+ We are grateful for the efforts of YALO/SMARTS service observers D. Couzalez auc J. Espinoza. who collected the data reported here.," We are grateful for the efforts of YALO/SMARTS service observers D. Gonzalez and J. Espinoza, who collected the data reported here."
+ This work was supported by NSF erant. AST 01-07063 to C.D.B., This work was supported by NSF grant AST 04-07063 to C.D.B.
+Stellar. population svuthesis (SPS) models are au iuportanut astrophysical tool for iuterpreting photometric and spectroscopic observations of e@alaxics.,Stellar population synthesis (SPS) models are an important astrophysical tool for interpreting photometric and spectroscopic observations of galaxies.
+ Alocdern surveys can detect vast umubers of galaxies across a range of redshifts and can provide importan observational coustraints on star formation historics and galaxw evolution., Modern surveys can detect vast numbers of galaxies across a range of redshifts and can provide important observational constraints on star formation histories and galaxy evolution.
+" Ilowewer. it is eenerally no possible to resolve the stellar populations of these walasios,"," However, it is generally not possible to resolve the stellar populations of these galaxies."
+ Instead. their properties are offen estimated by conrpariue their iutegrated euission with that predicte from composite stellar population models.," Instead, their properties are often estimated by comparing their integrated emission with that predicted from composite stellar population models."
+ Such moclels conibine our knowledge of stellar evolutiou with stellar spectral libraries to predict the iuteerated omission of an unresolved stellar population., Such models combine our knowledge of stellar evolution with stellar spectral libraries to predict the integrated emission of an unresolved stellar population.
+ Iu their simplest fori. simple stellar population (SSP) models. SPS models xedict the integrated cussion frou a stellar population with a sinele metallicity aud star formation epoch.," In their simplest form, simple stellar population (SSP) models, SPS models predict the integrated emission from a stellar population with a single metallicity and star formation epoch."
+ The ability of SPS models to accurately represent real stellar »opulatious has a direct influence on ealactic studies., The ability of SPS models to accurately represent real stellar populations has a direct influence on galactic studies.
+ For many vears. star clusters lave provided some of he best locations to calibrate paracters i SPS models (not all of which can be derived theoretically) aud to test heir ability to reproduce the observed properties.," For many years, star clusters have provided some of the best locations to calibrate parameters in SPS models (not all of which can be derived theoretically) and to test their ability to reproduce the observed properties."
+ This is because the relatively simple stellar populations of these clusters can be directly compared with such models., This is because the relatively simple stellar populations of these clusters can be directly compared with such models.
+" Photometric and spectroscopic observations of the Milky Wavs (old) globular chisters. (vouug) open clusters and the Larec/Small Magellanic Cloud’s clusters. have previously been used to calibrate aud test such mo06dels (e.g,3227)."," Photometric and spectroscopic observations of the Milky Way's (old) globular clusters, (young) open clusters and the Large/Small Magellanic Cloud's clusters have previously been used to calibrate and test such models \citep[e.g.][]{Renzini88,VandenBerg03,Maraston05,An08,Lyubenova10,Riffel11}."
+" These studies. which often consider observations obtained in the Johuson-Cousius VRB.I, filters.. have sheAVL reasonable agreciucnt between models and. observations,"," These studies, which often consider observations obtained in the Johnson-Cousins $_{c}$ $_{c}$ filters, have shown reasonable agreement between models and observations."
+" Iknever, several discrepancies. nxst notably iu the BV (eg.77). and near-infrared have demonstrated. colors."," However, several discrepancies, most notably in the B-V \citep[e.g.][]{Worthey94,Maraston05} and near-infrared \citep[e.g.][]{Cohen07,Conroy10} colors, have been demonstrated."
+"Despite the heen popularity aud use of. the SDS5EIN5 ae (T)... ο...relatively little work exists which filterSoD models to star clusters iu these filters,"," Despite the increasing popularity and use of the SDSS's filter set \citep{Fukugita96}, relatively little work exists which compares SSP models to star clusters in these filters."
+ This is due conparesprimarily to a relative lack of cluster photometry hrough these. compared with the Johuson-Cousius. filters.," This is due primarily to a relative lack of cluster photometry through these, compared with the Johnson-Cousins, filters."
+ Iowever. there have been some sugeestions of offsets between SPS models aud SDSS observations.," However, there have been some suggestions of offsets between SPS models and SDSS observations."
+ Iu uticular. studies of luminous red galaxies (LRGs. which are expected to have a predominantly old. metal rich. stellar population) have shown siguificant offsets frou he model predictions. particularly around the gri--bauds (e.g.22).," In particular, studies of luminous red galaxies (LRGs, which are expected to have a predominantly old, metal rich, stellar population) have shown significant offsets from the model predictions, particularly around the -bands \citep[e.g.][]{Eisenstein01,Wake06}."
+ Several explanations have beeu proposed to rectify this offset., Several explanations have been proposed to rectify this offset.
+ ?. proposed that slight discrepancies vetween theoretical and cuipivical stellar libraries may, \citet{Maraston09} proposed that slight discrepancies between theoretical and empirical stellar libraries may
+The isothermal plasma was chosen to have logZ7=6.0 (in K) and logEM.=43.0 (in em).,The isothermal plasma was chosen to have $\log T=6.0$ (in K) and $\log EM=43.0$ (in $^{-3}$ ).
+ This temperature was chosen as it allowed us to include lines routinely observed by available spectrometers such as SOHO/CDS. SOHO/SUMER. and Hinode/EIS.," This temperature was chosen as it allowed us to include lines routinely observed by available spectrometers such as SOHO/CDS, SOHO/SUMER, and Hinode/EIS."
+ Figure 2 shows the EM loci technique applied to the calculated intensities. without random errors included: the original temperature and EM values are recovered without problems.," Figure \ref{1-isot-1} shows the EM loci technique applied to the calculated intensities, without random errors included: the original temperature and EM values are recovered without problems."
+ Figure 3 shows the MCMC reconstruction obtained without random errors and varying the bin size W., Figure \ref{1-isot-3} shows the MCMC reconstruction obtained without random errors and varying the bin size $W$ .
+ Results are summarized in Table 2.., Results are summarized in Table \ref{metrics}.
+ In all cases the MCMC technique is able to provide an isolated peak several order of magnitude larger than the background., In all cases the MCMC technique is able to provide an isolated peak several order of magnitude larger than the background.
+ —In the case of bin width W=0.01 the shape of the peak is irregular and one additional satellite peak is present at higher temperatures: this feature 1s an artifact and it disappears when the bir width increases., In the case of bin width $W=0.01$ the shape of the peak is irregular and one additional satellite peak is present at higher temperatures: this feature is an artifact and it disappears when the bin width increases.
+ In all cases the peak maintains a non-negligible width Adog7)=0.03—0.05 that is approximately constant as the bin width W increases.," In all cases the peak maintains a non-negligible width $\Delta{\left({\log T}\right)} 
+\approx 0.03-0.05$ that is approximately constant as the bin width $W$ increases."
+ We interpret this as the ultimate resolution capability of the MCMC technique to resolve a purely isothermal plasma: the result is indistinguishable from a multithermal plasma with a very narrow width., We interpret this as the ultimate resolution capability of the MCMC technique to resolve a purely isothermal plasma: the result is indistinguishable from a multithermal plasma with a very narrow width.
+ Even if the y of the solution with W=0.1 is the lowest. the solution with W0.05 ts to be preferred because the bin width is closer to the apparently intrinsic resolving power of the MCMC technique.," Even if the $\chi^2$ of the solution with W=0.1 is the lowest, the solution with W=0.05 is to be preferred because the bin width is closer to the apparently intrinsic resolving power of the MCMC technique."
+ In the case with W=0.05 the two bins at the side of the peak show EM values of «15€ of the peak value. and might be taken as real.," In the case with $W=0.05$ the two bins at the side of the peak show EM values of $\approx$ of the peak value, and might be taken as real."
+ Therefore. care should be taken in interpreting such features.," Therefore, care should be taken in interpreting such features."
+ The solutions provided by the MCMC technique always have a background at very low EM values at all temperatures (logEM=33— 37)., The solutions provided by the MCMC technique always have a background at very low EM values at all temperatures $\log EM \simeq 33-37$ ).
+ This background represents the maximum values of the EM at all temperatures outside the isothermal, This background represents the maximum values of the EM at all temperatures outside the isothermal
+cells since this is the configuration to which they tend to evolve.,cells since this is the configuration to which they tend to evolve.
+ We therefore set up the particles in this manner. with the (hexagonal) lattice centred on the origin and again with the particle masses set according to the equilibrium density profile.," We therefore set up the particles in this manner, with the (hexagonal) lattice centred on the origin and again with the particle masses set according to the equilibrium density profile."
+ 1045 particles are used., 1045 particles are used.
+" In this case the lattice is able to damp to an equilibrium state since the particles do not shift from the lattice apart from a small initial acljustment at the outer edges (note that we also use à p,;, in (23)) as will be discussed in refsec:axisvm)).", In this case the lattice is able to damp to an equilibrium state since the particles do not shift from the lattice apart from a small initial adjustment at the outer edges (note that we also use a $\rho_{min}$ in \ref{eq:hrho}) ) as will be discussed in \\ref{sec:axisym}) ).
+ The equilibrium distribution in this case is shown in Figure 3. and the density profile is shown in Figure 4.., The equilibrium distribution in this case is shown in Figure \ref{fig:static_varm} and the density profile is shown in Figure \ref{fig:static_varm_rho}.
+ The relaxation procedure for the Toy Star also revealed some pertinent aspects of the reversible integration algorithm., The relaxation procedure for the Toy Star also revealed some pertinent aspects of the reversible integration algorithm.
+ In particular if the timestep changes rapidly between steps. the energy in the reversible integrator could. exhibit. [large oscillations and under some circumstances become unstable.," In particular if the timestep changes rapidly between steps, the energy in the reversible integrator could exhibit large oscillations and under some circumstances become unstable."
+ This is a result of the averaging procedure used to determine the timestep., This is a result of the averaging procedure used to determine the timestep.
+ In this paper we have used the arithmetic average where the superscripts 0.3 ancl 1 refer to the stepsize calculated. at. the. beginning. middle and end of a given timestep.," In this paper we have used the arithmetic average where the superscripts $0,\frac12$ and $1$ refer to the stepsize calculated at the beginning, middle and end of a given timestep."
+" This expression is rearranged in order to determine the next timestep 73 from the previous timestep of"" and the timestep determined from the Courant condition at the half step 0/7.", This expression is rearranged in order to determine the next timestep $\delta t^{1}$ from the previous timestep $\delta t^{0}$ and the timestep determined from the Courant condition at the half step $\delta t^{1/2}$.
+" ""Phe pitfall of this method is that if the timestep changes rapidly (for example. from a very short step to a very long step) then the timestep begins to oscillate between two values (short-Iong-short-lIong)."," The pitfall of this method is that if the timestep changes rapidly (for example, from a very short step to a very long step) then the timestep begins to oscillate between two values (short-long-short-long)."
+ Also there is nothing to prevent the timestep o/* from taking a negative value., Also there is nothing to prevent the timestep $\delta t^{1}$ from taking a negative value.
+ In this situation the time advance is achieved by taking a long step Forward and a slightly shorter step back to achieve a small net. forward step., In this situation the time advance is achieved by taking a long step forward and a slightly shorter step back to achieve a small net forward step.
+ Whilst the specific problem. of negative timesteps can be fixed by using a different averaging xocedure (for example using the geometric mean). we have ound that the oscillations can also cause instability if o greatly exceeds the Courant condition.," Whilst the specific problem of negative timesteps can be fixed by using a different averaging procedure (for example using the geometric mean), we have found that the oscillations can also cause instability if $\delta t^{1}$ greatly exceeds the Courant condition."
+ This can occur oeause a very short df)σι impliesH a very long afcel (where wv definition. οἱ2 meets the Courant conclition)., This can occur because a very short $\delta t^{0}$ implies a very long $\delta t^{1}$ (where by definition $\delta t^{1/2}$ meets the Courant condition).
+ If the Courant constraint also changes between steps the result can be that particles are evolved on a timestep which is too arge to capture the change in physical quantities accurately. resulting in instability.," If the Courant constraint also changes between steps the result can be that particles are evolved on a timestep which is too large to capture the change in physical quantities accurately, resulting in instability."
+ In practice these oscillations can be largely avoided ηον]ος that some care is taken in the setup and evolution., In practice these oscillations can be largely avoided provided that some care is taken in the setup and evolution.
+ For example in the initial runs we began the integration as we had. done for the predictor-corrector method using a timestep of zero (which simply runs through the force evaluation without evolving the particles)., For example in the initial runs we began the integration as we had done for the predictor-corrector method using a timestep of zero (which simply runs through the force evaluation without evolving the particles).
+ In the reversible integrator this immediately. leads to timestep oscillations since the zero timestep is then used in the averaging procedure., In the reversible integrator this immediately leads to timestep oscillations since the zero timestep is then used in the averaging procedure.
+ A second example is that it is common practice to reduce the timestep before output. dumps in order to reach the specified. time of output οχαστν., A second example is that it is common practice to reduce the timestep before output dumps in order to reach the specified time of output exactly.
+ Ες means taking a shorter step just. before an output. dump. but, This means taking a shorter step just before an output dump but
+AHPS22530 both show significant masses of cold CO gas.,MIPS22530 both show significant masses of cold CO gas.
+ Unfortunately both of these are among the sources. where the feature falls in the unreliable parts of the co-adcded spectra.," Unfortunately both of these are among the sources, where the feature falls in the unreliable parts of the co-added spectra."
+ Moreover. the presence of CO does not automatically imply the presence of the CCO feature which arises from the much warmer regions in the vicinity of the galactic nuclei.," Moreover, the presence of CO does not automatically imply the presence of the CO feature which arises from the much warmer regions in the vicinity of the galactic nuclei."
+ We would expect to see this feature in our sources as il is seen in some deeply obscured local ULIRGs (Spoonetal.2005)., We would expect to see this feature in our sources as it is seen in some deeply obscured local ULIRGs \citep{spoon_co}.
+. We do show some evidence for the CCO absorption feature in three of these sources., We do show some evidence for the CO absorption feature in three of these sources.
+ Although tentative. this is the first suggestion of this feature in such high redshift svstems.," Although tentative, this is the first suggestion of this feature in such high redshift systems."
+ However this feature is verv weak even in these sources and only seen in 3 of the 7 sources will usable spectra (and is nol present in the stacked spectra)., However this feature is very weak even in these sources and only seen in 3 of the 7 sources with usable spectra (and is not present in the stacked spectra).
+ We can therefore conclude that it is either not ubiquitous even alone sources will silicate absorption. or more realistically that the feature manifests wilh a range of absorption depths. and we only see the most pronounced features.," We can therefore conclude that it is either not ubiquitous even among sources with silicate absorption, or more realistically that the feature manifests with a range of absorption depths, and we only see the most pronounced features."
+ As discussed in 44.5. à 7 sstarburst continuum would effectively weaken the observed CO feature depth (as that likely originates close to the AGN itself).," As discussed in 4.5, a $\sim$ starburst continuum would effectively weaken the observed CO feature depth (as that likely originates close to the AGN itself)."
+ For weak-PAII sources. the ~ cecontinuum is most likely dominated bv the AGN (Sajinaetal.2007a).. however a starburst contribution. sufficient to affect the CO depth. cannot be excluded.," For weak-PAH sources, the $\sim$ continuum is most likely dominated by the AGN \citep{sajina07}, however a starburst contribution, sufficient to affect the CO depth, cannot be excluded."
+ The fact that one of our two strone-PAIT sources (MIPS22432) is also one of three sources showing some hints of CO absorption. suggests that (Bis is much more complicated that a simple correction scaled {ο the starburst strength (as revealed by the PAIL).," The fact that one of our two strong-PAH sources (MIPS22432) is also one of three sources showing some hints of CO absorption, suggests that this is much more complicated that a simple correction scaled to the starburst strength (as revealed by the PAH)."
+ Drawing a parallel with the dust-embecded Galactic YSOs. perhaps there is a range from purelv solid form CO (o fully gaseous CO as the temperatures rise and the CO previously deposited on grain surfaces sublimates (seeWhittet.2003.:andreferencestherein)..," Drawing a parallel with the dust-embedded Galactic YSOs, perhaps there is a range from purely solid form CO to fully gaseous CO as the temperatures rise and the CO previously deposited on grain surfaces sublimates \citep[see][; and references therein]{whittet03}."
+ phase CO shows a much narrower feature (han ils gas-phase counterpart. and in our spectra would be undetectable. unless exceptionally strong.," Solid-phase CO shows a much narrower feature than its gas-phase counterpart, and in our low-resolution spectra would be undetectable, unless exceptionally strong."
+ Ii local ULIBGs. this problem is present as well. and only the high resolution spectrum of RA8S085724-3915 has allowed to rule out CO ice in this source (Geballeetal. 2006)..," In local ULIRGs, this problem is present as well, and only the high resolution spectrum of IRAS08572+3915 has allowed to rule out CO ice in this source \citep{geballe06}. ."
+ To date. only NGC4945 has a CO ice detection (Spoonetal.2003).," To date, only NGC4945 has a CO ice detection \citep{spoon03}."
+. Following the argument in Lutzetal.(2004).. we might also conclude that we are not seeing (his feature due (ο dust tori obscuration rather than buried nuclei.," Following the argument in \citet{lutz04}, we might also conclude that we are not seeing this feature due to dust tori obscuration rather than `buried nuclei'."
+ However. this runs counter to our conclusions [rom the strong IHLAC: absorption features (see 55.3).," However, this runs counter to our conclusions from the strong HAC absorption features (see 5.3)."
+ Ultimately. there are too few CCO detections even locally (see Introduction) aud thereforethe statistics on what [raction," Ultimately, there are too few CO detections even locally (see Introduction) and thereforethe statistics on what fraction"
+pe) suggests that the supersoft channel is almost an order of magnitude more important than the red-giant channel. though presumably some single LMWDs destined to be single UCWDs have not yet cooled sufficiently to become UCWDs. and some UCWDs may not be LMWDs.,") suggests that the supersoft channel is almost an order of magnitude more important than the red-giant channel, though presumably some single LMWDs destined to be single UCWDs have not yet cooled sufficiently to become UCWDs, and some UCWDs may not be LMWDs."
+ The relative importance of these formation channels will be an interesting quantity to constrain with future data and compare with binary population synthesis models., The relative importance of these formation channels will be an interesting quantity to constrain with future data and compare with binary population synthesis models.
+ Perhaps more decisively. the estimates of Hansen (2003) and section 4.4. imply that the total number of observed remnants is consistent with the expected number of past SNe Ia in our Galaxy.," Perhaps more decisively, the estimates of Hansen (2003) and section \ref{sec:PopNumb} imply that the total number of observed remnants is consistent with the expected number of past SNe Ia in our Galaxy."
+ If this is confirmed by future work. it would leave room for only a minority of SNe Ia to result from double-degenerate systems.," If this is confirmed by future work, it would leave room for only a minority of SNe Ia to result from double-degenerate systems."
+ One curiosity of our study of UCWDs is that the range of orbital periods we infer for the SN la progenitors at explosion is only broadly consistent. with that predicted by Hachisu Kato (2001) for systems with red-giant donors., One curiosity of our study of UCWDs is that the range of orbital periods we infer for the SN Ia progenitors at explosion is only broadly consistent with that predicted by Hachisu Kato (2001) for systems with red-giant donors.
+ Their models produce no SNe Ia with final orbital periods 25d<Powx60d: most of the final parameter space for red-giant donors has 100d<Poyx1000d.," Their models produce no SNe Ia with final orbital periods ${\rm 2.5~d~\lesssim}~P_{\rm orb}~{\rm
+\lesssim~60~d}$; most of the final parameter space for red-giant donors has ${\rm 100~d~\lesssim}~P_{\rm orb}~{\rm
+\lesssim~1000~d}$."
+ Our favoured assumptions seem to indicate orbital periods on the shorter side of 100 d. re. the kinematic signal seems to be a little stronger than expected.," Our favoured assumptions seem to indicate orbital periods on the shorter side of 100 d, i.e. the kinematic signal seems to be a little stronger than expected."
+ A possible resolution could be that our use of UCWDs is biased and preferentially selects systems with shorter orbital periods because they produce lower-mass WDs. which cool more rapidly.," A possible resolution could be that our use of UCWDs is biased and preferentially selects systems with shorter orbital periods because they produce lower-mass WDs, which cool more rapidly."
+ Or perhaps the progenitor population is preferentially from the thick disc. and hence the remnants falsely appear kinematically hotter than in our models.," Or perhaps the progenitor population is preferentially from the thick disc, and hence the remnants falsely appear kinematically hotter than in our models."
+ It may also be that some non-LMWD single UCWDs from the thick disc or Galactic halo could be contaminating these results., It may also be that some non-LMWD single UCWDs from the thick disc or Galactic halo could be contaminating these results.
+ Alternatively. this may be an indication that the models need some modification: for example the systems containing redgiant donors might produce SNe Ia earlier than expected. or perhaps a subset of systems from the supersoft channel does not explode until their donor stars have evolved more than the current models predict.," Alternatively, this may be an indication that the models need some modification; for example the systems containing red-giant donors might produce SNe Ia earlier than expected, or perhaps a subset of systems from the supersoft channel does not explode until their donor stars have evolved more than the current models predict."
+ We encourage further work on the nature of UCWDs., We encourage further work on the nature of UCWDs.
+ It is Important to confirm that the apparently single UCWDs really are single., It is important to confirm that the apparently single UCWDs really are single.
+ High signal-to-noise searches for any spectral features that allow radial-velocity measurements would be worthwhile., High signal-to-noise searches for any spectral features that allow radial-velocity measurements would be worthwhile.
+ GAIA astrometry is a long-term hope for examining the single status of UCWDs. as well as providing accurate distances to all these objects.," GAIA astrometry is a long-term hope for examining the single status of UCWDs, as well as providing accurate distances to all these objects."
+ It is also important to understand the mass distribution of the UCWD population., It is also important to understand the mass distribution of the UCWD population.
+ The inferred mass of the best-studied UCWD. LHS 3250. (0.23M. [Bergeron Leggett. 2002]) makes it a clear LMWD and hence a good candidate SN Ia remnant. but no other apparently single UCWD has a constrained mass estimate that we are aware of.," The inferred mass of the best-studied UCWD, LHS 3250, $\rm 0.23~M_{\odot}$ [Bergeron Leggett, 2002]) makes it a clear LMWD and hence a good candidate SN Ia remnant, but no other apparently single UCWD has a well-constrained mass estimate that we are aware of."
+ A sample of non-UCWD single LMWDs has distinct advantage over the UCWDs in that the prescence of spectral lines allows for radial velocity measurements (see. e.g.. Maxted et 22000).," A sample of non-UCWD single LMWDs has distinct advantage over the UCWDs in that the prescence of spectral lines allows for radial velocity measurements (see, e.g., Maxted et 2000)."
+ The significant difficulty for these objects is in selecting a large sample: whilst UCWDs stand out from survey photometry. hotter LMWD do not.," The significant difficulty for these objects is in selecting a large sample: whilst UCWDs stand out from survey photometry, hotter LMWD do not."
+ Eisenstein et (2006) used the Sloan digital sky survey to produce a catalogue of over 9000 white dwarfs. and identify 13 WDs with masses <0.3M. of which 7 have masses €0.2M4.," Eisenstein et (2006) used the Sloan digital sky survey to produce a catalogue of over 9000 white dwarfs, and identify 13 WDs with masses $\rm < 0.3 ~{\rm M_{\odot}}$, of which 7 have masses $\rm
+\leq 0.2 ~{\rm M_{\odot}}$."
+ These objects should be investigated for signs of a companion NS. or for radial velocity variations.," These objects should be investigated for signs of a companion NS, or for radial velocity variations."
+ With a substantial sample of single LMWDs. it would make sense to use directional information about the space velocities of the objects rather than just the magnitudes of the transverse velocities.," With a substantial sample of single LMWDs, it would make sense to use directional information about the space velocities of the objects rather than just the magnitudes of the transverse velocities."
+ We expect that the use of such information will help distinguish between a large initial scale-height. long orbital period population and a thinner. shorter-period population.," We expect that the use of such information will help distinguish between a large initial scale-height, long orbital period population and a thinner, shorter-period population."
+ Our conclusions would be considerably stronger if we were sure about the initial kinematies of these stars; it is important to try and determine whether the remnants come from the thick disc., Our conclusions would be considerably stronger if we were sure about the initial kinematics of these stars; it is important to try and determine whether the remnants come from the thick disc.
+ Once the single LMWD sample becomes large enough. it could potentially be split into sub-samples with different WD masses.," Once the single LMWD sample becomes large enough, it could potentially be split into sub-samples with different WD masses."
+ If our model is correct. then less massive single LMWDs should be kinematically hotter than more massive single LMWDs (see section 3.4)).," If our model is correct, then less massive single LMWDs should be kinematically hotter than more massive single LMWDs (see section \ref{sec:OrbitalVelocities}) )."
+ Since we began this work. van Leeuwen et ((2007) and Kilic et al. (," Since we began this work, van Leeuwen et (2007) and Kilic et al. ("
+2007a) have searched for companions to LMWDs in the radio and optical wavebands. respectively.,"2007a) have searched for companions to LMWDs in the radio and optical wavebands, respectively."
+" The search by van Leeuwen et oof LMWDs for radio pulsations found none ‘down to flux densities of 00.8 mJy Κροτ, and concluded that “a given low-mass helitum-core white dwarf has a probability of < 0.18+0.05 of being in a binary with a radio pulsar’."," The search by van Leeuwen et of LMWDs for radio pulsations found none `down to flux densities of 0.8 mJy $\rm{^{-2}}$ ', and concluded that `a given low-mass helium-core white dwarf has a probability of $<$ $\pm$ 0.05 of being in a binary with a radio pulsar'."
+" For four WDs with masses <0.2M... Kilic et ffound: ""None of these white dwarfs show excess emission from a binary companion. and radial velocity searches will be necessary to constrain the nature of the unseen companions,"," For four WDs with masses $\rm < 0.2 ~{\rm M_{\odot}}$, Kilic et found: `None of these white dwarfs show excess emission from a binary companion, and radial velocity searches will be necessary to constrain the nature of the unseen companions.'"
+ Our paper suggests that the that there are unseen companions Is not nessecary., Our paper suggests that the that there are unseen companions is not nessecary.
+ Perhaps the most interesting possibility for the evolutionary state of the donor in a single-degenerate SN Ia is that of a hot subdwarf (sdO or sdB) star., Perhaps the most interesting possibility for the evolutionary state of the donor in a single-degenerate SN Ia is that of a hot subdwarf (sdO or sdB) star.
+ This would naturally allow short orbital periods (~| h) and also naturally produce extremely low-mass WDs. as recently observed in the runaway WD LP 400-22 (Kawka et 22006). which is inferred to have a mass of 0.17. M. and a tangential velocity of 414 + 43 km κ.," This would naturally allow short orbital periods $\rm
+\sim 1~h$ ) and also naturally produce extremely low-mass WDs, as recently observed in the runaway WD LP 400-22 (Kawka et 2006), which is inferred to have a mass of 0.17 $\rm ~{\rm M_{\odot}}$ and a tangential velocity of 414 $\rm \pm$ 43 km $\rm {\rm s^{-1}}$."
+ This would also provide a natural explanation for stars like the runaway hot subdwarf US 708 (discovered by Hirsch et 22005)., This would also provide a natural explanation for stars like the runaway hot subdwarf US 708 (discovered by Hirsch et 2005).
+ We feel that a SN Ia origin for this object is more satisfying than a scenario combining dynamical ejection from the supermassive black hole in the Galactic centre with the simultaneous merger of two helium WDs (as speculated by, We feel that a SN Ia origin for this object is more satisfying than a scenario combining dynamical ejection from the supermassive black hole in the Galactic centre with the simultaneous merger of two helium WDs (as speculated by
+"We note that. in a nonmagnetized disc. the azimuthal perturbed velocity varies as low with =fron,|fry,0 near corotation (ΓΕτο). whereas here it varies as 1/7 near the magnetic resonances.","We note that, in a nonmagnetized disc, the azimuthal perturbed velocity varies as $\ln x$ with $x= \left| r - r_p \right| / r_p
+\rightarrow 0$ near corotation (GT79), whereas here it varies as $1/x$ near the magnetic resonances."
+ At the locations where the solutions of equation (26)) change from wavelike to non wavelike (turning points). the disc response varies on a large scale. and therefore it couples well to the perturbing potential.," At the locations where the solutions of equation \ref{EQ2}) ) change from wave–like to non wave–like (turning points), the disc response varies on a large scale, and therefore it couples well to the perturbing potential."
+ Εις results in a large contribution to the tidal torque (see CE79 and 6))., This results in a large contribution to the tidal torque (see GT79 and \ref{sec:torque}) ).
+ IH is therefore important to calculate the location of these turning points., It is therefore important to calculate the location of these turning points.
+ The WKB analysis above (see 5.2)) indicates that the dilferential equation (26)) has turning points at the locations of the socallednominat Lindblad resonances r=rg., The WKB analysis above (see \ref{sec:wkb}) ) indicates that the differential equation \ref{EQ2}) ) has turning points at the locations of the so–called Lindblad resonances $r=r_L$.
+ Waves propagate outside these resonances. which are located on both sides of the corotation radius.," Waves propagate outside these resonances, which are located on both sides of the corotation radius."
+ In the nonmagnetic case. it is known that the turning points do no coinncide with the nominal Lindblad resonances for values of m in excess o£ e/4 (Coldreich Tremaine 1980. Xrtvmowicz 1993).," In the nonmagnetic case, it is known that the turning points do no ncide with the nominal Lindblad resonances for values of $m$ in excess of $r/H$ (Goldreich Tremaine 1980, Artymowicz 1993)."
+ LW they cid comncide. the turningpoints (also called. Lindblad resonances) would converge toward ry as m increases.," If they did ncide, the turningpoints (also called Lindblad resonances) would converge toward $r_p$ as $m$ increases."
+ Enstead. they converge toward a radius which is located a finite distance (equal to 2/273) away from ry.," Instead, they converge toward a radius which is located a finite distance (equal to $2H/3$ ) away from $r_p$."
+" Since the potential is more and more localized around, at large m. coupling between the perturbation and the dise response is lost. and the contribution to the total torque becomes negligible."," Since the potential is more and more localized around $r_p$ at large $m$, coupling between the perturbation and the disc response is lost, and the contribution to the total torque becomes negligible."
+ Vhis is the torque cutoll first discussed by CGoldreich. Tremaine (1980)., This is the torque cutoff first discussed by Goldreich Tremaine (1980).
+ The WIXD analysis above indicates that for values of m up to about kf. the location of the turning points does not depend on whether there is a magnetic field or not.," The WKB analysis above indicates that for values of $m$ up to about $r/H$, the location of the turning points does not depend on whether there is a magnetic field or not."
+ In both cases they colnneide with the nominal Lindblad: resonances., In both cases they ncide with the nominal Lindblad resonances.
+ However. for larger values of m. there are dillerences between a magnetized and a nonmagnetized disc. as the analysis belowshows.," However, for larger values of $m$, there are differences between a magnetized and a nonmagnetized disc, as the analysis belowshows."
+ We now calculate the location of the turning points., We now calculate the location of the turning points.
+ There are no singularities in the coellicientsof equation (26)) at r=rg., There are no singularities in the coefficientsof equation \ref{EQ2}) ) at $r=r_L$.
+" This is similar to the nonmagnetic case. where only the second.order cifferential equation for WY, has (removable) singularities there (Naravan. Croldreich Goodman 1987)."," This is similar to the nonmagnetic case, where only the second–order differential equation for $W'_{m}$ has (removable) singularities there (Narayan, Goldreich Goodman 1987)."
+ ALL the physical quantities are regular at these locations., All the physical quantities are regular at these locations.
+ To identifv the turning points. we eliminate the first.order derivative term in this equation.," To identify the turning points, we eliminate the first–order derivative term in this equation."
+ We define c=rrp)£rg. re being either the inner or outer nominal Lindblad resonance. and introduce the new dependent variable: so that the homogeneous equation associated. with equation (26)) can be written under the exact form: where: We now expand K around ap—0.," We define $x \equiv (r-r_L)/r_L$, $r_L$ being either the inner or outer nominal Lindblad resonance, and introduce the new dependent variable: so that the homogeneous equation associated with equation \ref{EQ2}) ) can be written under the exact form: where: We now expand ${\cal K}$ around $x=0$."
+ Abr=rp. s?=m?67.," At $r=r_L$, $\kappa^2 = m^2
+\sigma^2$."
+ We neglect here the departure [romKeplerian rotation (which is second order in 4/7) so that &=O., We neglect here the departure fromKeplerian rotation (which is second order in $H/r$ ) so that $\kappa = \Omega$.
+ Then σον)N/m. with Y=1153emrí/2 and εξ| lor 1 depending on whether we consider the outer or inner resonance. respectively.," Then $\sigma = \epsilon \Omega(r_L) X /m$, with $X=1 + 3 \epsilon m x /2$ and $\epsilon=+1$ or $-1$ depending on whether we consider the outer or inner resonance, respectively."
+ We assume 2 and h—οΓΩ)Hr constant.," We assume $\beta$ and $h \equiv
+c/(r \Omega) \sim H/r$ constant."
+ Note that for values of mlíh and For p~h. we cannot approximate e by its value at ig.," Note that for values of $m \sim 1/h$ and for $x \sim h$, we cannot approximate $\sigma$ by its value at $r_L$."
+ After some algebra. we find. to the first non zero order: where we have assumed mx1. since these are the values of interest for the tidal torque (see CET9).," After some algebra, we find, to the first non zero order: where we have assumed $m \gg 1$, since these are the values of interest for the tidal torque (see GT79)."
+" is valid even for m larger than or very large compared to L/h. and whatever the values of 64.55.64.d, and d? compatible with ;; and h being constant and provided they are not large compared to unity."," This expression is valid even for $m$ larger than or very large compared to $1/h$, and whatever the values of $b_1, b_2, c_1, d_1$ and $d_2$ compatible with $\beta$ and $h$ being constant and provided they are not large compared to unity."
+ Note that the expression above can also be seen as an expansion around the corotation radius. which corresponds to XN= 0.," Note that the expression above can also be seen as an expansion around the corotation radius, which corresponds to $X=0$ ."
+ The turning points are the radii where AC= 0., The turning points are the radii where ${\cal K}=0$ .
+ We solve this equation numerically and display the results in figure 1.. where we plot νοrH Owhere 4= ¢/O) as a function of m for the turning points located beyond corotation. and fort= 0.1. 1 and," We solve this equation numerically and display the results in figure \ref{fig1}, , where we plot $(r-r_p)/H$ (where $H \equiv c/\Omega$ ) as a function of $m$ for the turning points located beyond corotation, and for $\beta=0.1$ , 1 and"
+feld is in waves of some Xiid. these waves might dap. providing au alternative nechanisii for the ris (Polletal.2005)..,"field is in waves of some kind, these waves might damp, providing an alternative mechanism for the rims \citep{pohl05}. ."
+ Figure 23 slows a typical rini but it Is present in radio as weLo where svuchrotrou losses cannot possibly be a factor.," Figure \ref{tycho} shows a typical rim – but it is present in radio as well, where synchrotron losses cannot possibly be a factor."
+ It is possible that both mechanisius are required to explain some rimus., It is possible that both mechanisms are required to explain some rims.
+ An additional argueit for the amplification of magnetic fell comes from the variations secu iu periods of a few vears in X-ray features in C317.3-0.5 (Uchivamaetal.2007) aud Cas A (Patuaude&Fe-sen2007)., An additional argument for the amplification of magnetic field comes from the variations seen in periods of a few years in X-ray features in G347.3-0.5 \citep{uchiyama07} and Cas A \citep{patnaude07}.
+. Both brightenime and fading are seen., Both brightening and fading are seen.
+ [f the relevant timescales are those of particle acceleration aud svuchrotrou losses. in either case one obtains values of the iiagnetic field of order Lane. IToxcever. if the post-shock magnetic feld is hiehlv turbulent. oue expects such “twinklue” due to fluctuations even if the mean magnetic field is lower (Bykovetal.2008).," If the relevant timescales are those of particle acceleration and synchrotron losses, in either case one obtains values of the magnetic field of order 1 mG. However, if the post-shock magnetic field is highly turbulent, one expects such “twinkling” due to fluctuations even if the mean magnetic field is lower \citep{bykov08}."
+. The discovery of a remmant only about 100 vr old. from expansion between radio observations made in 1985 and N-ray observations iu 2007 (Revuoldsctal. 2008b).. opens a new window on the very carly development of a SNR. aud of the physics of nuprecedentedly fast SNR shocks.," The discovery of a remnant only about 100 yr old, from expansion between radio observations made in 1985 and X-ray observations in 2007 \citep{reynolds08b}, opens a new window on the very early development of a SNR, and of the physics of unprecedentedly fast SNR shocks."
+ At an asstuect distance of the Galactic Center. 8.5 kpe (the very high absorbing+ column. Wy-~6⋅«E107s cimiP7. miens. the distance cannot be much less). the mean expuusiou speed is about 11.000 kin s.," At an assumed distance of the Galactic Center, 8.5 kpc (the very high absorbing column, $N_H \sim 6 \times 10^{22}$ $^{-2}$, means the distance cannot be much less), the mean expansion speed is about 14,000 km $^{-1}$."
+ The integratedC» svuchlirotrou] spectruui las a rolloff frequency of fy=2.2 keV. one of the highest ever measured.," The integrated synchrotron spectrum has a rolloff frequency of $h\nu = 2.2$ keV, one of the highest ever measured."
+ A recent longer observation has revealed the preseuce of spectral lines from the radio-brieht N riu aud interior (Figure 1))., A recent longer observation has revealed the presence of spectral lines from the radio-bright N rim and interior (Figure \ref{g1.9spec}) ).
+ The figure shows clear Hines of heliumu-like states of $i. S. and Av from the ceutral region.," The figure shows clear lines of helium-like states of Si, S, and Ar from the central region."
+" Ca. while visible frou, the N riu. is absent in the center. — but a line at L1 keV. which we identify as due to !!Sc, is clearly present."," Ca, while visible from the N rim, is absent in the center – but a line at 4.1 keV, which we identify as due to $^{44}$Sc, is clearly present."
+ !Se is produced by electron capture, $^{44}$ Sc is produced by electron capture
+"event is assumed to have been at the level of the 1998 season then the height of the peak maximum above the previous local minimum is My,0.5 magnitudes.",event is assumed to have been at the level of the 1998 season then the height of the peak maximum above the previous local minimum is $M_{peak}\sim 0.5$ magnitudes.
+ We calculate the cumulative probability for the difference between peak maxima and the preceding local minima ( ΕΕ a, We calculate the cumulative probability for the difference between peak maxima and the preceding local minima $M_{height}$ ).
+ndMpicigni is a function of sampling rate for sharp peaks sparse samplings., $M_{height}$ is a function of sampling rate for sharp peaks and sparse samplings.
+ We have therefore used a sampling rate of1 point per week and integrated the probability over οas well as 5 and im As before we have computed. the function. for. cusps GPOMneigh< Mycok}). |ee Caustic crossings Μα” Mysaa)) DC caustic Crossings (Pap(Mncigne€ Adpean)).," We have therefore used a sampling rate of 1 point per week and integrated the probability over $v_{eff}$as well as $S$ and $\langle m\rangle$: As before we have computed the function for cusps $P_{M_{C}}(M_{height}<M_{peak})$ ), $+ve$ caustic crossings $P_{M_{+}}(M_{height}<M_{peak})$ ), and $-ve$ caustic crossings $P_{M_{-}}(M_{height}<M_{peak})$ )."
+ Fig., Fig.
+ ui mMdisplavs the resulting Curves (assuming no systematic in the source size estimate)., \ref{max_height} displays the resulting curves (assuming no systematic uncertainty in the source size estimate).
+ ελΑνω~0.5. magnitudes is typical if he event is due to a cusp. and is ruled out at the evel if the event is. |re caustic crossing.," $\Delta M_{peak}\sim 0.5$ magnitudes is typical if the event is due to a cusp, and is ruled out at the level if the event is $+ve$ caustic crossing."
+ Lf the event is a re caustic crossing the results are inconclusive., If the event is a $-ve$ caustic crossing the results are inconclusive.
+ Ao very similar conclusion is reached if the source size has been uncderestimated by a factor of 2., A very similar conclusion is reached if the source size has been underestimated by a factor of 2.
+ The results are summarised in column 6 of Tab. 1, The results are summarised in column 6 of Tab. \ref{trigger_probs}.
+ WAVTAOO describes ai general function to determine how long one should wait (27) for a LAE following a hypothetical observed. light-curve derivative (2)., WWTA00 describes a general function to determine how long one should wait $P$ ) for a HME following a hypothetical observed light-curve derivative $T$ ).
+" Ifa light-curve derivative Zin,cMb. ds observed. predictions FOR,€x) can » made about lortheoming events specificallv for the current data using the sampling rate identical to that of the observations."," If a light-curve derivative $T_{obs}\pm\Delta T_{obs}$ is observed, predictions $F(P|T_{obs}\pm\Delta T_{obs})$ can be made about forthcoming events specifically for the current data using the sampling rate identical to that of the observations."
+ On the 19th of June 1999 monitoring from the OGLE collaboration (OGLE web page) showed image € rising at a rate of 1.21-1.78 nmiags/vear., On the 19th of June 1999 monitoring from the OGLE collaboration (OGLE web page) showed image C rising at a rate of 1.21-1.78 mags/year.
+ Since at this level 2 is not sensitive to the sampling rate for sampling spacings smaller than two weeks. we used observations on the 10th June. = June and. Ist July.," Since at this level $P$ is not sensitive to the sampling rate for sampling spacings smaller than two weeks, we used observations on the 10th June, 19th June and 1st July."
+ For the present calculation. the triggering function discussed. in. WW'EAO0 is slightly modified., For the present calculation the triggering function discussed in WWTA00 is slightly modified.
+ Rather than searching for peaks following a derivative larger than some hypothetical value. we search for peaks following light-curve derivatives in the measured range.," Rather than searching for peaks following a derivative larger than some hypothetical value, we search for peaks following light-curve derivatives in the measured range."
+ Thus since the algorithm steps along the curve in steps equal to the sampling rate. certain events (in. particular DC Caustic Crossings) are missed if the derivative jumps from below the lower bound o above the upper bound during one sampling interval.," Thus since the algorithm steps along the curve in steps equal to the sampling rate, certain events (in particular $+ve$ caustic crossings) are missed if the derivative jumps from below the lower bound to above the upper bound during one sampling interval."
+" ""Therefore the results describe the relative likelv-hood. of observing the dillerent. tvpes of event. following a current ieht-curve derivative.", Therefore the results describe the relative likely-hood of observing the different types of event following a current light-curve derivative.
+ As a consequence of the triggering algorithm searching for derivatives along the light-curve in one direction. then the mocdel lisht-curve derivative at the volt of the trigger is svstematically biased toward 2 AY.," As a consequence of the triggering algorithm searching for derivatives along the light-curve in one direction, then the model light-curve derivative at the point of the trigger is systematically biased toward $T-\Delta T$ ."
+ P ds therefore an overestimate., $P$ is therefore an overestimate.
+ This bias is minimised by reduced. photometric error., This bias is minimised by reduced photometric error.
+ lig., Fig.
+ 5 shows the resulting triggering functions £j," \ref{obs_trigger} shows the resulting triggering functions $F_{+}$ ,"
+Large Magellanic Cloud (LAIC). whose distance is taken as HO dpe. aud au error of 10 kpe for a couservative estimation is adapted (Bradley 2007).,"Large Magellanic Cloud (LMC), whose distance is taken as 50 kpc, and an error of 10 kpc for a conservative estimation is adapted (Bradley 2007)."
+ We sununarize the properties of all the 27 RPPs aud 2| PWNe in Tables 1 and 2. respectively.," We summarize the properties of all the 27 RPPs and 24 PWNe in Tables 1 and 2, respectively."
+ In Table 2 ouly 22 PWNe are associated with the RPPs listed in Table l., In Table 2 only 22 PWNe are associated with the RPPs listed in Table 1.
+ The other two pulsars are excluded. from Table 1: (1) Camilo et al. (, The other two pulsars are excluded from Table 1: (1) Camilo et al. (
+2001) suggested that PSR 5857 is oo faint to be resolved frou its background PWN in theChandra Observation: (2) Hoessels et al. (,2004) suggested that PSR $-$ 5857 is too faint to be resolved from its background PWN in the Observation; (2) Hessels et al. (
+2001) found that he spectrum of PSR J2021|3651 canbe ft by a DD model and is thus dominated by thermal components.,2004) found that the spectrum of PSR J2021+3651 can be fit by a BB model and is thus dominated by thermal components.
+ Ou the other haud. 5 out of the 27 RPP samples in Table 1 are not isted in Table 2 for PWNe. because the following 5 pulsars rave no PWN reported: 28. BOGSG|11. DOS23|26. D0950|08 and 52 (Ovechuan Tepedoeleuliosllu 2001. Becker et al.," On the other hand, 5 out of the 27 RPP samples in Table 1 are not listed in Table 2 for PWNe, because the following 5 pulsars have no PWN reported: $-$ 28, B0656+14, B0823+26, B0950+08 and $-$ 52 (Öggelman Tepedelenlioǧllu 2004, Becker et al."
+ 2001. De Luca et al.," 2004, De Luca et al."
+ 2005)., 2005).
+ The pulsar plotou indices are distributed iu a range of LsZSos (as shown in Fig 1))., The pulsar photon indices are distributed in a range of $1\la\Gamma_{\rm psr}\la3$ (as shown in Fig \ref{fig:Lxvs.Gamma}) ).
+ It should be noted that a Dasignificant fraction. about 154 (L out of 27). of the sources lave soft spectra of Ty72. which may raise problems for current models as discussed later in &1.," It should be noted that a significant fraction, about $\sim15\%$ (4 out of 27), of the sources have soft spectra of $\Gamma_{\rm psr}>2$, which may raise problems for current models as discussed later in 4."
+ The photon indices of the PWNe span a narrower range (Fig 1)), The photon indices of the PWNe span a narrower range (Fig \ref{fig:Lxvs.Gamma}) ).
+ As discussed. below (81). this is consistent. with tle pulsar wiud model.," As discussed below 4), this is consistent with the pulsar wind model."
+ We investigate below the correlations between the X-ray enission properties of RPPs aud PWNe. aud between the Cluission properties aud the pulsar rotational parameters.," We investigate below the correlations between the X-ray emission properties of RPPs and PWNe, and between the emission properties and the pulsar rotational parameters."
+ The rotation parameters include the period P. the period derivatives P. and some derived parameters. ce. the naenetic field B=3.3«101?(ppi ?. the characteristic age T=P 2P. aud the spin down power E-VpuIPIPS. where a typical moment of 7=10οen? is assuned.," The rotation parameters include the period $P$, the period derivatives $\dot{P}$, and some derived parameters, e.g., the magnetic field $B=3.3\times10^{19}(P\dot{P})^{1/2}$ G, the characteristic age $\tau=P/2\dot{P}$ , and the spin down power $\dot{E}=4\pi^2I\dot{P}/P^3$, where a typical moment of $I=10^{45}\rm g\;cm^{2}$ is assumed."
+ We rave taken the values of P aud P from the pulsar catalog wv Manchester et al., We have taken the values of $P$ and $\dot{P}$ from the pulsar catalog by Manchester et al.
+"(200510, Tn order to evaluate the significance level of the correlations of the twoparameters concerned we calculate he widely used Pearson correlation coefficient (r0). the Spearman rank correlation coefficient (74). and the two-sided significance level (po) of the Spearman rauk test."," In order to evaluate the significance level of the correlations of the twoparameters concerned, we calculate the widely used Pearson correlation coefficient $r$ ), the Spearman rank correlation coefficient $r_{\rm s}$ ), and the two-sided significance level $p_{\rm s}$ ) of the Spearman rank test."
+ The results are listed in Table 3., The results are listed in Table 3.
+ Iu addition to the correlation tests. we also perform a linear fit usine the least square method (LSM) to the relevant relations of the parameter pairs.," In addition to the correlation tests, we also perform a linear fit using the least square method (LSM) to the relevant relations of the parameter pairs."
+ Since the fitting results are usually dominated by a few data poiuts with much sinaller observational errors than the others. we also perform a linear fit without the observational errors for comparison.," Since the fitting results are usually dominated by a few data points with much smaller observational errors than the others, we also perform a linear fit without the observational errors for comparison."
+ The fitting results are all listed in Table 3. and shown iu the relevant figures.," The fitting results are all listed in Table 3, and shown in the relevant figures."
+ In the following we preseut the results in details., In the following we present the results in details.
+ We study the RPP emission. first., We study the RPP emission first.
+ Strong correlatious appear between the X-ray Inninosities of pulsars (Lx iu) and the pulsar rotational parameters (sce Table 3 aud Figs 2 aud 3)).," Strong correlations appear between the X-ray luminosities of pulsars $L_{\rm X, psr}$ ) and the pulsar rotational parameters (see Table 3 and Figs \ref{fig_le} and \ref{fig_Lx_psr}) )."
+" First. Lx is negatively correlated with 7 and positively correlated with E. which ave supported by the Spearman tests: r=Sl and p,«0.0001 between ExNoτωνpst and 7: and rr,O82 and p,«0.0001 between Ex and E."," First, $L_{\rm X,psr}$ is negatively correlated with $\tau$ and positively correlated with $\dot{E}$, which are supported by the Spearman tests: $r_{\rm s} = -0.81$ and $p_{\rm s}<0.0001$ between $L_{\rm X,psr}$ and $\tau$ ; and $r_{\rm s}=0.82$ and $p_{\rm s}<0.0001$ between $L_{\rm X,psr}$ and $\dot{E}$ ."
+ We also note that there are some hints of the correlation hold for Lx4: vs P aud P separately. with the relevant Spearman rank correlation cocticicuts of 0.06 and 0.69 respectively aud both significance levels <0.001.," We also note that there are some hints of the correlation hold for $L_{\rm X,psr}$ vs $P$ and $\dot{P}$ separately, with the relevant Spearman rank correlation coefficients of $-0.66$ and $0.69$ respectively and both significance levels $<0.001$ ."
+ The correlations between Lxpar vs P aud P will disappear when the sample includes both the AISPs aud the normal RPPs. just as shown by Posseuti et al. (," The correlations between $L_{\rm X,psr}$ vs $P$ and $\dot{P}$ will disappear when the sample includes both the MSPs and the normal RPPs, just as shown by Possenti et al. ("
+2002).,2002).
+ Despitethe Pearson aud Spearman correlation cocfücient support the existence of a correlation between Lx aud E (or 7). a simple linear fit to the logaritlin of the data poiuts with the observational errors included does not produce a statistically acceptable model.," Despitethe Pearson and Spearman correlation coefficient support the existence of a correlation between $L_{\rm X,psr}$ and $\dot{E}$ (or $\tau$ ), a simple linear fit to the logarithm of the data points with the observational errors included does not produce a statistically acceptable model."
+ Iu fact. it results (here Lx4 aud E ave in units of ere st and τ in vears: see also Figs 2 and 3).," In fact, it results (here $L_{\rm X,psr}$ and $\dot{E}$ are in units of erg $^{-1}$ and $\tau$ in years; see also Figs 2 and 3)."
+ Tere aud elsewhere in this paper. the uncertainties ou the linear fits are reported at confidence level.," Here and elsewhere in this paper, the uncertainties on the linear fits are reported at confidence level."
+ Previous authors (notably Posseuti et al., Previous authors (notably Possenti et al.
+ 2002) also noticed that a large scatter in the plot preveuts to obtain an acceptable fit of the data with a simple power law dependence of Lxpar on E.," 2002) also noticed that a large scatter in the plot prevents to obtain an acceptable fit of the data with a simple power law dependence of $L_{\rm X,psr}$ on $\dot{E}$."
+ ITeuce this relation must only be seeu as an ορΊσα average trend aud not suitable for predicting the huuinositv of auy specific source., Hence this relation must only be seen as an empirical average trend and not suitable for predicting the luminosity of any specific source.
+ We have also explored a linear fit which does not account for the uncertainties ou the values of Ly4.," We have also explored a linear fit which does not account for the uncertainties on the values of $L_{\rm X,psr}$."
+ It turus out (see Fies 2. and 3))., It turns out (see Figs \ref{fig_le} and \ref{fig_Lx_psr}) ).
+ A comparison between the curreut ESjy vorsus E velation aud the previous studies is slow in Fig 2..," A comparison between the current $L_{\rm
+X,psr}$ versus $\dot{E}$ relation and the previous studies is shown in Fig \ref{fig_le}."
+ It cau be seen that the relation we obtain above is close to the one between the pulsed X-ray cussion aud the spin down power in Cheng et al. (, It can be seen that the relation we obtain above is close to the one between the pulsed X-ray emission and the spin down power in Cheng et al. (
+200D). which iudicates that most of the non-thermal N-ray cussion frou a pulsar is pulsed.,"2004), which indicates that most of the non-thermal X-ray emission from a pulsar is pulsed."
+ As already doue by Posseuti et al. (, As already done by Possenti et al. (
+2002) using a saiple including also the MSPs (but not discutaneling PWN frou RPP euüssou). we also try to fit Lx with aPr.,"2002) using a sample including also the MSPs (but not disentangling PWN from RPP emission), we also try to fit $L_{\rm X}$ with $aP^{b}\dot{P}^c$."
+ This eives the relation The nominal result of this (still statistically unacceptable) fit would sugeest a preferred dependence of on L/P: however we note that. accounting for the Lxuncertainties ou the parameters.the simpler dependeuce on E (recovered in the work of Possenti ct al.," This gives the relation The nominal result of this (still statistically unacceptable) fit would suggest a preferred dependence of $L_{\rm X,psr}$ on $\dot{E}/P$: however we note that, accounting for the uncertainties on the parameters,the simpler dependence on $\dot{E}$ (recovered in the work of Possenti et al."
+ 2002) is also viable., 2002) is also viable.
+ We also study the pulsar spectral properties. aud check if there is anv correlation between the pulsar spectral iudex aud the pulsar rotational parameters.," We also study the pulsar spectral properties, and check if there is any correlation between the pulsar spectral index $\Gamma_{\rm psr}$ and the pulsar rotational parameters."
+ Tuspectionof DasFig. L. ," Inspectionof Fig. \ref{fig_Gamma_psr}, ,"
+may indicate the occurrence of a positive correlation of Ty. with P aud τ aud of a uegativecorrelation ofD; with Paud E., may indicate the occurrence of a positive correlation of $\Gamma_{\rm psr}$ with $P$ and $\tau$ and of a negativecorrelation of$\Gamma_{\rm psr}$ with $\dot{P}$and $\dot{E}$ .
+ However. a umnerical test indicates that all these correlations are too weak (the Spearman coefficients [rj are all S 0.60. see Table 3) for drawing anv frm conclusion with the available data.," However, a numerical test indicates that all these correlations are too weak (the Spearman coefficients $|r_{\rm s}|$ are all $\lesssim0.60$ see Table 3) for drawing any firm conclusion with the available data."
+planetary nebulae (PN) in M31 of LsLO?ereslo,"planetary nebulae (PN) in M31 of $4\times 10^{35}\,{\rm erg}\,{\rm s}^{-1}$."
+ For conrparisou. the IIo. luminosity for regions ionized by an O5 star is 6<LOeres+ (this corresponds tfo. the approximate limut of an regiou ionized by a single star or multiple stars). whereas for a BO.S star it is Έν10oresLO," For comparison, the $\alpha$ luminosity for regions ionized by an O5 star is $6\times 10^{37}\,{\rm erg}\,{\rm s}^{-1}$ (this corresponds to the approximate limit of an region ionized by a single star or multiple stars), whereas for a B0.5 star it is $1\times 10^{35}\,{\rm
+erg}\,{\rm s}^{-1}$."
+ The median Wa luminosities of NGC 598 and NGC 6207. the closest galaxies in our sample. are 4 and 5«LO?eres+. respectively. or close to the limit for PN. and we have thus excluded these two galaxies from our statistical analysis.," The median $\alpha$ luminosities of NGC 598 and NGC 6207, the closest galaxies in our sample, are 4 and $5\times
+10^{35}\,{\rm erg}\,{\rm s}^{-1}$, respectively, or close to the limit for PN, and we have thus excluded these two galaxies from our statistical analysis."
+ For the remaining galaxies. the Ia huninosities of the detected regions are always above the limit of PN. so they are likely to correspoucd to real roeglous.," For the remaining galaxies, the $\alpha$ luminosities of the detected regions are always above the limit of PN, so they are likely to correspond to real regions."
+ IE aualvzed the properties of disk regions of a siuuple of nearby spiral aud irregular galaxies., KEH analyzed the properties of disk regions of a sample of nearby spiral and irregular galaxies.
+ They fouud that the muuber of regions (normalized to the total D-baud huuinosity) increases by a factor of 6 frou IDIubble type Sb to Sc. aud by another from Sc to imregulars.," They found that the number of regions (normalized to the total $B$ -band luminosity) increases by a factor of 6 from Hubble type Sb to Sc, and by another from Sc to irregulars."
+ KEII observed a simular behavior with the number of regions per unit area. and no significant differences between the regious in the immer aud outer halves of the disk of a eiven galaxy.," KEH observed a similar behavior with the number of regions per unit area, and no significant differences between the regions in the inner and outer halves of the disk of a given galaxy."
+ From Table 2. if can be seen that within a eiven distance biu. there is a large dispersion in the umber of regions detected iu the circuunmuclear regious of ealaxies.," From Table 2, it can be seen that within a given distance bin, there is a large dispersion in the number of regions detected in the circumnuclear regions of galaxies."
+ We cau deteriune whether there is a relation between the mmuber of ceutral regions per unit area iu kpe? (so the distance effect ou the region covered by the images can be attenuated) aud the morphological tvpe (ve use the types T eiven in Table 1) for all ealaxies., We can determine whether there is a relation between the number of central regions per unit area in $^2$ (so the distance effect on the region covered by the images can be attenuated) and the morphological type (we use the types $T$ given in Table 1) for all galaxies.
+ The comparison is presented in the upper left panel of Figure lL. where galaxies have been divided iuto three iuorphologicaleroups: Z20 3.F=1 5and T>5.," The comparison is presented in the upper left panel of Figure 1, where galaxies have been divided into three morphological groups: $T = 0-3$, $T=4-5$ and $T>5$."
+ We oulv iuclude galaxies with distances 42Mpc. aud plot ealaxies with oulv uuclear cussion (see notes of Table 2) as having oue region detected.," We only include galaxies with distances $d \ge 2\,{\rm Mpc}$, and plot galaxies with only nuclear emission (see notes of Table 2) as having one region detected."
+ We also indicate the value of the median of the distribution., We also indicate the value of the median of the distribution.
+ This diagram does not show the tendency for later-type galaxies to lave more circrunnuuclear reeious per unit area than the earlier types. found by INET for regions im the disks of ealaxies.," This diagram does not show the tendency for later-type galaxies to have more circumnuclear regions per unit area than the earlier types, found by KEH for regions in the disks of galaxies."
+ Iu the upper right panel of Figure l1. we compare a related quantity. that is the wmuber of regions per nuit ucar-infrared Zf-band hpuuinositv. which represeuts he stellar mass in the region covered by the images.," In the upper right panel of Figure 1, we compare a related quantity, that is the number of regions per unit near-infrared $H$ -band luminosity, which represents the stellar mass in the region covered by the images."
+ From his figure it can be seen that although the values of the uecdian only rise slowly from carly to late types. the late-ype galaxies show a tailin the distribution reaching higher iunubers of eireummuclear reeious per unit stellar mass han earlicr-types.," From this figure it can be seen that although the values of the median only rise slowly from early to late types, the late-type galaxies show a tail in the distribution reaching higher numbers of circumnuclear regions per unit stellar mass than earlier-types."
+ This is caused by the fact that carly uorphological type galaxies tend to have more huuinous (and iore massive) bulges., This is caused by the fact that early morphological type galaxies tend to have more luminous (and more massive) bulges.
+ This in good agreement with he elobal behavior of disk regious iu ucarby galaxies (IKEIT)., This in good agreement with the global behavior of disk regions in nearby galaxies (KEH).
+ A similar comparison is carried out using the bar class (A. AB and D) to determine whether the presence of a bar in a galaxy influences the uumuber of cireumnmmclear regions.," A similar comparison is carried out using the bar class (A, AB and B) to determine whether the presence of a bar in a galaxy influences the number of circumnuclear regions."
+ The comparison is shown in the lower panucls of Figure 1 for the wmmber of regions per unit area. and per unit stellar mass.," The comparison is shown in the lower panels of Figure 1 for the number of regions per unit area, and per unit stellar mass."
+ There is no tendency for barred ealaxies to show more regious per unit stellar mass or uit area., There is no tendency for barred galaxies to show more regions per unit stellar mass or unit area.
+ The accepted view is that the global SF activity in ealaxies decreases systematically from late-type to carlytype spirals. a conclusion mainly based ou the integrated equivalent widths of Ta (e.g.. Nennicutt Invent 1983 aud Ikeunicutt 1998a for a review].," 	 The accepted view is that the global SF activity in galaxies decreases systematically from late-type to early-type spirals, a conclusion mainly based on the integrated equivalent widths of $\alpha$ (e.g., Kennicutt Kent 1983 and Kennicutt 1998a for a review)."
+ This view is also supported by the fact that the luminosities of the brightest regions in the disks of early-type spiral galaxies are on average lower than those of later types (IKEIT)., This view is also supported by the fact that the luminosities of the brightest regions in the disks of early-type spiral galaxies are on average lower than those of later types (KEH).
+ Moreover. Caldwell ct al. (," Moreover, Caldwell et al. ("
+1991) reported that in Sa galaxies there are no regions with Πα luuinositics greater than 107eresi1. although this view has been challenged by Devereux Tameed (1997) aud Hameed Devercux (1999. and references therein).,"1991) reported that in Sa galaxies there are no regions with $\alpha$ luminosities greater than $10^{39}\,{\rm erg}\,{\rm s}^{-1}$, although this view has been challenged by Devereux Hameed (1997) and Hameed Devereux (1999, and references therein)."
+ The SF activity of galaxy nuclei differs from the elobal behavior as found by Ilo. Filippeuko. Sargent (1997a) who from au optical spectroscopic survey showed that early-type (80She} ealaxies have nuclei with higher Ho Inuinosities than late-type (ScI0) galaxies. with an increase iu the median of about a factor of 9. for (starburst) tvpe galaxies.," The SF activity of galaxy nuclei differs from the global behavior as found by Ho, Filippenko, Sargent (1997a) who from an optical spectroscopic survey showed that early-type (S0–Sbc) galaxies have nuclei with higher $\alpha$ luminosities than late-type (Sc–I0) galaxies, with an increase in the median of about a factor of 9, for (starburst) type galaxies."
+ When we compare the median dDunuinosities of the circununuclear regions with the morphological tvpe for galaxies within distauce bius (see Table 2). we find no clear evideuce for later-type galaxies to show brighter eircuumnmuclear regions.," When we compare the median luminosities of the circumnuclear regions with the morphological type for galaxies within distance bins (see Table 2), we find no clear evidence for later-type galaxies to show brighter circumnuclear regions."
+ For mstauce. within the 4=17 Mpc distance bin. the median Ha huninosity of regious iu the earltype spirals IC 750 and NGC 1192 are above the average value within this bin (although we note that IC 750 is an interacting galaxy).," For instance, within the $d=17\,$ Mpc distance bin, the median $\alpha$ luminosity of regions in the early-type spirals IC 750 and NGC 4192 are above the average value within this bin (although we note that IC 750 is an interacting galaxy)."
+ A similar behavior is found within the MMpe distance bin. where the early-type galaxies NGC 3593 and NGC L826 host relatively bright median regions.," A similar behavior is found within the Mpc distance bin, where the early-type galaxies NGC 3593 and NGC 4826 host relatively bright median regions."
+ Although our ο of early-type spirals (that is. Sa aud Sab) is small. there is no clear trend for carlicr-type galaxies to show fainter circiunuuclear roeglous.," Although our number of early-type spirals (that is, Sa and Sab) is small, there is no clear trend for earlier-type galaxies to show fainter circumnuclear regions."
+ The huuinositv of the brightest cimemunuclear region can be used as an estinate of the age of the SE xocess (at least if the SE. occurred. iu an instantaneous must). although this quautitv is also depeucdent upon the uaximumn number of ioniziug stars per cluster. and upon he upper mass eutoff of the IME (see Feinstein 1997: Ocv Clark 1998).," The luminosity of the brightest circumnuclear region can be used as an estimate of the age of the SF process (at least if the SF occurred in an instantaneous burst), although this quantity is also dependent upon the maximum number of ionizing stars per cluster, and upon the upper mass cutoff of the IMF (see Feinstein 1997; Oey Clark 1998)."
+ An advantage of using the luminosity of he brightest region over the mean hnuuinositv is that is not scusitive o the niuinmuni size imposed for the detection of an region. although the blending effects will be more nuportaut for the more distaut galaxies.," An advantage of using the luminosity of the brightest region over the mean luminosity is that is not sensitive to the minimum size imposed for the detection of an region, although the blending effects will be more important for the more distant galaxies."
+ The left panel of Figure 2 shows the logarithiu of the Πα uuimnositv of the brightest eireumnuuclear region as a Muction of morphological type (erouped as in Figure 1) or galaxies with distances d72A\Ipe.," The left panel of Figure 2 shows the logarithm of the $\alpha$ luminosity of the brightest circumnuclear region as a function of morphological type (grouped as in Figure 1) for galaxies with distances $d> 2\,{\rm Mpc}$."
+ For a given group of morphological types. there is a spread of approximately wo orders of magnitude in the πιο ofthe brightest region. where part of the scatter is probably caused by lending effects in the more distaut galaxies.," For a given group of morphological types, there is a spread of approximately two orders of magnitude in the luminosity of the brightest region, where part of the scatter is probably caused by blending effects in the more distant galaxies."
+ Although the uecdian values of the Πα liminosity for the brightest region are simular for the earl-types (£=0 3) and late vpe (T> 5). the distribution for the late-tvpe galaxies shows a tail reaching fainter huuinosities.," Although the median values of the $\alpha$ luminosity for the brightest region are similar for the early-types $T=0-3$ ) and late type $T>5$ ), the distribution for the late-type galaxies shows a tail reaching fainter luminosities."
+ As the, As the
+"In this paper. we have presented multi-epoch. VLBA observations of EC 95, à young stellar object located in the SVS 4 sub-cluster of the Serpens molecular core.","In this paper, we have presented multi-epoch VLBA observations of EC 95, a young stellar object located in the SVS 4 sub-cluster of the Serpens molecular core."
+ Our data demonstrate that EC 95 is in facta üght binary system with a separation of about 15 mas., Our data demonstrate that EC 95 is in fact a tight binary system with a separation of about 15 mas.
+" The primary (EC 95a) appears to be a 4-5 M. intermediate-mass Herbig AeBe protostar, whereas the secondary (EC 95b) is most likely a low-mass T Tauri."," The primary (EC 95a) appears to be a 4–5 $M_\odot$ intermediate-mass Herbig AeBe protostar, whereas the secondary (EC 95b) is most likely a low-mass T Tauri."
+" At radio wavelengths, the secondary is on average brighter than the primary."," At radio wavelengths, the secondary is on average brighter than the primary."
+" It is, therefore, also likely to dominate the fairly bright X-ray emission [rom the system."," It is, therefore, also likely to dominate the fairly bright X-ray emission from the system."
+" The primary, on the other hand, contributes most of the infrared and of the bolometric luminosity of EC 95."," The primary, on the other hand, contributes most of the infrared and of the bolometric luminosity of EC 95."
+ This might naturally explain why the exünction to EC 95 based on infrared observations was systematically found to be much larger than the exünction based on X-ray data., This might naturally explain why the extinction to EC 95 based on infrared observations was systematically found to be much larger than the extinction based on X-ray data.
+" Interestingly, both members of EC 95 appear to be non-thermal radio emitters."," Interestingly, both members of EC 95 appear to be non-thermal radio emitters."
+" While a low-mass T Tauri star such as EC 95b is expected to generate that type of emission, an intermediate-mass such as EC 95a is not."," While a low-mass T Tauri star such as EC 95b is expected to generate that type of emission, an intermediate-mass such as EC 95a is not."
+" We discussed several mechanisms that could explain the presence of non-thermal emission on EC 95a, and argue that the observed properties of EC 95a might be most readily explained if it possessed a corona powered by a thin, rotation-driven convective layer."," We discussed several mechanisms that could explain the presence of non-thermal emission on EC 95a, and argue that the observed properties of EC 95a might be most readily explained if it possessed a corona powered by a thin, rotation-driven convective layer."
+" The trigonometric parallax of EC 95 appears to be x = 2.41 + 0.02 mas, corresponding to a distance of 414.9!"" pe."," The trigonometric parallax of EC 95 appears to be $\pi$ = 2.41 $\pm$ 0.02 mas, corresponding to a distance of $^{+4.4}_{-4.3}$ pc."
+ We argue that this implies a distance to the Serpens core of 415 + 5 pc. and to the Serpens molecular cloud of 415 + 25 pe.," We argue that this implies a distance to the Serpens core of 415 $\pm$ 5 pc, and to the Serpens molecular cloud of 415 $\pm$ 25 pc."
+ This is significantly larger than previous distance estimates based on measurements of the exancuion sullered by stars in the direction of Serpens (d. — 260 pc)., This is significantly larger than previous distance estimates based on measurements of the extinction suffered by stars in the direction of Serpens $d$ $\sim$ 260 pc).
+ A possible explanation for this discrepancy is that these measurements correspond to the distance to clouds associated with the Aquila Rilt located in the foreground of the Serpens cloud., A possible explanation for this discrepancy is that these measurements correspond to the distance to clouds associated with the Aquila Rift located in the foreground of the Serpens cloud.
+" Since our observations resolve EC 95 into a binary, future radio monitoring of that system will allow us to measure the dynamical mass of that system."," Since our observations resolve EC 95 into a binary, future radio monitoring of that system will allow us to measure the dynamical mass of that system."
+" To our knowledge, this will be the first üme that a dynamical mass is obtained for a young Herbig AeBe system, and it will enable us to uniquely constrain theoretical pre-main sequence evolutionary for intermediate-mass stars."," To our knowledge, this will be the first time that a dynamical mass is obtained for a young Herbig AeBe system, and it will enable us to uniquely constrain theoretical pre-main sequence evolutionary for intermediate-mass stars."
+" The orbital period of the system appears to be 10—20 years, so observations in the next two decades will be required to obtain an accurate mass estimate."," The orbital period of the system appears to be 10–20 years, so observations in the next two decades will be required to obtain an accurate mass estimate."
+" To Serpens cloud is about 20 pc across on the plane of the sky, so it is likely to be also about 20 pc deep."," To Serpens cloud is about 20 pc across on the plane of the sky, so it is likely to be also about 20 pc deep."
+" Observations similar to those presented here for a sample of young stars distributed across the cloud would help determine the structure of the region, and would allow us to further constrain the mean distance to this important region of star-lormation."," Observations similar to those presented here for a sample of young stars distributed across the cloud would help determine the structure of the region, and would allow us to further constrain the mean distance to this important region of star-formation."
+The slowly rotatiug active Gas IIL-IV. eiant El Eri (IIR 1362. ΠΟ 27536) is considered as a candidate for beiug the descendant of a strongly magnetic Ap star ( Stepien 1993. Strassincier et al.,"The slowly rotating active G8 III-IV giant EK Eri (HR 1362, HD 27536) is considered as a candidate for being the descendant of a strongly magnetic Ap star ( Stepień 1993, Strassmeier et al."
+ 1999)., 1999).
+ The detection of a strong surface naguctic field. dominated by a poloidal component. is consistent with this hvpothesis (Aurierre et al.," The detection of a strong surface magnetic field, dominated by a poloidal component, is consistent with this hypothesis (Aurièrre et al."
+ 2008)., 2008).
+ EK En has now been surveyed photometrically during nore than 30 vears and its maenetic activity studied spectroscopically (Dall ct al., EK Eri has now been surveyed photometrically during more than 30 years and its magnetic activity studied spectroscopically (Dall et al.
+ 2010. hereafter DBSS1O).," 2010, hereafter DBSS10)."
+ These authors demonstrate that. while the photometric variation and possibly the perio may change from season O season. an average period of 308 + 2.5 d provides an acceptable explanation of the data.," These authors demonstrate that, while the photometric variation and possibly the period may change from season to season, an average period of 308 $\pm$ 2.5 d provides an acceptable explanation of the data."
+ DDSS10 report a photometric ephemeris based on this period. confi hat the main helt variatious are due to dark spots aud chtatively sugeest that the star is a cdipole-dominatecd oblique rotator viewed close to equator-on.," DBSS10 report a photometric ephemeris based on this period, confirm that the main light variations are due to dark spots and tentatively suggest that the star is a dipole-dominated oblique rotator viewed close to equator-on."
+ Based ou this Xeture they propose that the rotation period is twice the photometric period., Based on this picture they propose that the rotation period is twice the photometric period.
+ Following the initial magnetic observations presented yw Awurerre et al. (, Following the initial magnetic observations presented by Aurièrre et al. (
+2008). we have now monitored Els Eu throughout 1 seasons. correspouding to about 1 photometric periods (308.8 d.. according to DBSS10).,"2008), we have now monitored EK Eri throughout 4 seasons, corresponding to about 4 photometric periods $308.8$ d., according to DBSS10)."
+" Iu lis paper we preseut our new data and the associated analysis, which enables us to establish the rotational CLIO’ of EI Eri aud to model its surface magnetic opologv."," In this paper we present our new data and the associated analysis, which enables us to establish the rotational period of EK Eri and to model its surface magnetic topology."
+ The 28 observations of EI& Exi were obtained between 20 September 2007 and 21 March 2011 at the 212 telescope Bernard Lvot (TBL) of Pic du Midi: Observatorv., The 28 observations of EK Eri were obtained between 20 September 2007 and 21 March 2011 at the 2m telescope Bernard Lyot (TBL) of Pic du Midi Observatory.
+ We used the NARVAL spectropolarimeer. Which is à copy of he ESPaDOuS iustruneut at the Canada-France-Tawaii Telescope (Donati et al.," We used the NARVAL spectropolarimeter, which is a copy of the ESPaDOnS instrument at the Canada-France-Hawaii Telescope (Donati et al."
+ 20063)., 2006a).
+ NARVAL was used in »olarimnetiic uode wihn a spectra resolution of about 65000. in the same confguration and with the Sue xocedure as describes by Aurore et al. (," NARVAL was used in polarimetric mode with a spectral resolution of about 65000, in the same configuration and with the same procedure as described by Aurièrre et al. ("
+2008).,2008).
+ To complete the Zeeman analysis. least-squares deconvolution (LSD. Donati ct al," To complete the Zeeman analysis, least-squares deconvolution (LSD, Donati et al."
+ 1997) was applied to all observations., 1997) was applied to all observations.
+ We used a line mask calculated for an effecive telpcrature of 5250 EK. logg=35. and a inicroturbulence of 2.0 m |J. consistent with plysical paraneters reported by DBSSLO.," We used a line mask calculated for an effective temperature of $5250$ K, $\log g =3.5$, and a microturbulence of 2.0 km $^{-1}$, consistent with physical parameters reported by DBSS10."
+ Iu the preseut case this method enabled us to average about 11.000 spectral lines and to derive Stokes V. profiles with a S/N that improved by a factor of about 10 in comparison with that for single lues;," In the present case this method enabled us to average about 11,000 spectral lines and to derive Stokes $V$ profiles with a S/N that improved by a factor of about 40 in comparison with that for single lines."
+" We then computed the longitudinal magnetic field D, m Gauss. using the first-order moment method (Recs Semel 1979. Donati ο al."," We then computed the longitudinal magnetic field $B_{\ell}$ in Gauss, using the first-order moment method (Rees Semel 1979, Donati et al."
+ 1997)., 1997).
+ Spectra obtained with NARVAL enable us to ΠΕ obtain naenetic data. leasurements of activity ludicators as well as radial velocity (RV) micastrecuts. so that studying correlations between these quautities is straigitforward.," Spectra obtained with NARVAL enable us to simultaneously obtain magnetic data, measurements of activity indicators as well as radial velocity $RV$ ) measurements, so that studying correlations between these quantities is straightforward."
+ Tjc activity of the star was monitored by using lineactivity ludicators measured from the NARVAL spectra., The activity of the star was monitored by using line–activity indicators measured from the NARVAL spectra.
+ We couputed the 5-index (defined from the Moui Wilsou strvev. Duncan et al.," We computed the $S$ -index (defined from the Mount Wilson survey, Duncan et al."
+Therefore in a magnetar density changes. due to the spin-down. take place in a short interval of time.,"Therefore in a magnetar density changes, due to the spin-down, take place in a short interval of time."
+ Hence they can provide enough enerev lo trigger spontaneous nucleation {ο the quark phase. once (he central density of the star goes above (he critical density.," Hence they can provide enough energy to trigger spontaneous nucleation to the quark phase, once the central density of the star goes above the critical density."
+ The nucleation of quark droplets in the presence of strong magnetic fields essentially depends on the strength. of the magnetic field., The nucleation of quark droplets in the presence of strong magnetic fields essentially depends on the strength of the magnetic field.
+" The energy of a quantum particle changes signilicantly if the magnetic field is of the order or exceeds the critical value B,=m?c*/eh. where i» is (he mass of the particle."," The energy of a quantum particle changes significantly if the magnetic field is of the order or exceeds the critical value $B_{c}=m^{2}c^{3}/e\hbar $, where $m$ is the mass of the particle."
+ For electrons the value of the critical field is ο In order to calculate the value of the critical magnetic field for quarks we have to use the estimations for the quark masses.," For electrons the value of the critical field is $%
+B_{c}^{(e)}=4.4\times 10^{13 }$ G. In order to calculate the value of the critical magnetic field for quarks we have to use the estimations for the quark masses."
+ However. there are many uncertainües in the nunmerical values of (he quark masses. which play an essential role in (he calculation of the critical magnetic fied and which are generally poorly known.," However, there are many uncertainties in the numerical values of the quark masses, which play an essential role in the calculation of the critical magnetic field and which are generally poorly known."
+" In the case of ree quarks. the current quark mass is considered to be in the range of m,~5—LO MeV. Then the corresponding− critical−∙ magnetic⋅ field⋅ is⋅ Be’(qi~(1—2)ETx10EBY(c)zz(4.488)5LOY G.- For the current quark mass. this is the tvpical strength of the magnetic field at which the cvclotron lines begin to occur."," In the case of free quarks, the current quark mass is considered to be in the range of $m_{q}\sim 5-10$ MeV. Then the corresponding critical magnetic field is $B_{c}^{(q)}\sim (1-2)\times 10^{2}\times
+B_{c}^{(e)}\approx (4.4-8.8)\times 10^{15}$ G. For the current quark mass, this is the typical strength of the magnetic field at which the cyclotron lines begin to occur."
+ In this limit the evelotron quantum is of the order or greater (han its rest enerey., In this limit the cyclotron quantum is of the order or greater than its rest energy.
+ This is equivalent to the requirement that the de Broglie wavelength is of (he order or greater (han the Larmor radius of the particle in the magnetic field1995)., This is equivalent to the requirement that the de Broglie wavelength is of the order or greater than the Larmor radius of the particle in the magnetic field.
+. However. from the physical point of view a better description of the quark mass can be obtained by using instead the current quark mass (he effective or constitute quark mass. which takes into account the effects of the strong interactions. and which is of the order mg~100—300 MeV1993).. The critical magnetic field corresponding to the effective quark mass could be as high as BU~10!—1055 G. The study of the behavior of the surface and curvature terms in strong magnetic have. shown that both quantities diverge⋅⋅ [or⋅ B.(qi withtheenergy.curvature term diverging .fields much faster.," However, from the physical point of view a better description of the quark mass can be obtained by using instead the current quark mass the effective or constitute quark mass, which takes into account the effects of the strong interactions, and which is of the order $m_q\sim
+100-300$ MeV. The critical magnetic field corresponding to the effective quark mass could be as high as $B_{c}^{(q)}\sim 10^{17}-10^{18}$ G. The study of the behavior of the surface and curvature energy terms in strong magnetic fields have shown that both quantities diverge for }$, with the curvature term diverging much faster."
+ Consequently. in (he presence of strong magnetic fields (he rate of stable quark droplet formation per unit volume 1~T'exp(-W/T) tends to zero. 1—0. and therefore there cannot be any thermal nucleation of quark droplets in neutron star cores1996).," Consequently, in the presence of strong magnetic fields the rate of stable quark droplet formation per unit volume $%
+I\sim T^{4}\exp \left( -W/T\right) tends to zero, $I\rightarrow
+0$, and therefore there cannot be any thermal nucleation of quark droplets in neutron star cores."
+. The value of D inside à magnetar is not. known. and it is not certain if the interior magnetic field is much stronger than the surface field.," The value of $B$ inside a magnetar is not known, and it is not certain if the interior magnetic field is much stronger than the surface field."
+" If in the interior of the magnetar the value of the magnetic ⋅∙field can exceed the critical−− value 2,7(ο)~5xLOY 5¢G. which represents je nost conservative estimation of the critical magnetic field. corresponding to the [ree Aquark mass m.=5 MeV. then the rate of droplet lormation can be considerably reduced aud —1e (ranisitioni from neutron matter to quark matter could not take place."," If in the interior of the magnetar the value of the magnetic field can exceed the critical value $B_{c}^{(q)}\sim
+5\times 10^{15}$ G, which represents the most conservative estimation of the critical magnetic field, corresponding to the free quark mass $m=5$ MeV, then the rate of droplet formation can be considerably reduced and the transition from neutron matter to quark matter could not take place."
+ On the other hand. by taking into account that the physical mass of (he quark is effective mass.. the corresponding critical. magnetic∙∙ field− could∙ be of the(qi order," On the other hand, by taking into account that the physical mass of the quark is the effective mass, the corresponding critical magnetic field could be of the order"
+ On the other hand. by taking into account that the physical mass of (he quark is effective mass.. the corresponding critical. magnetic∙∙ field− could∙ be of the(qi order ," On the other hand, by taking into account that the physical mass of the quark is the effective mass, the corresponding critical magnetic field could be of the order"
+ On the other hand. by taking into account that the physical mass of (he quark is effective mass.. the corresponding critical. magnetic∙∙ field− could∙ be of the(qi order ," On the other hand, by taking into account that the physical mass of the quark is the effective mass, the corresponding critical magnetic field could be of the order"
+This system is very similar to PTFEB28.235. with similar variability levels and a similar light curve shape.,"This system is very similar to PTFEB28.235, with similar variability levels and a similar light curve shape."
+ The system has a GALEX far-UV detection at mpuy=21.0+0.1 (with no near-UV coverage)., The system has a GALEX far-UV detection at $\rm{m_{FUV}=21.0\pm0.1}$ (with no near-UV coverage).
+ The optical light in these systems is dominated by the M-dwarfs. and the large radius ratios between the components produce very small (millimag-level) primary. eclipses.," The optical light in these systems is dominated by the M-dwarfs, and the large radius ratios between the components produce very small (millimag-level) primary eclipses."
+ For these reasons. we were able to measure precise eclipse shapes only for the white-dwarf occultation events. and radial velocities for the M-dwarf components.," For these reasons, we were able to measure precise eclipse shapes only for the white-dwarf occultation events, and radial velocities for the M-dwarf components."
+ However. the low-resolution spectra allow us to estimate the mass and radius of the M-dwarf component of the system and thus estimate the properties ofthe white dwarf component.," However, the low-resolution spectra allow us to estimate the mass and radius of the M-dwarf component of the system and thus estimate the properties ofthe white dwarf component."
+ Finally. cadence photometry of the ingress and egress of the white-dwarf occultation events allowed us to set stringent. upper limits on the radius of the white dwarf in PTFEBI1.441.," Finally, high-cadence photometry of the ingress and egress of the white-dwarf occultation events allowed us to set stringent upper limits on the radius of the white dwarf in PTFEB11.441."
+ We model the low-resolution spectra of our targets as a combination of à M-dwarf and a white dwarf spectrum., We model the low-resolution spectra of our targets as a combination of a M-dwarf and a white dwarf spectrum.
+ Template M-dwarf spectra are taken from the HILIB stellar flux library (2). and have steps of a single spectral subclass., Template M-dwarf spectra are taken from the HILIB stellar flux library \citep{Pickles1998} and have steps of a single spectral subclass.
+ The white dwarf were kindly provided by Detlev Koester and are on a grid with 250K steps between 6000K and 10000K (and longer steps up to 20000K). and 0.25 dex steps between log g = 6.0 and log g = 9.5.," The white dwarf were kindly provided by Detlev Koester and are on a grid with 250K steps between 6000K and 10000K (and longer steps up to 20000K), and 0.25 dex steps between log g = 6.0 and log g = 9.5."
+ We simultaneously fit the white dwarf and M-dwarf spectra using a downhill simplex algorithm and a bootstrap (e.g. ?)) method to find the uncertainties in the fit., We simultaneously fit the white dwarf and M-dwarf spectra using a downhill simplex algorithm and a bootstrap (e.g. \citealt{Press1992}) ) method to find the uncertainties in the fit.
+ Figure 8 shows the fitted models and each component of the binaries with the model for the other component subtracted., Figure \ref{fig:specs} shows the fitted models and each component of the binaries with the model for the other component subtracted.
+ In each case subtraction of the template M-dwarf spectrum revealed a (high-noise) DA white dwarf spectrum with clear Balmer lines., In each case subtraction of the template M-dwarf spectrum revealed a (high-noise) DA white dwarf spectrum with clear Balmer lines.
+ PTFEB11.441 is well-fit by a simple combined white-dwarf and M-dwarf spectrum. suggesting that the third component in the system ts also an M-dwarf. and that the two M-dwarfs have similar spectral types.," PTFEB11.441 is well-fit by a simple combined white-dwarf and M-dwarf spectrum, suggesting that the third component in the system is also an M-dwarf, and that the two M-dwarfs have similar spectral types."
+ In the following sections. we use the measured spectral type for the M-dwarfs in PTFEBI1.441. assuming that it is the spectral type of the M-dwarf closest to the white dwarf.," In the following sections, we use the measured spectral type for the M-dwarfs in PTFEB11.441, assuming that it is the spectral type of the M-dwarf closest to the white dwarf."
+ The M-dwarf spectral types measured from the two-component fit to the system spectra closely agree with those estimated from infrared photometry of the systems (table 3))., The M-dwarf spectral types measured from the two-component fit to the system spectra closely agree with those estimated from infrared photometry of the systems (table \ref{tab:targets}) ).
+ This suggests that reddening is not a significant problem for the M-dwarf spectral types and thus mass and radius estimations., This suggests that reddening is not a significant problem for the M-dwarf spectral types and thus mass and radius estimations.
+ However. we note that reddening will affect the estimated white dwarf temperature more severely as its flux is mostly in the blue end of the spectrum.," However, we note that reddening will affect the estimated white dwarf temperature more severely as its flux is mostly in the blue end of the spectrum."
+ The estimated galactic extinction along the line of sight to PTFEBI1.441 is (?).. and will be lower for the target itself which is located at a distance of =400pc.," The estimated galactic extinction along the line of sight to PTFEB11.441 is E(B-V)=0.02 \citep{Schlegel1998}, and will be lower for the target itself which is located at a distance of $\approx400pc$."
+ The other two targets are located in regions of higher line-of-sight extinction. at E(B-Vo=1.5 and 1.1 mag for PTFEB28.235 and PTFEB28.852 respectively.," The other two targets are located in regions of higher line-of-sight extinction, at E(B-V)=1.5 and 1.1 mag for PTFEB28.235 and PTFEB28.852 respectively."
+ The actual extinction. for the two targets ts. however. likely to be much lower as they are located at distances of only «20θρο.," The actual extinction for the two targets is, however, likely to be much lower as they are located at distances of only $\approx$ 200pc."
+ This is supported by several different arguments: the very good match of the M-dwarf model spectrum to the observed spectrum: the match between the infrared-color-estimated spectral types and the spectrally-measured spectral types: and the GALEX detections of both objects. which require a small near and far-UV extinction.," This is supported by several different arguments: the very good match of the M-dwarf model spectrum to the observed spectrum; the match between the infrared-color-estimated spectral types and the spectrally-measured spectral types; and the GALEX detections of both objects, which require a small near and far-UV extinction."
+ We fitthe M-dwarf masses. radii and distances in the following manner: we first fit survey photometry of the system with the stellar SEDs listed in 2.. yielding a photometric spectral type and distance.," We fitthe M-dwarf masses, radii and distances in the following manner: we first fit survey photometry of the system with the stellar SEDs listed in \citet{Kraus2007a}, yielding a photometric spectral type and distance."
+ Despite the availability of other photometric data such as USNO-BI. we restrict our fits to 2MASS J. H and K colors that are unlikely to be affected by the light from the white dwarf.," Despite the availability of other photometric data such as USNO-B1, we restrict our fits to 2MASS J, H and K colors that are unlikely to be affected by the light from the white dwarf."
+ We assume a distance uncertainty (??).. and +1-2 spectral subclass uncertainty (2); the SED estimates for the spectral types are consistent. with those derived. from the two-component fits to our low-resolution spectroscopy.," We assume a distance uncertainty \citep{Kraus2007a, Law2010}, and $\pm$ 1-2 spectral subclass uncertainty \citep{Kraus2007a}; the SED estimates for the spectral types are consistent with those derived from the two-component fits to our low-resolution spectroscopy."
+ We also attempted to determine the M-dwarf spectral type from the Τον narrowband spectroscopic index (?).. but found implausibly high values of the index. suggesting possible contamination by white dwarf emission.," We also attempted to determine the M-dwarf spectral type from the TiO5 narrowband spectroscopic index \citep{Gizis1997}, but found implausibly high values of the index, suggesting possible contamination by white dwarf emission."
+ In the following sections we adopt the spectral types estimated from the two-component fits to the low-resolution spectra. as that method yields the lowest spectral type uncertainty.," In the following sections we adopt the spectral types estimated from the two-component fits to the low-resolution spectra, as that method yields the lowest spectral type uncertainty."
+ We estimate the M-dwarf masses from the spectral type vs. mass calibrations detailed in ?.. assuming they are on the main-sequence and have solar metallicity and age.," We estimate the M-dwarf masses from the spectral type vs. mass calibrations detailed in \citet{Delfosse2000}, , assuming they are on the main-sequence and have solar metallicity and age."
+ All three of our targets have a measured spectral type of M3. giving an estimated mass of 0.3540.05 M...," All three of our targets have a measured spectral type of M3, giving an estimated mass of $\pm$ 0.05 $M_{\odot}$."
+ We note that a small fraction of the systems measured in ? have much lower masses at the M3 spectral type. for reasons that are still unclear.," We note that a small fraction of the systems measured in \citet{Delfosse2000} have much lower masses at the M3 spectral type, for reasons that are still unclear."
+ To estimate the masses in an alternate manner. we estimated the M-dwarf effective temperatures from the relation described in ?..," To estimate the masses in an alternate manner, we estimated the M-dwarf effective temperatures from the relation described in \citet{Luhman1999}."
+ We then combined those estimates with the SGYr isochrones of solar-metallicity stars in ? to estimate the stellar masses., We then combined those estimates with the 5GYr isochrones of solar-metallicity stars in \citet{Baraffe1998} to estimate the stellar masses.
+" These relations also predict a much lowermass for our M-dwarf targets. of 0.23, "," These relations also predict a much lowermass for our M-dwarf targets, of $^{+0.06}_{0.04} \rm{M_{\odot}}$ ."
+In the following sections we adopt the higher masses M...measured in ?.., In the following sections we adopt the higher masses measured in \citet{Delfosse2000}. .
+ We. however. note that two- RV measurements for our systems giving masses," We, however, note that two-component RV measurements for our systems giving masses"
+11. 13 ancl March 6 using the Coolecl Infrared Speetrograph and Camera for the Oll-suppression spectrograph. CISCO (Alotohara οἱ al.,"11, 13 and March 6 using the Cooled Infrared Spectrograph and Camera for the OH-suppression spectrograph, CISCO (Motohara et al."
+ 1998)., 1998).
+ Vhe 2.12-5m image shown in Vie., The $\mu$ m image shown in Fig.
+ 6 is a section of the original 5« mosaic (provided by Masa Lavashi)., \ref{subaru} is a $\times$ section of the original $5\arcmin\times5\arcmin$ mosaic (provided by Masa Hayashi).
+ In order to overlay the ΟΙ maser positions onto the Subaru image. it was first necessary to astrometricallv register the image with the 2000 coordinate svstem.," In order to overlay the OH maser positions onto the Subaru image, it was first necessary to astrometrically register the image with the J2000 coordinate system."
+ This was done using stars from the catalogue. of 4Hillenbrand Carpenter (2000) ancl the StarlinkGALA software package., This was done using stars from the catalogue of Hillenbrand Carpenter (2000) and the Starlink software package.
+ A computer programme was then written to read. the 110500 OLL maser positions and convert them to J2000 coordinates using the Starlink Astrometry Library ane produce a catalogue that was readable by theGALA software package., A computer programme was then written to read the B1950 OH maser positions and convert them to J2000 coordinates using the Starlink Astrometry Library and produce a catalogue that was readable by the software package.
+ The positions were then overlaid usingGAIA., The positions were then overlaid using.
+ The 1665- and 1667-MlIZz masers (shown by dot. and cross symbols. respectively) are concentrated: towards. the hiehly obscured Hic2 region., The 1665- and 1667-MHz masers (shown by dot and cross symbols respectively) are concentrated towards the highly obscured IRc2 region.
+ In general they show Little correspondence with features in the near-Ht map., In general they show little correspondence with features in the near-IR map.
+ However the 1612-MlIIz masers (1) correspond with fingers of Lo emission in the molecular outllow., However the 1612-MHz masers (+) correspond with fingers of $_{2}$ emission in the molecular outflow.
+ This association is scen more clearlv in Fig. 7.," This association is seen more clearly in Fig. \ref{subcore},"
+ which shows the central region on an expanded. scale., which shows the central region on an expanded scale.
+" Phe 1612-MlIe masers show a strong association with fingers of ο emission in the molecular outflow. particularly in the South. near RA (J2000) 35"" L277. Dec. (12000).446"". where 6 (half) of the 1612-Alllz masers are found."," The 1612-MHz masers show a strong association with fingers of $_{2}$ emission in the molecular outflow, particularly in the South, near RA (J2000) $^{\rm h}$ $^{\rm m}$ 7, Dec. (J2000), where 6 (half) of the 1612-MHz masers are found."
+ Three. 1665R masers (numbers 63. 64 and 76 in Table 1)) ave found nearby.," Three 1665R masers (numbers 63, 64 and 76 in Table \ref{OHtab}) ) are found nearby,"
+"universe with p,=0.021. /pc. Fig.9 compares directly. the predictions of the model (solid. lines) to the mass profiles for 29 clusters derived. by weak lensing in a published: sample (points)(Wu et al.","universe with $\rho_c=0.02 
+M_{\odot}$ $^3$, Fig.9 compares directly the predictions of the model (solid lines) to the mass profiles for 29 clusters derived by weak lensing in a published sample (points)(Wu et al."
+ L998)., 1998).
+ Models. are computed for the same mass range listed above: Alii3botoe.107:101. Ar., Models are computed for the same mass range listed above: $M_{vir}=3 \ 10^{14};10^{15};3 \ 10^{15};10^{16} \ M_{\odot}$ .
+ In this case the agreemen with the observations appears a little better than for the Ixing model analysed in the previous section., In this case the agreement with the observations appears a little better than for the King model analysed in the previous section.
+" Phe correlation οποσα £245, and Ade, derived. within the hierarchica jxiceture is also plotted in Fig.9 as a long-dashed line.", The correlation between $R_{vir}$ and $M_{vir}$ derived within the hierarchical picture is also plotted in Fig.9 as a long-dashed line.
+ X goo agreement. between the theoretical curves ancl estimates of virial masses of observed. clusters (open circles) as inferrec ov Carlberg (1996) is clearly. seen., A good agreement between the theoretical curves and estimates of virial masses of observed clusters (open circles) as inferred by Carlberg (1996) is clearly seen.
+ ow. by adopting this model we are able to investigate he thermal radius {δν," Now, by adopting this model we are able to investigate the thermal radius $R_{th}$."
+ This length is not necessarily equa o the core radius., This length is not necessarily equal to the core radius.
+" In CL0024|1654 we predict a core radius roughly of A,=100 kpe but the thermal radius within which the thermalization is working is shorter. nearly 40 kpc."," In CL0024+1654 we predict a core radius roughly of $R_c=100$ kpc but the thermal radius within which the thermalization is working is shorter, nearly $40$ kpc."
+ This implies that the core is spherical for a few dozen of kiloparsec: bevond this radius the thermalization does not work and anisotropies will dominate., This implies that the core is spherical for a few dozen of kiloparsec; beyond this radius the thermalization does not work and anisotropies will dominate.
+ This is consistent with Miralca-IEscudé (2000) who finds evidence of elliptical rather than spherical potentials on scales of the order of 70 kpe. as for instance in MS21137-23.," This is consistent with Miralda-Escudé (2000) who finds evidence of elliptical rather than spherical potentials on scales of the order of $70$ kpc, as for instance in MS21137-23."
+ The great acvantage ollered by assuming afocal thermal equilibrium is that it allows to estimate the self-interaction cross-section cirectly from. the observational cata., The great advantage offered by assuming a thermal equilibrium is that it allows to estimate the self-interaction cross-section directly from the observational data.
+" Lo is the dark particle number density. 0 the cross section and e the dispersion velocity. assuming in the core a collision time T-l(nσc) close to the Hubble time. we derive: with nmi, the mass of the dark matter particle ancl p, 1e central density."," If $n$ is the dark particle number density, $\sigma$ the cross section and $v$ the dispersion velocity, assuming in the core a collision time $\tau = 1/(n \ \sigma \ v)$ close to the Hubble time, we derive: with $m_x$ the mass of the dark matter particle and $\rho_c$ the central density."
+ Even if we ignore the nature of such dark matter particles it is important to point out that our model of selt-interaction is characterized by a cross section that depends on the halo dispersion velocity Le. as in other classical physical interactions the cross section is a function ofthe particle energy.," Even if we ignore the nature of such dark matter particles it is important to point out that our model of self-interaction is characterized by a cross section that depends on the halo dispersion velocity i.e., as in other classical physical interactions the cross section is a function of the particle energy."
+" Lt is important to point out that if one assumes ofr, constant (or equivalentIv pec|cz constant). then the halo core racius will scale with v as A2,xe27. as proposed by Miralda-Escudé (2000)."," It is important to point out that if one assumes $\sigma/m_x$ constant (or equivalently $\rho_c \ v \ t \approx$ constant), then the halo core radius will scale with v as $R_c \propto v^{3/2}$, as proposed by Miralda-Escudé (2000)."
+ This predicted: trend is not [avored. by the present observational picture. even dough some uncertainty is present.," This predicted trend is not favored by the present observational picture, even though some uncertainty is present."
+" Assuming m/m, asa iminishing function of the dispersion. velocity allows the seli-interaction theory to predict: Z2,xe. in agreement with observations."," Assuming $\sigma/ m_x$ as a diminishing function of the dispersion velocity allows the self-interaction theory to predict: $R_c \propto v$, in agreement with observations."
+ During the refereeing of this paper. Yoshida et al. (," During the refereeing of this paper, Yoshida et al. ("
+2000) and Dave et al. (,2000) and Davé et al. (
+2000) have carried out cosmological body simulations for ai sell-interacting CDM moclel with a constant cross-section.,2000) have carried out cosmological N-body simulations for a self-interacting CDM model with a constant cross-section.
+ ln a more recent study (Firmani. D'Onshia & Chincarini 2000). using an original cosmological code based. on the collisional Boltzmann equation. we have confirmed that self-interaction creates in he halo a non-sngular isothermal core. the outer region remaining as a NEW profile.," In a more recent study (Firmani, D'Onghia $\&$ Chincarini 2000), using an original cosmological code based on the collisional Boltzmann equation, we have confirmed that self-interaction creates in the halo a non-singular isothermal core, the outer region remaining as a NFW profile."
+ Assuming @xl/r we jwe proved that a nearly central density scale invariance is reproduced., Assuming $\sigma \propto 1/v$ we have proved that a nearly central density scale invariance is reproduced.
+ Further. à study. of Wyithe et al... (," Further, a study of Wyithe et al. ("
+2000),2000)
+Si: 1108.36. 1100.91. 1113.17. 1113.20. 1113.23 blend: 1125.6 + 1127.,": 1108.36, 1109.94, 1113.17, 1113.20, 1113.23 blend: 1125.6 + 1127."
+ +SUV 1128.3 +4+ blend: 1138-1111 deud: 11L0.5-1115.7, + 1128.3 + blend: 1138-1141 blend: 1140.5-1145.7.
+ While it is obvious that there must be some broad emission around 1175 there is also absorption present., While it is obvious that there must be some broad emission around 1175 there is also absorption present.
+ Since we do not know a priori how iuuch emission aud absorption there is. we cannot model this region.," Since we do not know a priori how much emission and absorption there is, we cannot model this region."
+ Decreasing C aud Si further than 0.1 aud 0.3) (respectively) led to improvement of the fit in some parts of the spectrum but the fit also deteriorated iu other regions., Decreasing C and Si further than 0.1 and 0.3 (respectively) led to improvement of the fit in some parts of the spectrum but the fit also deteriorated in other regions.
+ It is also important to note that the depth of the broad absorption features which actually form part of the continuum (e.g. as [or Si) depends uot ouly on the abtuucances but also on the rotational velocity., It is also important to note that the depth of the broad absorption features which actually form part of the continuum (e.g. as for Si) depends not only on the abundances but also on the rotational velocity.
+ Lucreasing the rotational velocity has the same effect as reduciug the abuudauces. as the depth of the absorptiou features decreases.," Increasing the rotational velocity has the same effect as reducing the abundances, as the depth of the absorption features decreases."
+ Overall. the (around 1065 )). theSII (arouud 1110 )) and the Si-C blend (in the range 1138-1116 )) were the main absorption features that drove both the Si aud the C to sub-solar values.," Overall, the (around 1065 ), the (around 1110 ) and the Si-C blend (in the range 1138-1146 ) were the main absorption features that drove both the Si and the C to sub-solar values."
+ As these abundances were reduced. the fit of the theoretical spectrum also slightly improved tn other regions (such as around 1055 aand 1080 )).," As these abundances were reduced, the fit of the theoretical spectrum also slightly improved in other regions (such as around 1055 and 1080 )."
+ We feel coufideut that the carbon aud silicon WD photospheric abundances are actually well represented by the non-solar best fit models. as these were derived [rom fitting these regions of the spectrum that were not contaminated by broad emission features and/or sharp interstellar absorption features.," We feel confident that the carbon and silicon WD photospheric abundances are actually well represented by the non-solar best fit models, as these were derived from fitting these regions of the spectrum that were not contaminated by broad emission features and/or sharp interstellar absorption features."
+ For Nitrogen. tlie results are iuconclusive as the Nitrogen features were coutamiuated by sharp absorptious possibly originating [rom the ISM or from the immediate surroundings of the systel.," For Nitrogen, the results are inconclusive as the Nitrogen features were contaminated by sharp absorptions possibly originating from the ISM or from the immediate surroundings of the system."
+ Iu Figure | we show the residual emission. namely. we subtracted the theoretical spectrum from the observed one.," In Figure 4 we show the residual emission, namely, we subtracted the theoretical spectrum from the observed one."
+ The apparent emission features in the residual spectrum are the doublet (with an interstellar absorption). the two lines (977 aand 1175 )) ancl possibly alsoNIT.NIIL.NIV audSUT as marked in Figure L," The apparent emission features in the residual spectrum are the doublet (with an interstellar absorption), the two lines (977 and 1175 ) and possibly also, and as marked in Figure 4."
+ Such emission features are olten seen in theFUSE spectra of IP systems or other DNe systems aud it implies that gas is possibly, Such emission features are often seen in the spectra of IP systems or other DNe systems and it implies that gas is possibly
+profile to linearize the lensing transformation with respect to the shapelets basis and takes advantage of the well-described mathematics of the quantum simple harmonic oscillator.,profile to linearize the lensing transformation with respect to the shapelets basis and takes advantage of the well-described mathematics of the quantum simple harmonic oscillator.
+" This linear approximation of the surface brightness profile is how the shapelets method fundamentally differs from moment methods, which use ratios of measured moments to approximate the flexion fields."," This linear approximation of the surface brightness profile is how the shapelets method fundamentally differs from moment methods, which use ratios of measured moments to approximate the flexion fields."
+" In both approaches, assumptions on the unlensed surface brightness profile are made in order to measure the lensing fields."," In both approaches, assumptions on the unlensed surface brightness profile are made in order to measure the lensing fields."
+" These assumptions are that odd-order shapelets/moments in observed images are sourced by lensing flexion alone, meaning that the intrinsic two-dimensional skewness of unlensed galaxies is assumed to be, on average, zero."," These assumptions are that odd-order shapelets/moments in observed images are sourced by lensing flexion alone, meaning that the intrinsic two-dimensional skewness of unlensed galaxies is assumed to be, on average, zero."
+" The magnitude of the flexion scatter from intrinsic galaxy profile shapes is an important quantity to understand, as it sets the finite resolution of flexion mass measurements."," The magnitude of the flexion scatter from intrinsic galaxy profile shapes is an important quantity to understand, as it sets the finite resolution of flexion mass measurements."
+ This limit is analogous to how the distribution of intrinsic ellipticities limits the spatial resolution of weak lensing shear mass measurements., This limit is analogous to how the distribution of intrinsic ellipticities limits the spatial resolution of weak lensing shear mass measurements.
+" In each of the current flexion measurement methods, there are systematic issues that must be addressed."," In each of the current flexion measurement methods, there are systematic issues that must be addressed."
+ Shapelet-based lensing measurements have been shown to be unreliable and biased for large shears ∙, Shapelet-based lensing measurements have been shown to be unreliable and biased for large shears .
+"∙ This is a significant concern for measuring flexion, since the shear can not be assumed to be small in the regime where flexion fields are measurable."," This is a significant concern for measuring flexion, since the shear can not be assumed to be small in the regime where flexion fields are measurable."
+" Moment-based measurements of flexion rely on high-order surface brightness moments, making the strength of the lensing signal extracted very sensitive to the window functions used in the moment calculations: though the flexion symmetries under coordinate transformations are matched by third-order moments of the surface brightness, the normalization of the moments required to isolate the flexion signal includes both fourth- and sixth-order moments ∙"," Moment-based measurements of flexion rely on high-order surface brightness moments, making the strength of the lensing signal extracted very sensitive to the window functions used in the moment calculations: though the flexion symmetries under coordinate transformations are matched by third-order moments of the surface brightness, the normalization of the moments required to isolate the flexion signal includes both fourth- and sixth-order moments ."
+∙ Characterizing the noise properties of these moments and how the noise propagates into uncertainties in the measured flexion fields is a complex problem., Characterizing the noise properties of these moments and how the noise propagates into uncertainties in the measured flexion fields is a complex problem.
+" Additionally, the implementations of both these methods and their application to observational data have not addressed the well-established mass-sheet degeneracy in a consistent"," Additionally, the implementations of both these methods and their application to observational data have not addressed the well-established mass-sheet degeneracy in a consistent"
+annulus is given by where NyGc.0) denotes the number of observable modes in the annulus (only modes whose line-of-sight components fit within the observed. bandpass are included).,"annulus is given by where $N_{\rm m}(k,\theta)$ denotes the number of observable modes in the annulus (only modes whose line-of-sight components fit within the observed bandpass are included)."
+ In terms of the A-vector components & and 8. the number of independent Fourier modes within an annulus of racial width dA and angular width d6 is Αμ.8)=2xkVsin(6)dk d0/(2x)*. where V=DANDI£A). is. the observed volume.," In terms of the $k$ -vector components $k$ and $\theta$, the number of independent Fourier modes within an annulus of radial width $dk$ and angular width $d\theta$ is $N_{\rm m}(k,\theta) = 2\pi k^2 V \sin(\theta) dk\,d\theta/(2\pi)^3$ , where $V = D^2\Delta D(\lambda^2/A_{\rm e})$ is the observed volume."
+ Averaging ap(hk.@) over @ gives the spherically averaged sensitivity to the 21-em power spectrum aps). In order to find the sensitivity in terms of AS) we use Equation (70)). which gives In this section. we present the results of our simulations and analysis.," Averaging $\sigma_P(k,\theta)$ over $\theta$ gives the spherically averaged sensitivity to the 21-cm power spectrum $\sigma_P(k)$, In order to find the sensitivity in terms of $\Delta^2_{21}$ we use Equation \ref{Deltasq}) ), which gives In this section, we present the results of our simulations and analysis."
+ Section δι shows the result of applying the polarised. foreground. removal algorithm for a single line of sight., Section \ref{Single line of sight} shows the result of applying the polarised foreground removal algorithm for a single line of sight.
+ Section S.2. gives the results for a full three-dimension data set., Section \ref{Full data cube} gives the results for a full three-dimension data set.
+ Figure 7 shows the resulting spherically averaged 2]-cm power spectra., Figure \ref{Spherically averaged power spectra} shows the resulting spherically averaged 21-cm power spectra.
+ Figure 7 demonstrates the removal of a single polarised foreground component from a noiscless Stokes £ signal., Figure \ref{clean_1d_eg} demonstrates the removal of a single polarised foreground component from a noiseless Stokes $I$ signal.
+ For this example. we assume that any continuum DOSE has already. been removed. leaving only the cosmic signal and leaked polarised foreground.," For this example, we assume that any continuum DGSE has already been removed, leaving only the cosmic signal and leaked polarised foreground."
+ Pwo dilferent values of Faracay depth have been moclellecl in order to show the dillerence in cleaning performance., Two different values of Faraday depth have been modelled in order to show the difference in cleaning performance.
+ The left-hand. cxample uses oO=35.5 mm.7 with pO)z176.8 Ix and. xoz0.794., The left-hand example uses $\phi = 35.5$ $^{-2}$ with $p(\lambda^2_0) \approx 176.8$ K and $\chi_0 \approx 0.794$.
+ The right-hand. example uses ὁ=5.5 2 with 161.6 Ix and xoz—1.52., The right-hand example uses $\phi = 5.5$ $^{-2}$ with $p(\lambda^2_0) \approx 161.6$ K and $\chi_0 \approx -1.52$.
+ The top Stokes £ visibility versus frequency. plots in Figure 7 show the signal to be cleaned (blachk)) ancl the model cosmic signal (5fue)) for reference., The top Stokes $I$ visibility versus frequency plots in Figure \ref{clean_1d_eg} show the signal to be cleaned ) and the model cosmic signal ) for reference.
+" Note that. the frequeney axis is linear ancl therefore the apparent ""period? of the contaminated. signal containing the leakecl polarised foreground is non-uniform.", Note that the frequency axis is linear and therefore the apparent `period' of the contaminated signal containing the leaked polarised foreground is non-uniform.
+ This is because of the A7-dependence of the plane of polarisation as per Equation (8))., This is because of the $\lambda^2$ -dependence of the plane of polarisation as per Equation \ref{eq:chi}) ).
+ The second row of plots show the ellective Stokes £ Faraday dispersion. function. P;(ó). of the contaminated signal. including the real and imaginary part as well as its modulus.," The second row of plots show the effective Stokes $I$ Faraday dispersion function, $F_I(\phi)$, of the contaminated signal, including the real and imaginary part as well as its modulus."
+ The fit to each of these components (dashed)) are shown for comparison., The fit to each of these components ) are shown for comparison.
+ The third row of plots show the cleaned effective Stokes £ Faraday dispersion function with the contaminated ellective Stokes £ Faraday dispersion function (fain/)) shown [or comparison., The third row of plots show the cleaned effective Stokes $I$ Faraday dispersion function with the contaminated effective Stokes $I$ Faraday dispersion function ) shown for comparison.
+ The bottom row of plots show the cleaned signal (black)) after being Fourier inverted. and the moclel cosmic signal (bfie)) for reference., The bottom row of plots show the cleaned signal ) after being Fourier inverted and the model cosmic signal ) for reference.
+ ]t is evident that the instrumental polaristion leakage due to the Faraday. sereen at. the smaller. depth: Faraday screen) is less effectively removed than the deeper screen., It is evident that the instrumental polaristion leakage due to the Faraday screen at the smaller depth Faraday screen) is less effectively removed than the deeper screen.
+ This is due in part to a residual long-wavelength Iluctuation in the signal after continuum foreground removal which leaves small-ó power in the effective Stokes £ Faraday dispersion function., This is due in part to a residual long-`wavelength' fluctuation in the signal after continuum foreground removal which leaves $\phi$ power in the effective Stokes $I$ Faraday dispersion function.
+ ltesiduals such as these can make a small-ó peak in the Stokes £ Faraday. dispersion function asvinmetric and therefore more οΠο to fit. accurately., Residuals such as these can make a $\phi$ peak in the Stokes $I$ Faraday dispersion function asymmetric and therefore more difficult to fit accurately.
+ Deeper peaks are not alfected in this way as the miocelled continuum foreground. does not contain Ductuations that masquerade at these Faraday depths. therefore no residuals ab these depths are present after continuum. subtraction making more accurate fitting possible.," Deeper peaks are not affected in this way as the modelled continuum foreground does not contain fluctuations that masquerade at these Faraday depths, therefore no residuals at these depths are present after continuum subtraction making more accurate fitting possible."
+ Unwanted small-o resicuals have been removed. by applving a high-pass filter to the cleaned. cllective Stokes { Faraday dispersion. function., Unwanted $\phi$ residuals have been removed by applying a high-pass filter to the cleaned effective Stokes $I$ Faraday dispersion function.
+ The extent of this filter is 0xὡς<21.2 mm27. which has an equivalent width in frequeney-space of approximately 4.7 MIIz for the correlator specifications used here.," The extent of this filter is $0 \leq \phi_{\rm f} \lsim\,\,21.2$ $^{-2}$, which has an equivalent width in frequency-space of approximately 4.7 MHz for the correlator specifications used here."
+" Negligible cosmic signal is removed by this procedure since the Lok. Huctuations are of shorter ""wavelength! at the corresponding redshift at this frequency.", Negligible cosmic signal is removed by this procedure since the EoR fluctuations are of shorter `wavelength' at the corresponding redshift at this frequency.
+ Figure SN. shows a simulated. three-dimensional Stokes / data cube over a 32 MlIIz band. centred on 157.8. ΑΔΗ >< 7.8) using 128 channels.," Figure \ref{plots_z7-0} shows a simulated three-dimensional Stokes $I$ data cube over a 32 MHz band centred on 157.8 MHz $6.3 \lsim\,\,z \lsim\,\,7.8$ ) using 128 channels."
+ This simulation includes both continuum and polarised foreground removal., This simulation includes both continuum and polarised foreground removal.
+ Faraday depth has been modelled as a gradient. over skv-plane with 1lxox 5radmm and the intrinsic polarisation angle for each line of sight has been randomly assigned. —2/2<vXxm/2 rad.," Faraday depth has been modelled as a gradient over sky-plane with $1 \leq \phi \leq 5$ $^{-2}$ and the intrinsic polarisation angle for each line of sight has been randomly assigned, $-\pi/2 < \chi \leq \pi/2$ rad."
+ The upper-left cube shows the full Stokes { signal before cleaning., The upper-left cube shows the full Stokes $I$ signal before cleaning.
+ The upper-right cube has had. the continuum DOSE fitted and removed in each line of sight. leaving the cosmic signal and the elfects of instrumental polarisation leakage.," The upper-right cube has had the continuum DGSE fitted and removed in each line of sight, leaving the cosmic signal and the effects of instrumental polarisation leakage."
+" The banding along the 8,- ancl 8,- frequency planes result. from the oscillatory behaviour of the leaked polarisecl signal.", The banding along the $\theta_x$ - and $\theta_y$ -frequency planes result from the oscillatory behaviour of the leaked polarised signal.
+ The micelle-lef cube is the raw Lully cleaned: signal having had the leaked polarised signal removed., The middle-left cube is the raw fully cleaned signal having had the leaked polarised signal removed.
+ The pre-Ieakage-cleaned bancing has mostly disappeared. with the exception of channels close to the edge of the frequeney band.," The pre-leakage-cleaned banding has mostly disappeared, with the exception of channels close to the edge of the frequency band."
+ This is a svstematic effect. due in xwt to the continuum foreground removal., This is a systematic effect due in part to the continuum foreground removal.
+ The central-right cube shows the fully cleaned signal after high-pass filtering of the elfective Stokes £ Faraday dispersion function of each ine of sight., The central-right cube shows the fully cleaned signal after high-pass filtering of the effective Stokes $I$ Faraday dispersion function of each line of sight.
+ The same filter width has been usec as for he one-dimensional example in Section S.1.., The same filter width has been used as for the one-dimensional example in Section \ref{Single line of sight}.
+ The centra requeney slice of the fully cleaned and. high-pass filterec cube is shown at the lower-right. with the model Eolt signa shown at the lower-left for comparison.," The central frequency slice of the fully cleaned and high-pass filtered cube is shown at the lower-right, with the model EoR signal shown at the lower-left for comparison."
+ The model oR. cube was not had each line of sights mean subtracted. which would occur as a consequence ofthe continuum foregroune subtraction.," The model EoR cube has not had each line of sight's mean subtracted, which would occur as a consequence ofthe continuum foreground subtraction."
+ Figure 9 shows a selection. of line-of-sight signa profiles from the three-dimoensional simulation for the moce cosmic signal (6lack)). fully cleaned signal with no hish- filtering of their effective Stokes £ Faraday. dispersion function (red)) and fully cleaned. signal after. high-pass filtering of their effective. Stokes £ Faraday. dispersion function (bfuec)).," Figure \ref{plots_z7-2} shows a selection of line-of-sight signal profiles from the three-dimensional simulation for the model cosmic signal ), fully cleaned signal with no high-pass filtering of their effective Stokes $I$ Faraday dispersion function ) and fully cleaned signal after high-pass filtering of their effective Stokes $I$ Faraday dispersion function )."
+OIL iuaser cinission from the four groui-state transitions at 1612.231. 1665.102. 1667.359 aud 1720.530 MIIz occurs in a wide range ofobjects.,"OH maser emission from the four ground-state transitions at 1612.231, 1665.402, 1667.359 and 1720.530 MHz occurs in a wide range of objects."
+ Extragalactic OTlaser enission is associated with active ealactic nuclei aud chhanced ceutral star formation (Baan Taschick 1987: BRaudell 1995). with strongest cussion. iu general. from the main lines at 1665/1667 MITz.," Extra–galactic OH--maser emission is associated with active galactic nuclei and enhanced central star formation (Baan Haschick 1987; Randell 1995), with strongest emission, in general, from the main lines at 1665/1667 MHz."
+ ΟΠ maser ciuiisson from galactic star forming reeious (SFRs) is also strongest in the main lines. with dominant cussion mostly at 1665 MIIz.," OH maser emisson from galactic star forming regions (SFRs) is also strongest in the main lines, with dominant emission mostly at 1665 MHz."
+ The satellite line at 1720 ΑΠ is best known for if» occurence din supernova remnants. whereas that at 1612 MIIz is strongest iu the opticallythick cuvelopes of oxvgenvich asvinptoticgiantbrauch (ACB) stars.," The satellite line at 1720 MHz is best known for its occurence in supernova remnants, whereas that at 1612 MHz is strongest in the optically–thick envelopes of oxygen–rich asymptotic–giant–branch (AGB) stars."
+ Mainline ΟΠ emission may be preseut in opticallythin cuvelopes: frou nearby. Mira. variables. with low mass-loss rates around * ML. vr|. only mainline cmission is detected.," Mainline OH emission may be present in optically–thin envelopes; from nearby Mira variables, with low mass-loss rates around $^{-7}$ $_\odot$ $^{-1}$, only mainline emission is detected."
+ At the end of the ACB evolutionary pliase. the star uudergoes a transformation which takes a couple of thousand vears. duriug which it evolves to become a planctary nebula (PN).," At the end of the AGB evolutionary phase, the star undergoes a transformation which takes a couple of thousand years, during which it evolves to become a planetary nebula (PN)."
+ During this post-AGD stage. the tass-loss rate decreases by a factor of 1000 or more whilst the stellar surface temperature rises by a factor of LO or 20 to above 50.000 Ik. The star changes frou losing mass in a cool. deuse wind to losing mass in a niuch hotter. lower deusitv wiud.," During this post-AGB stage, the mass-loss rate decreases by a factor of 1000 or more whilst the stellar surface temperature rises by a factor of 10 or 20 to above 50,000 K. The star changes from losing mass in a cool, dense wind to losing mass in a much hotter, lower density wind."
+ Strong changes also occur in the envelope kinematics and morphologies., Strong changes also occur in the envelope kinematics and morphologies.
+ Iu particular. whilst the οποιοματ envelopes of ACB stars are generally spherically svinmnetric. tages of planetary nebulae show axi-sviunetric morphologies in over of cases (ATauchacdo 22000).," In particular, whilst the circumstellar envelopes of AGB stars are generally spherically symmetric, images of planetary nebulae show axi-symmetric morphologies in over of cases (Manchado 2000)."
+ Diving the initial searches in the 19708 it was quickly found that ecluission was not detected from AGB stars;, During the initial searches in the 1970s it was quickly found that emission was not detected from AGB stars.
+ As this agreed with theoretical expectations (ce., As this agreed with theoretical expectations (eg.
+" Elitzur. Coldreich Scoville 1976). there have been no systematic searches for this maser in stellar sources,"," Elitzur, Goldreich Scoville 1976), there have been no systematic searches for this maser in stellar sources."
+ Towever. it is imaginable that the shocks in some postACB stars could teiiporarilv induce nunaser enissjon.," However, it is imaginable that the shocks in some post–AGB stars could temporarily induce maser emission."
+ We therefore decided to observe. for the first time. a modest sample of selected sources. inchiding postAGB objects. at 1720 MIIz.," We therefore decided to observe, for the first time, a modest sample of –selected sources, including post–AGB objects, at 1720 MHz."
+ A total of 18 objects were selected from Seveuster ((1997a.b) on the basis of their unusual OT 1612MUTz profiles.," A total of 18 objects were selected from Sevenster (1997a,b) on the basis of their unusual OH 1612–MHz profiles."
+ Observations in all four eroundstate OTL transitions were taken with the Australia Telescope Compact Array (ATCA) in Noveiiber 1998 (1612. 1665 and 1667 MIIZ) and September 1999 (1720 MITz).," Observations in all four ground–state OH transitions were taken with the Australia Telescope Compact Array (ATCA) in November 1998 (1612, 1665 and 1667 MHz) and September 1999 (1720 MHz)."
+ Eleven of the sources are likely postACD objects., Eleven of the sources are likely post–AGB objects.
+" The full results from this study will be published separately,", The full results from this study will be published separately.
+ Hore we report the detection of 1720 MIIz frou the source (IRAS 2116)., Here we report the detection of 1720 MHz from the source (IRAS $-$ 2116).
+ In 2. we present the observations aud in 3 the evidence that jis a postACD object.," In 2, we present the observations and in 3 the evidence that is a post–AGB object."
+ In Eve discuss a possible scenario for the cexcitation., In 4 we discuss a possible scenario for the excitation.
+ Conclusions are geiveu in 5., Conclusions are given in 5.
+ Figue sshows the ΟΠ spectra for (CI8:07:20.9. 21:16:10.9 (J2000)). observed with a chaunel width of 0.7ο," Figure shows the OH spectra for (18:07:20.9, $-$ 21:16:10.9 (J2000)), observed with a channel width of 0.7."
+", The maser cuiission is stronecst at 1612 MIIz. with a peak flux deusitv of 11.0 Jv at 71.2L."," The maser emission is strongest at 1612 MHz, with a peak flux density of 11.0 Jy at 71.2."
+ The spectral profile has a one-sided' peak and a πιο weaker feature at LOL1., The spectral profile has a `one-sided' peak and a much weaker feature at 101.
+. The eenission covers the same velocity range as the eecniüssion 2). with a peak flux density of 1.61 Jv at 71.3L.," The emission covers the same velocity range as the emission ), with a peak flux density of 1.64 Jy at 71.3."
+. No emission was detected at 1667 ΑΠΣ with a detection nut of ~ SO mJy., No emission was detected at 1667 MHz with a detection limit of $\sim$ 80 mJy.
+ The eenission Is seen as a suele narrow peak at a slieltly more bluc-slüfted velocity than the, The emission is seen as a single narrow peak at a slightly more blue-shifted velocity than the
+Observations of the Cosmic Microwave. Background. ancl other astronomical distance determinations have convincnglv demonstrated that (he expansion of the universe started accelerating at a recent epoch (see IXomatsu 2010 and references therein).,Observations of the Cosmic Microwave Background and other astronomical distance determinations have convincingly demonstrated that the expansion of the universe started accelerating at a recent epoch (see Komatsu 2010 and references therein).
+" All of the available observations are consistent with a cosmological constant. or the effect of a constant negalive pressure with value (we use Planck units with ο—h=G1 throughout this paper) lere. Z4, is the value of the Hubble constant al present ancl Qyy is the ratio of the negalive pressure to the critical density of the universe. and we have used the values obtained"," All of the available observations are consistent with a cosmological constant, or the effect of a constant negative pressure with value (we use Planck units with $c=\hbar=G=1$ throughout this paper) Here, $H_0$ is the value of the Hubble constant at present and $\Omega_{\Lambda 0}$ is the ratio of the negative pressure to the critical density of the universe, and we have used the values obtained"
+One can iniuedia‘ly obtain the cosmological paraietcryf ouce a function for P(f) is chosen.,One can immediately obtain the cosmological parameters once a function for $\Gamma(t)$ is chosen.
+" Alternatively, one can obtain the parameters by introducing one mere constraint ou one of the cosmological parameters: for instance. on the energy density of the fluid p or ou the EoS paraicter."," Alternatively, one can obtain the parameters by introducing one more constraint on one of the cosmological parameters; for instance, on the energy density of the fluid $\rho$ or on the EoS parameter."
+" Trivially, one may also assume that he fluid is a perfect fluid (i6. 5=0: thus P is null) and would then obtain the solutions eiveu by Iunar&Singh(2007) for au exponentially expanding Bianchi vpe-I universe."," Trivially, one may also assume that the fluid is a perfect fluid (i.e., $\gamma=0$; thus $\Gamma$ is null) and would then obtain the solutions given by \cite{kumar07} for an exponentially expanding Bianchi type-I universe."
+ One can observe that for anv perfect Huid. i.c.. in case of P=0. the initial anisotropy of the expansion dies away monotonically as £ increases. while it may exhibit various non-trivial behaviors according o the function chosen for Et(f).," One can observe that for any perfect fluid, i.e., in case of $\Gamma=0$, the initial anisotropy of the expansion dies away monotonically as $t$ increases, while it may exhibit various non-trivial behaviors according to the function chosen for $\Gamma(t)$."
+ Note that we need one more constraint in addition to the one eiven in (35) to fully determine the system., Note that we need one more constraint in addition to the one given in (35) to fully determine the system.
+ Since the first assumption eiven by (35) causes de Sitter volumetric expansion. it secs reasouable to choose our last assumption as If we solve the «μοι of equations which consists of the three Einstein feld equations (28)-(30) and the two constraints given by (35) aud (13). we obtain the following exact expressions for the scale factors: The directional IIbubble paramcters are obtained as follows: The anisotropy parameter of the expansion aud thus the anisotropy energv density are found to be cynical. The eunerev density of the fluid is also found to be dynamical+ in. a way so as to secure p|pa=.P3&7. aud the skewness parameter of the EoS parameter is obtained as follows: One can check that this solutiou vields (20).," Since the first assumption given by (35) causes de Sitter volumetric expansion, it seems reasonable to choose our last assumption as If we solve the system of equations which consists of the three Einstein field equations (28)-(30) and the two constraints given by (35) and (43), we obtain the following exact expressions for the scale factors: The directional Hubble parameters are obtained as follows: The anisotropy parameter of the expansion and thus the anisotropy energy density are found to be dynamical, The energy density of the fluid is also found to be dynamical in a way so as to secure $\rho+\rho_{\beta}=3k^{2}$, and the skewness parameter of the EoS parameter is obtained as follows: One can check that this solution yields (20)."
+ The positiveness condition on the enerev deusity of the fluid (p>O) imposes the restriction that the anisotropy of the expansion is smaller than 2 (A<2) aud that iurposes A to be positive., The positiveness condition on the energy density of the fluid $\rho>0$ ) imposes the restriction that the anisotropy of the expansion is smaller than 2 $\Delta<2$ ) and that imposes $\lambda$ to be positive.
+ Thus. only those models with positive A values are viable. aud hence in the following we will consider only the positive A values.," Thus, only those models with positive $\lambda$ values are viable, and hence in the following we will consider only the positive $\lambda$ values."
+ A decreases inonotonicallv as f£ iucreases and A> ft>|o., $\Delta$ decreases monotonically as $t$ increases and $\Delta\rightarrow 0$ as $t\rightarrow+\infty$.
+ Thus. the space approaches isotropy as foo|x in this model.," Thus, the space approaches isotropy as $t\rightarrow+\infty$ in this model."
+ Ou the other hand. ast >0.," On the other hand, $\Delta\rightarrow 2\left[1-\kappa/(\kappa+3\lambda k)\right]^{2}$ as $t\rightarrow 0$."
+ po»ο.P aud >028 f»|x., $\rho\rightarrow 3k^{2}$ and $\gamma\rightarrow 0$ as $t\rightarrow+\infty$.
+" That is. the EoS paranueter of the finid isotropizes as f increases and muundes the cosmnological constaut as t+|xX,"," That is, the EoS parameter of the fluid isotropizes as $t$ increases and mimics the cosmological constant as $t\rightarrow+\infty$."
+ On the other hand. p>3h?(1(lL|4)*) aud STIGNA which is always a negative nuniber. as foO ," On the other hand, $\rho\rightarrow 3k^{2}\left({1-\left(1+\frac{\kappa}{\lambda k}\right)^{-2}}\right)$ and $\gamma\rightarrow -\frac{27{\lambda}^{2}k^{2}}{{\kappa}^{2}+6\lambda\kappa k}$, which is always a negative number, as $t\rightarrow 0$."
+Therefore. as expected from the isotropizatiou condition discussed in Sect.," Therefore, as expected from the isotropization condition discussed in Sect."
+ 2.5.," 2.5.,"
+ p behaves like a phantom cnerev and increases as V dueroases, $\rho$ behaves like a phantom energy and increases as $V$ increases.
+ This result is consistent with the following situation: 5 is always negative aud thus EoS parameters on the y and 2 axes are passing the phantom divide liue. ic. πο," This result is consistent with the following situation: $\gamma$ is always negative and thus EoS parameters on the $y$ and $z$ axes are passing the phantom divide line, i.e., $w_{y,z}<-1$."
+ While e ds assumed to be. 1 so as to close the svstem in the preceding section. it is allowed to be a function of the cosmic time £ in this section. but. this time. the energv deusitv of the fiuid is assumed to be constant: If we solve the system of equations which consists of the three Einstein field equations (28)-(30) and the two constraints given by (35) aud (51). we obtain the following exact expressions for the scale factors: The directional IIubble parameters are obtained as follows:," While $w$ is assumed to be $-1$ so as to close the system in the preceding section, it is allowed to be a function of the cosmic time $t$ in this section, but, this time, the energy density of the fluid is assumed to be constant: If we solve the system of equations which consists of the three Einstein field equations (28)-(30) and the two constraints given by (35) and (51), we obtain the following exact expressions for the scale factors: The directional Hubble parameters are obtained as follows:"
+"A further partial transit was obtained on 2010 June 30 with the camera (FS01) on the LCOGT 2.0-m Faulkes Telescope South (FTS, Siding Spring. Australia),","A further partial transit was obtained on 2010 June 30 with the camera (FS01) on the LCOGT 2.0-m Faulkes Telescope South (FTS, Siding Spring, Australia)."
+" The Pan-STARRS-Z filter was used with the instrument in binning 2x2 mode, giving 0.303 aresec/pixel, and no delocussing."," The Pan-STARRS-Z filter was used with the instrument in binning $\times$ 2 mode, giving 0.303 arcsec/pixel, and no defocussing."
+" The exposure time was 45 s, and 169 images were taken in the 3.5 h period."," The exposure time was 45 s, and 169 images were taken in the 3.5 h period."
+ Data were reduced using a written at Liverpool John Moores University. then differential photometry performed using the package.," Data were reduced using a written at Liverpool John Moores University, then differential photometry performed using the package."
+ Eight comparison stars were used with an 8 pixel radius aperture., Eight comparison stars were used with an 8 pixel radius aperture.
+ Points with error bars larger than 0.01 mag were removed [rom the second half of the dataset., Points with error bars larger than 0.01 mag were removed from the second half of the dataset.
+" These were likely caused by passing cloud, and the observation was stopped prematurely due to bad weather."," These were likely caused by passing cloud, and the observation was stopped prematurely due to bad weather."
+" In order to establish the planetary nature and determine the stellar parameters, we obtained spectroscopic observations."," In order to establish the planetary nature and determine the stellar parameters, we obtained follow-up spectroscopic observations."
+ Five spectra were taken between 2010 April 17 and 2010 May 28 with the stabilised echelle spectrograph SOPHIE at the 1.93-m telescope of Observatoire de Haute-Provence (22)..," Five spectra were taken between 2010 April 17 and 2010 May 28 with the stabilised echelle spectrograph SOPHIE at the 1.93-m telescope of Observatoire de Haute-Provence \citep{Perruchot08,Bouchy09}."
+ The observations were all made with a signal-to-noise ratio of S/N~20 in order to minimise the Charge Transfer Inelfficieney (CTD elfect (2).., The observations were all made with a signal-to-noise ratio of $\sim$ 20 in order to minimise the Charge Transfer Inefficiency (CTI) effect \citep{Bouchy09}.
+" Two 3 arc-second diameter optical fibres were used, the first centred on the target and the second on the sky to simultaneously measure the background to remove contamination [rom scattered moonlight."," Two 3 arc-second diameter optical fibres were used, the first centred on the target and the second on the sky to simultaneously measure the background to remove contamination from scattered moonlight."
+" A further 8 spectra were obtained with the CORALIE Fibre-Fed Echelle Spectrograph on the Swiss 1.2-m telescope at ESO-La Silla, Chile between 2010 April 19 and 2010 June 29 with S/N~ 10-20 in dark/erey üme to minimise contamination Irom scattered moonlight."," A further 8 spectra were obtained with the CORALIE Fibre-Fed Echelle Spectrograph on the Swiss 1.2-m telescope at ESO-La Silla, Chile between 2010 April 19 and 2010 June 29 with $\sim$ 10–20 in dark/grey time to minimise contamination from scattered moonlight."
+ The data were processed using the standard pipeline (?2).," The data were processed using the standard pipeline \citep{Baranne96,Pepe02}."
+. The radial velocities (RVs) and line bisector spans (Vipau. derived [rom cross correlation are shown in Table 1. and plotted with the best-fit model in Figure 3..," The radial velocities (RVs) and line bisector spans \citep*[$V_{\rm span}$, derived from cross correlation are shown in Table \ref{rv-data} and plotted with the best-fit model in Figure \ref{RV}."
+" No significant correlation is seen between the bisector span and radial velocity, with a supporting the signals origin as a planetary companion rather than a blended eclipsing binary system (?).."," No significant correlation is seen between the bisector span and radial velocity, with a supporting the signal's origin as a planetary companion rather than a blended eclipsing binary system \citep{Queloz01}."
+" A Spearman rank-order correlation test also indicates that any correlation between the radial velocities and bisector i5weak, with a probability of correlation of 0.12."," A Spearman rank-order correlation test also indicates that any correlation between the radial velocities and bisector isweak, with a probability of correlation of 0.12."
+where we assume that the expelled matter keeps its aneular momentum with respect to the primary as df rapidly claubs out of the gravitational poteutial.,where we assume that the expelled matter keeps its angular momentum with respect to the primary as it rapidly climbs out of the gravitational potential.
+ The last term on the right side accounts for the distance between the primary and the dinars ceuter of mass., The last term on the right side accounts for the distance between the primary and the binary's center of mass.
+ This vields with For «50.1 and 50.2 one obtains f=1.3., This yields with For $q$ =0.1 and $\beta$ =0.2 one obtains $f$ =1.3.
+ A graphical evaluation of trajectories calculate by Wrun et al. (, A graphical evaluation of trajectories calculated by Wynn et al. (
+1997) for AE Aqr taking iuto account the smaller mass ratio leads to f=l.5.,1997) for AE Aqr taking into account the smaller mass ratio leads to $f$ =1.5.
+ We emphasize that these estimates are rough., We emphasize that these estimates are rough.
+ With such values of f the increase of orbital aueular moment loss compared ο Jinup Changes the sien of X and then leads to an accelerated growth of the lnass transfer rate [Eq. (, With such values of $f$ the increase of orbital angular momentum loss compared to $\dot{J}_{\rm {spin-up}}$ changes the sign of $X$ and then leads to an accelerated growth of the mass transfer rate [Eq. (
+10)].,10)].
+ Whether this occurs depends ou the detailed xocess during the swing aroun the white cwart., Whether this occurs depends on the detailed process during the swing around the white dwarf.
+" For f=1.3 and our standard case the time to reach arbitrarily large ALis only 10:9, ", For $f$ =1.3 and our standard case the time to reach arbitrarily large $\dot M$ is only $10^{4.6}$ ys.
+The rate cannot erow indefinitely. finally the maeuetic field becomes unable to handle the ever C»erowing mass transfer.," The rate cannot grow indefinitely, finally the magnetic field becomes unable to handle the ever growing mass transfer."
+ The efficiency of acceleration diminishes aud augular momoeutun loss from the syste ects limited., The efficiency of acceleration diminishes and angular momentum loss from the system gets limited.
+ The question arises how such systems develop., The question arises how such systems develop.
+ Tt the «πο stabilizes at a high transfer rate the secondary niav lose all its mass in a relatively short time aud a suele fast rotating magnetic white dwart remains., If the system stabilizes at a high transfer rate the secondary may lose all its mass in a relatively short time and a single fast rotating magnetic white dwarf remains.
+ The observation of RE J0317-555 (Burleigh ct al., The observation of RE J0317-853 (Burleigh et al.
+ 1999) with these unusual characteristics seenis to support possibility., 1999) with these unusual characteristics seems to support such a possibility.
+ Another path of evolution might leac to the formation of an accretion disk. which results in return of ALLUMLALY luonientuii back to the secondary. connected initially with a strong increase of accretion onto the white dwarf. as J becomes 1. aud then a rapid decrease of AL.," Another path of evolution might lead to the formation of an accretion disk, which results in return of angular momentum back to the secondary, connected initially with a strong increase of accretion onto the white dwarf, as $\beta$ becomes 1, and then a rapid decrease of $\dot M$."
+ But the strong inagnetie pressure of the white dwarf removes such a disk before the mass accretion rate has dropped to the low stable value Moy(NX [Eq. (, But the strong magnetic pressure of the white dwarf removes such a disk before the mass accretion rate has dropped to the low stable value $\dot M_{\rm {GW}}/X$ [Eq. (
+10)].,10)].
+ This could lead to a cyclic spiu-up/spin-«down ou short timescales. iuvolviug plases of high M. (compare the model for AE Aqr 1w Wrun et al.," This could lead to a cyclic spin-up/spin-down on short timescales, involving phases of high $\dot M$ (compare the model for AE Aqr by Wynn et al."
+ 1997) aud also rapid depletion of the secondary., 1997) and also rapid depletion of the secondary.
+ These phases are probably too short to be observable., These phases are probably too short to be observable.
+ If expelled matter would form a circunibinaryv disk. such a disk could extract siguificaut aneular momentum from the orbit aud thereby even lead to a supersoft X-ray - state (van Tecescling Nine 1995).," If expelled matter would form a circumbinary disk, such a disk could extract significant angular momentum from the orbit and thereby even lead to a supersoft X-ray - state (van Teeseling King 1998)."
+ The loss of magnetic[m] couplingC» aud the degeneracyOo of the secondarv stars are the important features of tli late evolution of AM Που systems., The loss of magnetic coupling and the degeneracy of the secondary stars are the important features of the late evolution of AM Her systems.
+ We conclude that AM Ier svstems should exist bevoud the period turuiug poiut (except the brown dwarf would be voung. hot :id. still magnetic - searches or the secondary stars in AM Πα systems near period miuiuma would be very desirable).," We conclude that no AM Her systems should exist beyond the period turning point (except the brown dwarf would be young, hot and still magnetic - searches for the secondary stars in AM Her systems near period mimimum would be very desirable)."
+ Tustead many of the unusual svsteuis at low orbital periods iuieht be desceudauts of the AM Er systeiis., Instead many of the unusual systems at low orbital periods might be descendants of the AM Her systems.
+ There is one paraueter. the streugth of he white dwarf magnetic field. which may decide which way the laer evolution goes.," There is one parameter, the strength of the white dwarf magnetic field, which may decide which way the later evolution goes."
+ Dwart ova systems do no partake iu these evolutionary phases., Dwarf nova systems do not partake in these evolutionary phases.
+ When thev lose nmuagnetic braking and settle to gravitational wave braking only hei lass transfer rate could well become too low for allowing any outburst (soe the marginal case of WZ See. Ilofiuciser et al.," When they lose magnetic braking and settle to gravitational wave braking only their mass transfer rate could well become too low for allowing any outburst (see the marginal case of WZ Sge, Meyer-Hofmeister et al."
+ 1998) iux thus face from view., 1998) and thus fade from view.
+ The ow viscosity needed to model the outbursts evcles of the dwarf nova systems in late secular evolution aud the disappearance of the AM Ier svsteuis could thus both result from the loss of the magnetic field of the companion star. when the star evolves to a cool brown clwart.," The low viscosity needed to model the outbursts cycles of the dwarf nova systems in late secular evolution and the disappearance of the AM Her systems could thus both result from the loss of the magnetic field of the companion star, when the star evolves to a cool brown dwarf."
+Compare to other accretion-powered. X-ray. pulsars. GX 1)4 has an atypically hard spectrum extending out well past 100 keV (FronteraanclDalFiume1989).,"Compared to other accretion-powered X-ray pulsars, GX 1+4 has an atypically hard spectrum extending out well past 100 keV \cite{fro89}."
+. Historically the spectrum aas been fitted with thermal bremsstrahlung or power law models: more recent observations with iniproved. spectral resolution generally. favour a power law model with ex»onential cutoll., Historically the spectrum has been fitted with thermal bremsstrahlung or power law models; more recent observations with improved spectral resolution generally favour a power law model with exponential cutoff.
+ Evpical values for the cutoll power law mocel parameters are photon index a=1.1.2.5: cutoll energy 5l8 keV: e-folding energv 11.26 keV. For any spectral model covering he range 1-10 keV. it is also necessary to include a gaussian component representing iron line emission at zz6.4 keV. and a term to account [or the elfects. of hotoelectrie: absorption by cold. gas along the line of. sight with hwdrogen column density in the range ng=(4140)l077cm7.," Typical values for the cutoff power law model parameters are photon index $\alpha = 1.1-2.5$; cutoff energy $5-18$ keV; $e$ -folding energy $11-26$ keV. For any spectral model covering the range 1-10 keV, it is also necessary to include a gaussian component representing iron line emission at $\approx 6.4$ keV, and a term to account for the effects of photoelectric absorption by cold gas along the line of sight with hydrogen column density in the range $n_H = (4-140) \times 10^{22}\,{\rm cm^{-2}}$."
+ The source spectrum. and in particular the column density ag have previously exhibited. significant variability on time-scales as short as a dav (Beckeretal.1976)., The source spectrum and in particular the column density $n_H$ have previously exhibited significant variability on time-scales as short as a day \cite{beck76}.
+. Measurements of spectral variation with phase are few: one example of pulse-phase spectroscopy was undertaken with data from the satellite [rom LOST and LOSS (Dotanietal.1989)., Measurements of spectral variation with phase are few; one example of pulse-phase spectroscopy was undertaken with data from the satellite from 1987 and 1988 \cite{dot89}.
+. Only the column density and the iron line centre energy were allowed to vary with phase in the spectral fits. and no significant variation was observed.," Only the column density and the iron line centre energy were allowed to vary with phase in the spectral fits, and no significant variation was observed."
+ The Ghosh and. Lamb (1979) model. predicts a correlation between torque ancl mass transfer rate (and hence luminosity) for aceretion-driven X-rav sources., The Ghosh and Lamb \shortcite{gl79} model predicts a correlation between torque and mass transfer rate (and hence luminosity) for accretion-driven X-ray sources.
+ For most sources it is difficult to test such a relationship since the range of luminosities at which they are observed is limited., For most sources it is difficult to test such a relationship since the range of luminosities at which they are observed is limited.
+ Llowever the correlation between torque ancl Luminosity has been confirmed. at least approximately. for three transient sources using data from the Burst anc Transient. Source Experiment (BATSE) aboard. the Compton Gamma. lav Observatory (CORO) andEEXOSAT (Reynolds ct αἱ.," However the correlation between torque and luminosity has been confirmed, at least approximately, for three transient sources using data from the Burst and Transient Source Experiment (BATSE) aboard the Compton Gamma Ray Observatory ) and (Reynolds et al."
+ 1996. Aldsten et al.," 1996, Bildsten et al."
+ 1997)., 1997).
+ The situation for persistent. pulsars is. however. less straightforward.," The situation for persistent pulsars is, however, less straightforward."
+ Phe BATSE cata have demonstrated that in general the torque is in facted with Luminosity in these sources., The BATSE data have demonstrated that in general the torque is in fact with luminosity in these sources.
+ The spin-up or spin-down rate can remain almost constant over intervals (referred. to in this paper as a constant torque state or just torque state). which are long compared to other characteristic time-scales of the system. even when the luminosity varies by several orders of magnitude over that time.," The spin-up or spin-down rate can remain almost constant over intervals (referred to in this paper as a `constant torque state' or just `torque state') which are long compared to other characteristic time-scales of the system, even when the luminosity varies by several orders of magnitude over that time."
+ Transitions between these states of constant torque can be abrupt. with time-scales of < 1d when the two torque values have the same sien: alternatively when switching from spin-up to spin-down (or vice-versa) the switch generally occurs smoothly over à period of 10.50 d. lt ds possible that there. remains some connection between the torque ancl luminosity. since at. times the torque measured for GX 114 has beenendicorretaled with luminosity (Chakrabartyctal.1997)..," Transitions between these states of constant torque can be abrupt, with time-scales of $< 1$ d when the two torque values have the same sign; alternatively when switching from spin-up to spin-down (or vice-versa) the switch generally occurs smoothly over a period of $10 - 50$ d. It is possible that there remains some connection between the torque and luminosity, since at times the torque measured for GX 1+4 has been with luminosity \cite{chak97b}."
+. This behaviour has not been observed in other pulsars., This behaviour has not been observed in other pulsars.
+ One important. caveat regarding the BATSE measurements is that the instrument can only measure pulsed. (ux., One important caveat regarding the BATSE measurements is that the instrument can only measure pulsed flux.
+ Systematic variations in pulse profiles or pulse fraction could. introduce significant aliasing to the flux data. hence masking the true relationship between bolometric Dux and torque.," Systematic variations in pulse profiles or pulse fraction could introduce significant aliasing to the flux data, hence masking the true relationship between bolometric flux and torque."
+ Given twt pulse profile shape and torque state have shown evidence! for correlation in GX 114 (Greenhill.GallowayandStorey998) this could potentially be an important οσο., Given that pulse profile shape and torque state have shown evidence for correlation in GX 1+4 \cite{gre98} this could potentially be an important effect.
+ 1n this paper we present results from s»eetral analysis of data obtained from GX 1|4 during 1996 using the Itossi X-ray Timing Explorer satellite(RAPE Ciies et al., In this paper we present results from spectral analysis of data obtained from GX 1+4 during 1996 using the Rossi X-ray Timing Explorer satellite; Giles et al.
+ 1995)., 1995).
+ A companion Oper (Cilesetal.L999) entains detailed analysis of pulse arrival times ancl pulse proile changes., A companion paper \cite{gil99} contains detailed analysis of pulse arrival times and pulse profile changes.
+ The source was observed withTE between 1996 July. 19 16:437 UT and 1996 July 21 02:89 UT., The source was observed with between 1996 July 19 16:47 UT and 1996 July 21 02:39 UT.
+ Several interruptions were made during that time as a consequence of previously scheduled monitoring of other sources., Several interruptions were made during that time as a consequence of previously scheduled monitoring of other sources.
+ After screening the data to avoid periods contaminated by. Earth occultations. the passage of the satellite through the South Atlantic Anomaly (SAA). ancl periods of unstable pointing. the total on-source duration was 51 ks.," After screening the data to avoid periods contaminated by Earth occultations, the passage of the satellite through the South Atlantic Anomaly (SAA), and periods of unstable pointing, the total on-source duration was $51$ ks."
+RAL consists of three instruments. the proportional counter array (PCA) covering the encrev range 2-60 keV. the high-cnereyv X-ray timing experiment (LIENTIS) covering 16-250 keV. ancl the all-sky monitor (ASAI) which spans 2-10 keV. Pointec observations are performed: using the PCA and. HIENTIS. instruments. while the ASM regularly scans the entire visible sky.," consists of three instruments, the proportional counter array (PCA) covering the energy range 2-60 keV, the high-energy X-ray timing experiment (HEXTE) covering 16-250 keV, and the all-sky monitor (ASM) which spans 2-10 keV. Pointed observations are performed using the PCA and HEXTE instruments, while the ASM regularly scans the entire visible sky."
+ The background-subtracted total PCA count rate for 3 of the five proportional counter units (PCUs) comprising the PCA is shown in Fig., The background-subtracted total PCA count rate for 3 of the five proportional counter units (PCUs) comprising the PCA is shown in Fig.
+ laa. The other two PCUs were only active briellv at the beginning of the observation so those data are not included in the analysis., \ref{fig1}a a. The other two PCUs were only active briefly at the beginning of the observation so those data are not included in the analysis.
+ Phe phase-averaged PCA count rate was initially low at &SOcounts+.," The phase-averaged PCA count rate was initially low at $\approx 80\,{\rm
+count\,s^{-1}}$."
+ This corresponds to a [ux of =6LOores+ in the 2-60 keV energy range. using the spectral model discussed in section 3 and assuming a source distance of 10 kpe.," This corresponds to a flux of $\approx 6 \times
+10^{36}\,{\rm erg\,s^{-1}}$ in the 2-60 keV energy range, using the spectral model discussed in section \ref{spec} and assuming a source distance of 10 kpc."
+ Throughout this paper we shall use this value as the source distance unless otherwise specified: the actual distance is thought to be in the range 3-15 kpe (ChakrabartyandRoche1997)., Throughout this paper we shall use this value as the source distance unless otherwise specified; the actual distance is thought to be in the range 3-15 kpc \cite{chak97}.
+. During the course of the observation the count rate decreased. to D. ⋅ ⋜↧⊔↓↓⊔↓⊔↓⊔⊔↓∪⇂≈⋅↱⊔∙⋯⋯⋡∖⋡≼∙∪↓⋅↓⋅≺⋅⋡∖↓≻∪⊔∠∐⊔⋏∙≟∣∪⋜↧∐⊔⇀∖∪⇂≈1 . before partially recovering towards the end.," During the course of the observation the count rate decreased to a minimum of $\approx 5\,{\rm count\,s^{-1}}$, corresponding to a flux of $\approx 4 \times 10^{35}\,{\rm erg\,s^{-1}}$ before partially recovering towards the end."
+ ''he count rates are unusually low for this source. with other observations giving significantly higher rates: for example zz320⋅−counts1 and 230counts5 (equivalent rates for 32 CLs) in February 1996 and January 1997 respectively.," The count rates are unusually low for this source, with other observations giving significantly higher rates; for example $\approx 320\,{\rm count\,s^{-1}}$ and $230\,{\rm
+count\,s^{-1}}$ (equivalent rates for 3 PCU's) in February 1996 and January 1997 respectively."
+ A imes the backeround-subtracted countrate curing interva 2 drops significantly below zero., At times the background-subtracted countrate during interval 2 drops significantly below zero.
+ This is a consequence of the ow source to background signal ratio (around. 1:10) during his interval coupled with statistical variations in the binnec countrate values., This is a consequence of the low source to background signal ratio (around 1:10) during this interval coupled with statistical variations in the binned countrate values.
+ The observation is divided into three intervals on the xwis of the mean flux (Fig., The observation is divided into three intervals on the basis of the mean flux (Fig.
+ laa)., \ref{fig1}a a).
+ Interval 1 covers the star of the observation to just before the Hux minimum., Interval 1 covers the start of the observation to just before the flux minimum.
+ Interya 2 spans the period of minimum flux. during which time the lux was =30counts+ apart from z10 s during a Hare (scc," Interval 2 spans the period of minimum flux, during which time the flux was $\la 30\,{\rm count\,s^{-1}}$ apart from $\approx 10$ s during a flare (see"
+is no visible area.,is no visible area.
+ In the opposite case. (he emission Irom (he wall of a disk seen edge-on is completely extinguished by (he disk κο.," In the opposite case, the emission from the wall of a disk seen edge-on is completely extinguished by the disk itself."
+ Thus. the range in cos? relevant for modeling is small and (he variation can be neglected.," Thus, the range in $\cos i$ relevant for modeling is small and the variation can be neglected."
+ In (he present model. the height of the wall is taken as an arbitrary free parameter that can be adjusted to compensate lor any. variation ol ;.," In the present model, the height of the wall is taken as an arbitrary free parameter that can be adjusted to compensate for any variation of $i$."
+ Models showing changes in / aud  are presented in Table 5. as Series 7./h.," Models showing changes in $i$ and $h$ are presented in Table \ref{table-models} as Series $i,h$."
+ A complete analvsis of all the possibilities that describe the dust. composition of the wall is computationally demanding. so we decided to explore the changes on the SED Lor ivpical cases.," A complete analysis of all the possibilities that describe the dust composition of the wall is computationally demanding, so we decided to explore the changes on the SED for typical cases."
+ All these models ean be used as a starting point to tackle future mid-IBR. spectra modeling with observations of new binary (targets., All these models can be used as a starting point to tackle future mid-IR spectra modeling with observations of new binary targets.
+" Figure 23. shows the effect on the SED of changing the twpe of silicates. [or e=0.60 0.5.a— 0.cosi—0.5. Ru,=1.τα and h=0.28a."," Figure \ref{fig-4Myr-pyr-oliv} shows the effect on the SED of changing the type of silicates, for $e=0$ $\phi=0.5$ $\alpha=0$ $\cos i=0.5$ $R_{cb}=1.7a$ and $h=0.28a$."
+ We show results for the fiducial model. and for the case where the silicates component are changed οι pyroxenes GUgiFeu5704) to olivines GUgFe S/O1). model POI in Table 5. and the only member of Series pyr.oliv.," We show results for the fiducial model, and for the case where the silicates component are changed from pyroxenes $Mg_{0.8}\,Fe_{0.2}\,SiO_{3}$ ) to olivines $Mg\,Fe\,SiO_{4}$ ), model PO1 in Table \ref{table-models} and the only member of Series pyr,oliv."
+ The flux from the wall is larger when the dust is made by olivines (model POL) than when 1 is made by pyroxenes (fiducial model)., The flux from the wall is larger when the dust is made by olivines (model P01) than when it is made by pyroxenes (fiducial model).
+ This can be explained. because the former model has a higher temperature at the same distance from the central star.," This can be explained, because the former model has a higher temperature at the same distance from the central star."
+ Another difference is that the central wavelength of the LOjan feature due to olivines is slightly shifted to longer wavelengths than the feature for the pyroxenes., Another difference is that the central wavelength of the $10\mu$ m feature due to olivines is slightly shifted to longer wavelengths than the feature for the pyroxenes.
+" Figure 24. shows the effect on the SED of changing the twpe of carbon. lor e=0.0 0.5.a= 0.cosi;=0.5. Ra,=1.τα and h=0.28a."," Figure \ref{fig-4Myr-pyr-graf-carb} shows the effect on the SED of changing the type of carbon, for $e=0$ $\phi=0.5$ $\alpha=0$ $\cos i=0.5$ $R_{cb}=1.7a$ and $h=0.28a$."
+ We show results for the case where we take amorphous carbon (Mathis&Whillen1989) and graphite (Weigartner&Draine2001) erains., We show results for the case where we take amorphous carbon \citep{Mathis3} and graphite \citep{Weingartner} grains.
+ First. we compare (wo cases: a model with no graphite nor amorphous carbon (1.e.. only pyroxene grains in the wall. model PAL) and a model like the fichicial one. but where eraphite has been replaced by amorphous carbon (model PA2).," First, we compare two cases: a model with no graphite nor amorphous carbon (i.e., only pyroxene grains in the wall, model PA1) and a model like the fiducial one, but where graphite has been replaced by amorphous carbon (model PA2)."
+" We can see that model PAL has a lower flux than model DÀ2. due to a lower 75,5. as noted in Table 5.."," We can see that model PA1 has a lower flux than model PA2, due to a lower $T_{wall}$, as noted in Table \ref{table-models}."
+ The presence of amorphous carbon in model PÁ2 results in a larger temperature., The presence of amorphous carbon in model PA2 results in a larger temperature.
+ For the same abundances. the absorption coellicients for graphite and amorphous carbon are similar between 0.3 and 3 microns. wilh the graphite opacity slightly larger.," For the same abundances, the absorption coefficients for graphite and amorphous carbon are similar between 0.3 and 3 microns, with the graphite opacity slightly larger."
+ Outside (his interval. the amorphous carbon opacitv is higher.," Outside this interval, the amorphous carbon opacity is higher."
+ Thus. changes in the SED (see Figure 24)) are expected due to the differences in (he opacity for both components.," Thus, changes in the SED (see Figure \ref{fig-4Myr-pyr-graf-carb}) ) are expected due to the differences in the opacity for both components."
+ Also in Figure 24.. models with variations on the abundance of graphite ancl amorphous carbon are also analvzed. for e= 0.6= 0.5.a= O.cos?=0.5. Ry=L.fa and h=0.284.," Also in Figure \ref{fig-4Myr-pyr-graf-carb}, models with variations on the abundance of graphite and amorphous carbon are also analyzed, for $e=0$ $\phi=0.5$ $\alpha=0$ $\cos i=0.5$, $R_{cb}=1.7a$ and $h=0.28a$."
+ Results are shown lor the cases where there is no amorphous carbon ancl (he graphite abundance is halved ancl doubled (models PA3 and DÀ4 in table 5))., Results are shown for the cases where there is no amorphous carbon and the graphite abundance is halved and doubled (models PA3 and PA4 in table \ref{table-models}) ).
+ In another case. the graphite is replaced by amorphous carbon: model PAS corresponds to half the adopted standard value for amorphous carbon abundance anc PAG to the doubled value.," In another case, the graphite is replaced by amorphous carbon; model PA5 corresponds to half the adopted standard value for amorphous carbon abundance and PA6 to the doubled value."
+ For graphite, For graphite
+"1n summary our simulations have shown that in order to niaximize the ratio of Virgo members to background galaxies we should use a selection. criteria of (54,=23 and a scale length f z 3 aresec.",In summary our simulations have shown that in order to maximize the ratio of Virgo members to background galaxies we should use a selection criteria of $\mu_{o} \geq 23$ and a scale length $h$ $\geq$ 3 arcsec.
+ On a practical note we found that typically the seeing on our frames was à poor 2 arc sec., On a practical note we found that typically the seeing on our frames was a poor 2 arc sec.
+ Convolving a 3 are sec scale size galaxy with this seeing leads to à measured scale size closer to 4 are sec., Convolving a 3 arc sec scale size galaxy with this seeing leads to a measured scale size closer to 4 arc sec.
+ Thus in practice we selected objects with measured scale sizes greater than 53 are sec which. because of the template sizes used in the method. described below. means a minimum scale size of 4 arc sec.," Thus in practice we selected objects with measured scale sizes greater than 3 arc sec which, because of the template sizes used in the method described below, means a minimum scale size of 4 arc sec."
+ As a further check to validate our selection criteria. we have applied our. algorithm (see sec. 22))," As a further check to validate our selection criteria, we have applied our algorithm (see sec. \ref{sec:algo}) )"
+ to another set of INP WEC data using the same detector. exposure time and [filter and covering a region of skv at about the same Calactic latitude.," to another set of INT WFC data using the same detector, exposure time and filter and covering a region of sky at about the same Galactic latitude."
+" The cata we used are part of the Millennium Galaxy Survey that is a 36 aremin wide strip eoing from O58""28 to 1H446""45. J2000 (Liske et al.."," The data we used are part of the Millennium Galaxy Survey that is a 36 arcmin wide strip going from $9^{h}58^{m}28^{s}$ to $14^{h}46^{m}45^{s}$, J2000 (Liske et al.,"
+ 2002)., 2002).
+ From these data we chose a number of random fields to compare our predicted number of background detections with that of the model., From these data we chose a number of random fields to compare our predicted number of background detections with that of the model.
+ The number of galaxies. detected bv our algorithm is 4 galaxies per sq deg., The number of galaxies detected by our algorithm is 4 galaxies per sq deg.
+ This number is in agreement both with what our numerical simulations predicted (6 gal per sq deg) and from what we measure for the background. counts to be in the Virgo fields., This number is in agreement both with what our numerical simulations predicted $\sim6$ gal per sq deg) and from what we measure for the background counts to be in the Virgo fields.
+ Lt is much less than the actual number of galaxies ( 720 gal/ sq deg) found in the Virgo cluster fields., It is much less than the actual number of galaxies ( $\approx 20$ gal/ sq deg) found in the Virgo cluster fields.
+ ALL these points are discussed further below., All these points are discussed further below.
+ Looking for objects with a surface Lux close to the sky noise requires the use of image enhancement: techniques in order to optimize their detection., Looking for objects with a surface flux close to the sky noise requires the use of image enhancement techniques in order to optimize their detection.
+ Standard: detection algorithms for connected. pixels in fact. often fail on these kind of objects because of their poor. pixel-to-pixel signal to noise ratio (S/N)., Standard detection algorithms for connected pixels in fact often fail on these kind of objects because of their poor pixel-to-pixel signal to noise ratio (S/N).
+ The algorithm we have developed uses frequency domain techniques ancl mainly consists of convolutions of the image with matched Lilters: the advantage in this case is the use of the total flux of the ealaxw to detect it. instead. of the very low S/N pixels at," The algorithm we have developed uses frequency domain techniques and mainly consists of convolutions of the image with matched filters: the advantage in this case is the use of the total flux of the galaxy to detect it, instead of the very low S/N pixels at"
+For example. a smaller core scatter will probably produce more 3c outliers than a larger core scatter.,"For example, a smaller core scatter will probably produce more $3\sigma$ outliers than a larger core scatter."
+ With no alternative at hand to condense the performance of one particular setup into a handful of numbers we can only refer to the zy VS. Spee plots shown in the following which give an uncompressed view of the data., With no alternative at hand to condense the performance of one particular setup into a handful of numbers we can only refer to the $z_{\mathrm{phot}}$ vs. $z_{\mathrm{spec}}$ plots shown in the following which give an uncompressed view of the data.
+ The VVDS ts complete down to {αμ=24 which corresponds to Ίνα=423.5., The VVDS is complete down to $I_\mathrm{AB}=24$ which corresponds to $I_\mathrm{Vega}\approx23.5$.
+ Thus. we decided to asses the photo-z accuracy for all objects with 17<R23 to achieve a reasonable level of completeness.," Thus, we decided to asses the $z$ accuracy for all objects with $17<R<23$ to achieve a reasonable level of completeness."
+ We also present the results for a fainter magnitude bin with objects in the range 23<R24 which are then possibly biased in terms of selection., We also present the results for a fainter magnitude bin with objects in the range $23<R<24$ which are then possibly biased in terms of selection.
+ The mean redshifts in the bright and the faint bins are z=0.55 and z=0.75. respectively.," The mean redshifts in the bright and the faint bins are $z=0.55$ and $z=0.75$, respectively."
+" The different setups are named with three-letter acronyms with the first letter denoting the code (""H™ forHyperz.. ""C7 for the COMBO code. and ""B"" for BPZ)). the second letter denoting the dataset (CB"" for GaBoDS and ""C"" for COMBO). and the digit at the third position denoting the filter set (757 for UBVRI. ""A"" for ""BVRI. and 7177 for the full COMBO-17 filter set including medium-band-filters)."," The different setups are named with three-letter acronyms with the first letter denoting the code (“H” for, “C” for the COMBO code, and “B” for ), the second letter denoting the dataset (“B” for GaBoDS and “C” for COMBO), and the digit at the third position denoting the filter set (“5” for $UBVRI$, “4” for “BVRI”, and “17” for the full COMBO-17 filter set including medium-band-filters)."
+ Since the FDF data are extremely deep and subtle trends in photometric errors make little difference to the = quality. we do not split the FDF sample into magnitude bins.," Since the FDF data are extremely deep and subtle trends in photometric errors make little difference to the $z$ quality, we do not split the FDF sample into magnitude bins."
+ Furthermore. given the selection choices made for the spectroscopic sample it is not complete at or representative for any particular magnitude limit.," Furthermore, given the selection choices made for the spectroscopic sample it is not complete at or representative for any particular magnitude limit."
+ Hence. we split the FDF spectroscopic catalogue into two samples at z=2 to show the effects of different filter sets and especially NIR bands on the performance at low and high redshift in comparison.," Hence, we split the FDF spectroscopic catalogue into two samples at $z=2$ to show the effects of different filter sets and especially NIR bands on the performance at low and high redshift in comparison."
+" The FDF setups are denoted by a second letter ""F for FDF and the filter set is spelled out.", The FDF setups are denoted by a second letter “F” for FDF and the filter set is spelled out.
+ In the following we report the results from our blind test of photo-z performance for the different setups., In the following we report the results from our blind test of $z$ performance for the different setups.
+ As a complete coverage of all possible data-code-parameter combinations would be beyond the scope of this paper we concentrate on some well-chosen setups to illustrate the effects of key parameters., As a complete coverage of all possible data-code-parameter combinations would be beyond the scope of this paper we concentrate on some well-chosen setups to illustrate the effects of key parameters.
+ The statisties for all photo-z setups that are compared to the VVDS spectroscopic catalogue are shown in Table 2. and Table 3..," The statistics for all $z$ setups that are compared to the VVDS spectroscopic catalogue are shown in Table \ref{tab:res_VVDS_COMBO}
+ and Table \ref{tab:res_VVDS_GaBoDS}."
+ Selected setups are also illustrated in Fig., Selected setups are also illustrated in Fig.
+ 5 and Fig., \ref{fig:zz_BCH_C5} and Fig.
+ by plots showing photo-z versus spectroscopic redshift.," \ref{fig:zz_BCH_B5}
+ by plots showing $z$ versus spectroscopic redshift."
+ Clearly. the COMBO code performs best in comparison to the two other codes with the 17-filter set as well as with the 4- and 5-filter sets.," Clearly, the COMBO code performs best in comparison to the two other codes with the 17-filter set as well as with the 4- and 5-filter sets."
+ The COMBO code shows the expected behaviour that the photo-z accuracy decreases when further filters are excluded., The COMBO code shows the expected behaviour that the $z$ accuracy decreases when further filters are excluded.
+ Completeness decreases while outlier rate and scatter increase., Completeness decreases while outlier rate and scatter increase.
+ No large biases are produced in any setup., No large biases are produced in any setup.
+ Owing to their greater depth the GaBoDS data mostly lead to better results than the shallower COMBO data. in UBVRI as well as in BVRI. As expected. the effect is much more pronounced in the faint bin. while the depth helps less in the estimation of redshifts for high S/N objects at the bright end of our catalogue.," Owing to their greater depth the GaBoDS data mostly lead to better results than the shallower COMBO data, in $UBVRI$ as well as in $BVRI$ As expected, the effect is much more pronounced in the faint bin, while the depth helps less in the estimation of redshifts for high $S/N$ objects at the bright end of our catalogue."
+ Nevertheless. also bright objects with >2007 detections in the R-band can benefit from the depth in the other bands.," Nevertheless, also bright objects with $>20\sigma$ detections in the $R$ -band can benefit from the depth in the other bands."
+ The negative biases in the photo-z estimation with and is also present in GaBoDS setups with the U-band included., The negative biases in the $z$ estimation with and is also present in GaBoDS setups with the $U$ -band included.
+At this point. it is important to mentionagain that the,"At this point, it is important to mentionagain that the"
+which was computed analytically.,which was computed analytically.
+ AC L. the faint source contribution in each region is the 2.2 ya value multiplied by the calibration ratio of 0.491.," At L, the faint source contribution in each region is the 2.2 $\mu$ m value multiplied by the calibration ratio of 0.491."
+" These E, are listed in Tables 1.2 and 3 for K. J and L. An uncertainty of of the total prediction is assigned to this correction. and is listed in Table 5 under “Faint Source.”"," These $_q$ are listed in Tables \ref{ktable}, \ref{jtable} and \ref{ltable} for K, J and L. An uncertainty of of the total prediction is assigned to this correction, and is listed in Table \ref{error} under “Faint Source.”"
+ The CIRD in each region is then (he derived intercept. DZ(0). minus the faint star contribution Fy. C = DZ(0) - F. These values are also listed in Tables 1.. 2. and 3..," The CIRB in each region is then the derived intercept, DZ(0), minus the faint star contribution $_q$, C = DZ(0) - F. These values are also listed in Tables \ref{ktable}, \ref{jtable} and \ref{ltable}."
+ The mean of these CIRD estimates are 14.59 c 0.14. 8.83 4 0.51 and 15.57 dE 0.20 kJv ! for Ix. J and L. These standard deviations of the means are listed in Table !-) as “Scatter.”," The mean of these CIRB estimates are 14.59 $\pm$ 0.14, 8.83 $\pm$ 0.51 and 15.57 $\pm$ 0.20 kJy $^{-1}$ for K, J and L. These standard deviations of the means are listed in Table \ref{error} as “Scatter.”"
+ PATASS magnitudes at IX and J were converted into fluxes using ΕΝ) = 614 Jy and F.(J) = 1512 Jv which were derived by Wright(2001) ancl Gorjianetal.(2000)., 2MASS magnitudes at K and J were converted into fluxes using $_\circ$ (K) = 614 Jy and $_\circ$ (J) = 1512 Jy which were derived by \citet{elw01} and \citet{gor00}.
+. However. the derived slopes at IX and. J indicate DIRBE fluxes for a zero magnitude 2MASS star of 540 and 1467 Jv respectively.," However, the derived slopes at K and J indicate DIRBE fluxes for a zero magnitude 2MASS star of 540 and 1467 Jy respectively."
+ The ratio of the calibration factors atReese 3.5 and 2.2 jam is 0.491. consistent with those found by Dwek&Arendt(1998).. Wright&(2000) and &Johnson (2001).," The ratio of the calibration factors at 3.5 and 2.2 $\mu$ m is 0.491, consistent with those found by \citet{dwe98}, \citet{wrr00} and \citet{wrj01}."
+. Uncertainties in (he CIRD due to calibration error were estimated using the change in the median DZ(0) when the slopes. &. were forced to change by &5% at J and L. or X105 at K due to the large difference (Wright&Johnson2001) between the fitted value of 0.38 and the expected value of 1.," Uncertainties in the CIRB due to calibration error were estimated using the change in the median DZ(0) when the slopes, $\kappa$, were forced to change by $\pm$ at J and L, or $\pm$ at K due to the large difference \citep{wrj01} between the fitted value of 0.88 and the expected value of 1."
+" The change in the medians are x 2.24. 1.77. and i0 kJv tat IX. J and L respectively and are listed in Table 5. under ""Calibration."""," The change in the medians are $\mp$ 2.24, 1.77, and 0.60 kJy $^{-1}$ at K, J and L respectively and are listed in Table \ref{error} under “Calibration.”"
+ There is a small correction for faint galaxies (hat appear in the 2AIASS PSC., There is a small correction for faint galaxies that appear in the 2MASS PSC.
+ These have been subtracted along with the Galactic stars. but should be included in the CIRB.," These have been subtracted along with the Galactic stars, but should be included in the CIRB."
+ Wright(2001) estimates this correction is 0.1 and 0.05 kJv tat 2.2 and 1.25 yan. The 0.1 kJv | correction at 2.2 jan implies a 0.05 kJv |1 correction at 3.5 jan due to the relative calibration [actor of 0.491., \citet{elw01} estimates this correction is 0.1 and 0.05 kJy $^{-1}$ at 2.2 and 1.25 $\mu$ m. The 0.1 kJy $^{-1}$ correction at 2.2 $\mu$ m implies a 0.05 kJy $^{-1}$ correction at 3.5 $\mu$ m due to the relative calibration factor of 0.491.
+ These corrections have been added back after taking the mean of C = DZ(0) - F in the 40 regions., These corrections have been added back after taking the mean of C = DZ(0) - F in the 40 regions.
+ Thus. the final reported. values of the CIRD are ihe mean of the values Cin Tables 1.. 2 and 3. plus this correction.," Thus, the final reported values of the CIRB are the mean of the values C in Tables \ref{ktable}, \ref{jtable} and \ref{ltable} plus this correction."
+ An uncertainty of of this correction is included in Table 5 under “Galaxies.”, An uncertainty of of this correction is included in Table \ref{error} under “Galaxies.”
+ Gorjianοἱal.(2000) adopt an uncertainty of of the zodiacal intensity at the ecliptic poles., \citet{gor00} adopt an uncertainty of of the zodiacal intensity at the ecliptic poles.
+ These uncertainties are listed in Table 5 under ~Zodiacal.”, These uncertainties are listed in Table \ref{error} under “Zodiacal.”
+ After adding errors in quadrature. we obtain a CIRB of 14.69 + 449 kJy | at2.2 ja. a weak limit of 8.88 + 6.26 kJy tat 125 pan and a CIRB of 15.62 4 3.34 kJv t αἱ 3.5 yan. Figures 4.. 5.. and 6. show histograms of the residuals DZ;—&D;DZ(0) lor all 2971 pixels in IX. J and L with interquartile ranges of2.85. 3.57 and 2.11 respectively.," After adding errors in quadrature, we obtain a CIRB of 14.69 $\pm$ 4.49 kJy $^{-1}$ at 2.2 $\mu$ m, a weak limit of 8.88 $\pm$ 6.26 kJy $^{-1}$ at 1.25 $\mu$ m and a CIRB of 15.62 $\pm$ 3.34 kJy $^{-1}$ at 3.5 $\mu$ m. Figures \ref{kres}, \ref{jres}, and \ref{lres} show histograms of the residuals $DZ_i- \kappa B_i-DZ(0)$ for all 2971 pixels in K, J and L with interquartile ranges of 2.85, 3.57 and 2.11 respectively."
+ Divicine these residuals by o(2;)Dj al each pixel gives a non-Gaussian distribution which is tightlySL packed near zero wilh a few pixels extending out to large values., Dividing these residuals by $\sigma(B_i)$ at each pixel gives a non-Gaussian distribution which is tightly packed near zero with a few pixels extending out to large values.
+ The number of pixels with, The number of pixels with
+Llicrarchical. aggregation. seems to be at the heart of. galaxy evolution.,Hierarchical aggregation seems to be at the heart of galaxy evolution.
+ In a cold dark matter universe. as depicted by numerical simulations. its structure grows. through subsequent. mergers ancl zero fragmentations.," In a cold dark matter universe, as depicted by numerical simulations, its structure grows through subsequent mergers and zero fragmentations."
+ The growth and evolution of galaxies. which are thought to use dark matter as scalfolding. is channeled through this hierarchical aggregation. at least for the most massive structures (Springeletal.2006).," The growth and evolution of galaxies, which are thought to use dark matter as scaffolding, is channeled through this hierarchical aggregation, at least for the most massive structures \citep{LSS_SFW}."
+. Notwithstanding all the complexity in the process of ealaxy formation and evolution. galaxies still are the most basic. population. unit MENin the description.2. of. large scale structure in the Universe.," Notwithstanding all the complexity in the process of galaxy formation and evolution, galaxies still are the most basic population unit in the description of large scale structure in the Universe."
+ And. still nowadays much work is being invested in galaxy formation to disentangle the influence of the hierarchical context setup by dark matter rom the secular barvonic processes on small scales., And still nowadays much work is being invested in galaxy formation to disentangle the influence of the hierarchical context setup by dark matter from the secular baryonic processes on small scales.
+ ‘Tackling the problem theoretically implies numerical experiments following large structure dynamics and. at ie same time. a description. of baryonic processes. such wdrodynamies and radiative cooling.," Tackling the problem theoretically implies numerical experiments following large structure dynamics and, at the same time, a description of baryonic processes such hydrodynamics and radiative cooling."
+ This is still. very hallenging in the usual numerical approach that discretizes space and time. and try to solve a relevant. set of equations ο capt the physics (Abeletal.3002:Got 2006).. u," This is still very challenging in the usual numerical approach that discretizes space and time, and try to solve a relevant set of equations to capture the physics \citep{2002Sci...295...93A,2006AIPC..878....3G}."
+rFrom the computational point of view it oberinvolves achievingS6 an elfective∙⋠ resolution| spanning. at least 5 orders of noonmagnitude Mqin mass.κ ne κand lengthtimeime (Norman(Norms:et4albal.2OO0T2007).," From the computational point of view it involves achieving an effective resolution spanning at least $5$ orders of magnitude in mass, length and time \citep{2007arXiv0705.1556N}."
+. To overcome this barrier the semi-analvtic moclel (hereafter SAM) approach to describe proposesfirst the non-Ιinear clustering[ustο of clarkdark matterει on lglarge scales.Ι and Lddescribe| later the small scale barvonic physics through5 analytical prescriptions.," To overcome this barrier the semi-analytic model (hereafter SAM) approach proposes to describe first the non-linear clustering of dark matter on large scales, and describe later the small scale baryonic physics through analytical prescriptions."
+ The connection between the two. scales. 18 provided through the dark matter halo. which is the most," The connection between the two scales is provided through the dark matter halo, which is the most"
+addition. a faint structure appears ahead. of NGC 2208s motion. which may be the leading tail.,"addition, a faint structure appears ahead of NGC 2298's motion, which may be the leading tail."
+ Phe two structures found at the extreme North and. South most. likely are boundary effect: introduced. by the smoothing process., The two structures found at the extreme North and South most likely are boundary effect introduced by the smoothing process.
+ At last. a strong Northeast structure appears in the direction perpencdicuar to the Gaactic disk.," At last, a strong Northeast structure appears in the direction perpendicular to the Galactic disk."
+ This may be interpreted as stars tha have left he cluster although have not had time to fall behind. or ahead of the orbit., This may be interpreted as stars that have left the cluster although have not had time to fall behind or ahead of the orbit.
+ Another [aint structure opposite to the previous one is present. although not connected to any other enhancement.," Another faint structure opposite to the previous one is present, although not connected to any other enhancement."
+ Our fincings are similar to those predicted by Combes.Leon.&Aevlan (1999)., Our findings are similar to those predicted by \citet{combes}.
+. The authors simulations preclict a formation of multiple perpendicular tails in a cross-like pattern resulting from multiple disk crossines., The authors simulations predict a formation of multiple perpendicular tails in a cross-like pattern resulting from multiple disk crossings.
+ We do not discard the possibility that some of the features detected. are. in [act due to wrong extinction corrections., We do not discard the possibility that some of the features detected are in fact due to wrong extinction corrections.
+ They could result from hieh [reeuenevy structures on the dust filaments close to the disk., They could result from high frequency structures on the dust filaments close to the disk.
+ These structures are not properly sampled: due to. resolution limitations ofthe Schlegel.Finkbeiner.&Marc(1998). maps built using LRAS that has a £M1ΑΛ=6.1., These structures are not properly sampled due to resolution limitations of the \citet{schlegel} maps built using IRAS that has a $FWHM = 6\arcmin.1$.
+ However. most of the structures seen in figure S. extend along many. HUAS EWLLMs and tjus are likely to be real.," However, most of the structures seen in figure \ref{2298tail} extend along many IRAS FWHMs and thus are likely to be real."
+ So far. the AIF technique has been applied without any assessment of the influence of improperly builtfy.," So far, the MF technique has been applied without any assessment of the influence of improperly built."
+ That is. i£ does not reflect the distribution of cluster stars in the colour-magnitude space. as well as its variations. throughout the entire FOV. the application of the ME may lead. to miss identifications of tical structures.," That is, if does not reflect the distribution of cluster stars in the colour-magnitude space, as well as its variations, throughout the entire FOV, the application of the MF may lead to miss identifications of tidal structures."
+ One phenomenon that may give rise to an improper is mass segregation., One phenomenon that may give rise to an improper is mass segregation.
+ We here attempt to quantify its effect on a model cluster consistent with NGC 2298., We here attempt to quantify its effect on a model cluster consistent with NGC 2298.
+ ‘To assess what is the effect of mass segregation on the ALP technique. we made another simulation of a GC similar to NGC 2298. but now using the present clay mass function (PDME) from the literature.," To assess what is the effect of mass segregation on the MF technique, we made another simulation of a GC similar to NGC 2298, but now using the present day mass function (PDMF) from the literature."
+ Since the slope of the PDME is only determined to a distance of 1.8’. we extrapolate it to 3.67 by assuming the sume growth rate of the PDME slope as in the inner parts of NGC 2298.," Since the slope of the PDMF is only determined to a distance of $1.8\arcmin$, we extrapolate it to $3.6\arcmin$ by assuming the same growth rate of the PDMF slope as in the inner parts of NGC 2298."
+ For the outermost parts of the cluster. we assume tiu the PDME slope saturates at the value at 3.67.," For the outermost parts of the cluster, we assume that the PDMF slope saturates at the value at $3.6\arcmin$."
+ This later value. therefore. is the one that applies to most extra-tidal stars in the model.," This latter value, therefore, is the one that applies to most extra-tidal stars in the model."
+ ‘Table l lists the model slopes at different distances from cluster centre., Table \ref{slopes} lists the model slopes at different distances from cluster centre.
+ We also run another simulation. which is identical to the previous one. but wihout mass segregation.," We also run another simulation, which is identical to the previous one, but without mass segregation."
+ In this second. non-scereeatecl case. the PDAME slope used is the one corresponding to the ouermost bin in Table 1..," In this second, non-segregated case, the PDMF slope used is the one corresponding to the outermost bin in Table \ref{slopes}."
+ The simulation parameters and number of simulated stars in these two extra simulations are the same as in $3.1... except for the PDME slopes. as described.," The simulation parameters and number of simulated stars in these two extra simulations are the same as in \ref{sec:simulations}, except for the PDMF slopes, as described."
+ The background of stars is also the same one used in that section. built using," The background of stars is also the same one used in that section, built using."
+TriLegal.. In figure 9. we show the comparison of the true number of cluster stars to those detected. using the ME. in. cases when mass-sceregation is present ancl absent.," In figure \ref{finalsim} we show the comparison of the true number of cluster stars to those detected using the MF, in cases when mass-segregation is present and absent."
+ Notice that there are no systematic changes in the number of stars obtained in either situation., Notice that there are no systematic changes in the number of stars obtained in either situation.
+ However there is a significant difference on the dispersion. in that the ME reconstructed star counts have larger scatter in the cluster which is subject to mass segregation.," However there is a significant difference on the dispersion, in that the MF reconstructed star counts have larger scatter in the cluster which is subject to mass segregation."
+ We thus conclude that. even. though 16 detection of a tidal tail is still possible in. presence of mass segregation ellects. the limiting distance out to which 1e tail may. be detected. as well as some of its Iow-density substructure. may be alfected if does not properly take segregation into account.," We thus conclude that, even though the detection of a tidal tail is still possible in presence of mass segregation effects, the limiting distance out to which the tail may be detected, as well as some of its low-density substructure, may be affected if does not properly take segregation into account."
+part of the disk the cloud. assumed not uniform in temperature and clensily. becomes more and more compact. following the inward motion of the disk edge.,"part of the disk the cloud, assumed not uniform in temperature and density, becomes more and more compact, following the inward motion of the disk edge."
+ Initially the cloud is large and marginally thick (ο electron-seattering., Initially the cloud is large and marginally thick to electron-scattering.
+ Photons generated by the disk are upscattered in (he inner and hotter regions. emerging with a positive time lag.," Photons generated by the disk are upscattered in the inner and hotter regions, emerging with a positive time lag."
+ In later stages. the cloud is much more dense. with the consequence (hat in ils central regions photons reach a Bose-Einstein equilibrium. and are then downscattered when they propagate through the outer and cooler shells.," In later stages, the cloud is much more dense, with the consequence that in its central regions photons reach a Bose-Einstein equilibrium, and are then downscattered when they propagate through the outer and cooler shells."
+ In this latter configuration lags turn out to be negalive., In this latter configuration lags turn out to be negative.
+ Their calculations succeeded (ο reproduce some of the features of GRS 1915-105. but cannot explain. e.g.. the apparent alternating negative and positive phase lags between ihe harmonics of the QPOs observed in some microquasars (Wijnandsetal.1999).," Their calculations succeeded to reproduce some of the features of GRS 1915+105, but cannot explain, e.g., the apparent alternating negative and positive phase lags between the harmonics of the QPOs observed in some microquasars \citep{wij99}."
+. In (his paper we propose a simple alternative model. showing (hat also dvnamical Conmptonization by the jet on the radiation emitted by an accretion disk can account for some of the observed properties of the source.," In this paper we propose a simple alternative model, showing that also dynamical Comptonization by the jet on the radiation emitted by an accretion disk can account for some of the observed properties of the source."
+ The model is alternative also to other interpretations of the (üming properties of microquasars based on dvnamical Conmptonization effects in converging [lows (LaurentandTitarchuk2001)., The model is alternative also to other interpretations of the timing properties of microquasars based on dynamical Comptonization effects in converging flows \citep{lau01}.
+. The geometrical setting of our model is discussed in (he next section. while in section 32 we present a imunerical simulation based on a Monte Carlo technique.," The geometrical setting of our model is discussed in the next section, while in section 3 we present a numerical simulation based on a Monte Carlo technique."
+ One of the defining characteristics of microquasars is (he fact (hat they exhibit relativistic ejection of matter., One of the defining characteristics of microquasars is the fact that they exhibit relativistic ejection of matter.
+ From the observations of superluminal expansion in GRS 1915-105. Fender(1999) inferred a bulk ejecta velocity 9]=v/e20.95 and a mass outflow rate mnmlOestsMaga.," From the observations of superluminal expansion in GRS 1915+105, \citet{fen99} inferred a bulk ejecta velocity $\,\beta=v/c \gtrsim 0.95\, $ and a mass outflow rate $\dot m \gtrsim 10^{18}\, {\rm g\, s^{-1}}\approx {\dot M}_{Edd} $."
+ The question of what the accelerating mechanism really is. is still unanswered.," The question of what the accelerating mechanism really is, is still unanswered."
+ The most probable hypothesis is that the plasma ejection is driven and collimated by magnetic fields anchored to the disk (see. e.g.. HeinzandBegelman (2000))).," The most probable hypothesis is that the plasma ejection is driven and collimated by magnetic fields anchored to the disk (see, e.g., \citet[]{hei00}) )."
+ This implies (hat the accelerating engine should be effective near the central black hole. (hus making reasonable the assumption that the plasmoids do attain their maximum velocity already al a few gravitational radii.," This implies that the accelerating engine should be effective near the central black hole, thus making reasonable the assumption that the plasmoids do attain their maximum velocity already at a few gravitational radii."
+ A rough calculation shows that the inner portion of a jel with an opening angle Ο20.2 radians becomes optically thick to electron scattering for in=ο we then argue that during the active state a small but non-negligible fraction of the photons emitted by the disk is scattered by (he electrons on the jet.," A rough calculation shows that the inner portion of a jet with an opening angle $\Theta_J\approx 0.2 $ radians becomes optically thick to electron scattering for $\,\dot m \gtrsim 0.1\, {\dot M}_{Edd}$, we then argue that during the active state a small but non-negligible fraction of the photons emitted by the disk is scattered by the electrons on the jet."
+ H should be noted (hat. whether a photon acquires or loses energy. depends on the angles of scattering and ullimately. on (he posiüion of the source. the height where scattering occurs and the inclination with respect to the observer.," It should be noted that, whether a photon acquires or loses energy depends on the angles of scattering and ultimately, on the position of the source, the height where scattering occurs and the inclination with respect to the observer."
+ As an example. let us. consider," As an example, let us consider"
+initial nature of the two populations of plasma and the shape of the loss-cone. re. for the hot plasma: particle density ip. particle temperature 73: for the cold plasma: particle density ης. particle temperature Το: for the loss-cone: loss-cone angle a. pitch angle distribution slope N.,"initial nature of the two populations of plasma and the shape of the loss-cone, i.e. for the hot plasma: particle density $n_{\rm h}$, particle temperature $T_{\rm h}$; for the cold plasma: particle density $n_{\rm c}$ , particle temperature $T_{\rm c}$; for the loss-cone: loss-cone angle $\alpha_{\rm c}$ , pitch angle distribution slope $N$."
+" Rather than choowing values of the cold plasma density ης we fix its value via adoption of a value for uv=f;/ the ratio of plasma frequency (f,=9x10201) MHz) fk.and gyrofrequency (Cf.«2.86xB MHz) where B is the magnetic field strength in G. In the quasi-linear diffusion process. the existence of background electromagnetic wave energy Ty is taken into account."," Rather than choowing values of the cold plasma density $n_{\rm c}$ we fix its value via adoption of a value for $u=f_{\rm p}/f_{\rm c}$, the ratio of plasma frequency $f_{\rm p}\simeq 9\times10^{-3}(n_{\rm c})^{1/2}$ MHz) and gyrofrequency $f_{\rm c}\simeq2.86\times B$ MHz) where $B$ is the magnetic field strength in G. In the quasi-linear diffusion process, the existence of background electromagnetic wave energy $T_{\rm w}$ is taken into account."
+ As discussed by ? in the case of solar millisecond spikes. the initial brightness temperature could be as low as the level of thermal bremsstrahlung in the range of 10°~105 K. However. because of gyroresonance or eyrosynchrotron radiation. an enhanced photon level could exist in a flaring loop which would affect the ECM process significantly. yielding a wave turbulence at the level of ~10 Κ. We apply Τι in the range of 10°—10!° K in our model.," As discussed by \cite{Aschwanden90a} in the case of solar millisecond spikes, the initial brightness temperature could be as low as the level of thermal bremsstrahlung in the range of $10^{6}\sim10^{8}$ K. However, because of gyroresonance or gyrosynchrotron radiation, an enhanced photon level could exist in a flaring loop which would affect the ECM process significantly, yielding a wave turbulence at the level of $\sim10^{15}$ K. We apply $T_{\rm w}$ in the range of $10^{\rm 6}-10^{16}$ K in our model."
+ The free energy of the plasma with a loss-cone distribution can be converted into an equivalent electromagnetic energy., The free energy of the plasma with a loss-cone distribution can be converted into an equivalent electromagnetic energy.
+ The energy conversion factor ες can be defined as the ratio of the change of kinetic energy between the initial state and the final state of the plasma and the initial kinetic energy., The energy conversion factor $\varepsilon_{\rm c}$ can be defined as the ratio of the change of kinetic energy between the initial state and the final state of the plasma and the initial kinetic energy.
+ In the present paper. the amount of converted energy is about0.5%.. ie. the same as in ?..," In the present paper, the amount of converted energy is about, i.e. the same as in \citet{Aschwanden90a}."
+ ? investigated how the parameters influence the growth rate and brightness temperature generated by the ECM instability and gives the general formula for the key parameters as follows: We plot in Figure 2. the effect of various parameters on the growth rate and brightness temperature., \citet{Aschwanden90a} investigated how the parameters influence the growth rate and brightness temperature generated by the ECM instability and gives the general formula for the key parameters as follows: We plot in Figure \ref{fig_grbt} the effect of various parameters on the growth rate and brightness temperature.
+ Many parameters. not only those describing the nature of the plasma but also those associated with the properties of the UCD. can affect the observable flux density significantly.," Many parameters, not only those describing the nature of the plasma but also those associated with the properties of the UCD, can affect the observable flux density significantly."
+ In this section. we will discuss these parameters and compare our simulations with observations.," In this section, we will discuss these parameters and compare our simulations with observations."
+" In order to study the influence of the various physical quantities we adopt ""standard"" values for each parameter. listed in Table 1.."," In order to study the influence of the various physical quantities we adopt `standard' values for each parameter, listed in Table \ref{tab_parameters}."
+ Unless otherwise noted. the simulations are always for magneto-ionic X-mode c=-1. harmonie number s =1.," Unless otherwise noted, the simulations are always for magneto-ionic X-mode $\sigma=-1$, harmonic number $s=$ 1."
+ In order to see the effect of rotation of the UCD. we first start à set of simulations by fixing the initial plasma and loss-cone parameters as listed in Table 1..," In order to see the effect of rotation of the UCD, we first start a set of simulations by fixing the initial plasma and loss-cone parameters as listed in Table \ref{tab_parameters}."
+ In addition. as suggested by ?.. we adopt the radius (Rucp =0.1 Res) as a constant in ouPos simulations.," In addition, as suggested by \citet{Chabrier00}, we adopt the radius $R_{\rm UCD}$ =0.1 $R_{\odot}$ ) as a constant in our simulations."
+ We initially assume that the radio-emitting regio is close to the equator (latitude €= 30°), We initially assume that the radio-emitting region is close to the equator (latitude $\theta=30^{\circ}$ ).
+ Given the above parameters. the size (A) of the emitting region and the tube size of a small radio area. we show in Fig.," Given the above parameters, the size $A$ ) of the emitting region and the tube size of a small radio area, we show in Fig."
+ 3 the influence of rotation., \ref{fig_tucd} the influence of rotation.
+ From the figure. we can see that the radio light curve ts broadened with increasing rotation period (Tucp) while the intensity does not change at all.," From the figure, we can see that the radio light curve is broadened with increasing rotation period $T_{\rm UCD}$ ) while the intensity does not change at all."
+" This is because the intensity of the radio emission is only related tc the behavior of plasmaand thetotal size of the emitting region,", This is because the intensity of the radio emission is only related to the behavior of plasmaand thetotal size of the emitting region.
+ Figure 4. shows the effect of the total size (A) of the radio- region (left panel) and the radius (rune) of the flux, Figure \ref{fig_rtubev} shows the effect of the total size $A$ ) of the radio-emitting region (left panel) and the radius $r_{\rm tube}$ ) of the flux
+are capable of suspending dust clouds in their photospheres at cooler temperatures than field brown cdwarls.,are capable of suspending dust clouds in their photospheres at cooler temperatures than field brown dwarfs.
+ The authors (hank the referee for his/her review. as well as Kevin Flaherty. Tom Greene. Mark Marley. Nevin Zahlinle. Kees Dullemond ancl James Raclomski for useful discussions.," The authors thank the referee for his/her review, as well as Kevin Flaherty, Tom Greene, Mark Marley, Kevin Zahnle, Kees Dullemond and James Radomski for useful discussions."
+ We also thank Amanda Morrow and Nevin Luhiman for providing us the IRS spectrum of 2NLASS 1207 and Adam Durrows for supplving the thick-cloud model grid., We also thank Amanda Morrow and Kevin Luhman for providing us the IRS spectrum of 2MASS 1207 and Adam Burrows for supplying the thick-cloud model grid.
+ AJS acknowledges ihe NASA Graduate Student Research. Program (GSRP) for its generous support., AJS acknowledges the NASA Graduate Student Research Program (GSRP) for its generous support.
+ LMC is supported by an NSF Career award and the NASA Origins of Solar Systems Program., LMC is supported by an NSF Career award and the NASA Origins of Solar Systems Program.
+ LSz acknowledges support from the Spitzer Data Analyzes erant. 1348621. the Lunearian OTA erant K7GOS16 and the Student Union of Science and Informatics of the University of Szeged.," LSz acknowledges support from the Spitzer Data Analyzes grant 1348621, the Hungarian OTKA grant K76816 and the Student Union of Science and Informatics of the University of Szeged."
+"empirical ZAMS is shifted along the reddening vector (indicated by the arrow in the upper-left corner) to fit the bulk of the stars in each cluster, using the green color, whenever this is possible.","empirical ZAMS is shifted along the reddening vector (indicated by the arrow in the upper-left corner) to fit the bulk of the stars in each cluster, using the green color, whenever this is possible."
+" As for the fit, this is done by considering only the stars with spectral type earlier than AO, since for later spectral type stars there is not a unique fitting solution."," As for the fit, this is done by considering only the stars with spectral type earlier than A0, since for later spectral type stars there is not a unique fitting solution."
+" To facilitate viewing, we indicated two spectral types along the ZAMS and the paths, parallel to the reddening vector, along which stars are displaced by reddening."," To facilitate viewing, we indicated two spectral types along the ZAMS and the paths, parallel to the reddening vector, along which stars are displaced by reddening."
+" This procedure allows us to estimate the mean reddening for each cluster, that we report in Table 8 (third column)."," This procedure allows us to estimate the mean reddening for each cluster, that we report in Table 8 (third column)."
+ The associated uncertainties have been estimated visually and represent the range in reddening we can move back and forth the ZAMS keeping the fit acceptable., The associated uncertainties have been estimated visually and represent the range in reddening we can move back and forth the ZAMS keeping the fit acceptable.
+ This uncertainty basically comes from photometric errors and variable reddening across the cluster., This uncertainty basically comes from photometric errors and variable reddening across the cluster.
+" When the uncertainty is larger than photometric errors, which are much less than 0.03 mag in this magnitude range (see Fig 2), we are forced to conclude that differential reddening is present in the cluster."," When the uncertainty is larger than photometric errors, which are much less than 0.03 mag in this magnitude range (see Fig 2), we are forced to conclude that differential reddening is present in the cluster."
+" This is far from being un-expected, since these clusters are typically located at low latitude in the inner Galactic disk, inside or close to gas- and dust-rich spiral features."," This is far from being un-expected, since these clusters are typically located at low latitude in the inner Galactic disk, inside or close to gas- and dust-rich spiral features."
+ Distances are estimated in the CMDs in Figs., Distances are estimated in the CMDs in Figs.
+" 9 to 11, using the same ZAMS as in the previous Section, and adopting the reddening values already derived."," 9 to 11, using the same ZAMS as in the previous Section, and adopting the reddening values already derived."
+" In this process, therefore, distance is the only free parameter, having fixed metallicity and reddening."," In this process, therefore, distance is the only free parameter, having fixed metallicity and reddening."
+" In the same way as for reddening, the uncertainty in distance is estimated vertically shifting the ZAMS (green line) for the range of distance moduli which provides an acceptable fit."," In the same way as for reddening, the uncertainty in distance is estimated vertically shifting the ZAMS (green line) for the range of distance moduli which provides an acceptable fit."
+ In this process we pay attention that the reddening remains within the range of values independently estimated in the TCD., In this process we pay attention that the reddening remains within the range of values independently estimated in the TCD.
+" The final step is to estimate the age, and for this we make use of isochrones from the Padova suite of models (Marigo et al."," The final step is to estimate the age, and for this we make use of isochrones from the Padova suite of models (Marigo et al."
+" 2008,"," 2008,"
+to the lens. and the sign of the dellection angle is chosen to match the sign of 8.,"to the lens, and the sign of the deflection angle is chosen to match the sign of $\theta$."
+" Ehe source and image positions and the dellection angle are linked. via the lens equation. .. where D, and D, are angular diameter distances from the observer to the source and from the lens to the source. respectively."," The source and image positions and the deflection angle are linked via the lens equation, = - , where $D_{\rm s}$ and $D_{\rm ls}$ are angular diameter distances from the observer to the source and from the lens to the source, respectively."
+ ‘To develop a general understanding of how the angular beam separation depends on the lensing geometry and the physical properties of the lens. we begin with a simple power law mass profile.," To develop a general understanding of how the angular beam separation depends on the lensing geometry and the physical properties of the lens, we begin with a simple power law mass profile."
+ In length units we write AZ(/2)=eb with il some constant. so in angular units we have M(8)=A(Dj8)*.," In length units we write $M(R) = A\,R^\gamma$ with $A$ some constant, so in angular units we have $M(\theta) = A(D_{\rm l}\theta)^\gamma$."
+ The cases =0 and 5=I correspond to the familiar cases of a point mass lens (PAL) and a singular isothermal sphere (SIS). respectively.," The cases $\gamma=0$ and $\gamma=1$ correspond to the familiar cases of a point mass lens (PM) and a singular isothermal sphere (SIS), respectively."
+ “Phe lens equation takes the form5... where we use the minus sign when 670 and the plus sign when 6«0. and the angular Einstein radius is .," The lens equation takes the form, where we use the minus sign when $\theta>0$ and the plus sign when $\theta<0$, and the angular Einstein radius is _E =."
+ Note that if we consider the mass within some fixed physical radius we have Afxο and hence GpκAtt , Note that if we consider the mass within some fixed physical radius we have $M \propto A$ and hence $\theta_E \propto M^{1/(2-\gamma)}$.
+bor 0x>«l the lens equation (5)) formally has two solutions for all source positions. although when 3 gets large the image on the opposite side of the lens is faint.," For $0\le\gamma<1$ the lens equation \ref{eq:powlens}) ) formally has two solutions for all source positions, although when $\beta$ gets large the image on the opposite side of the lens is faint."
+ For ~=1 the lens equation has one or two solutions depending on the position of the source. while for 1«52 it has one or three solutions.," For $\gamma=1$ the lens equation has one or two solutions depending on the position of the source, while for $1<\gamma<2$ it has one or three solutions."
+ For each solution. the corresponding dellection angle is given hy ," For each solution, the corresponding deflection angle is given by | = ."
+As we study how the angular beam separation depends on the lens redshift ο. we want to keep thep/rysieal properties of the lens fixed. which is why we elected to write the mass profile as AL=A where vA is a constant.," As we study how the angular beam separation depends on the lens redshift $z_{\rm l}$, we want to keep the properties of the lens fixed, which is why we elected to write the mass profile as $M = A\,R^\gamma$ where $A$ is a constant."
+ To facilitate the comparison of models with dillerent. power law slopes 5. we choose the value of el such that the dilferent masses all have the same Einstein radius when the lens is halfway between the observer and source.," To facilitate the comparison of models with different power law slopes $\gamma$, we choose the value of $A$ such that the different masses all have the same Einstein radius when the lens is halfway between the observer and source."
+ shows the angular beam separation as a function of lens redshift or à point mass lens (5=0). considering two values of the source τούς ancl different values o£ the angle 3 of the source with respect to the optical axis.," shows the angular beam separation as a function of lens redshift for a point mass lens $\gamma=0$ ), considering two values of the source redshift and different values of the angle $\beta$ of the source with respect to the optical axis."
+ AX strikinge result is the steep increase of & with τι., A striking result is the steep increase of $\delta$ with $z_l$.
+ This and other scalings can be understood. as follows., This and other scalings can be understood as follows.
+" The two images are located atthe angular positions 8,»=(3+|32)/2. and the corresponding dellection angles are (D.EDj)07.8,ο."," The two images are located atthe angular positions $\theta_{1,2}=(\beta \pm \sqrt{4\theta_E^2+\beta^2})/2$ , and the corresponding deflection angles are $\hat{\alpha}_{1,2}=(D_{\rm s}/D_{\rm ls})\;\theta_E^2/\theta_{1,2}$."
+ Equation (2)) then vields, Equation \ref{eq:delta}) ) then yields.
+ This equation elucidates the trends. apparent in the ligure., This equation elucidates the trends apparent in the figure.
+. First.. since. 0577I'MxDU1/2 we see that op- formally. as the lens approaches the source (24ας ancl hence Dy+ 0).," First, since $\theta_E^{\rm PM}\propto D_{\rm ls}^{1/2}$ we see that $\delta_{\rm PM}$ formally as the lens approaches the source $z_{\rm l} \to z_{\rm s}$ and hence $D_{\rm ls} \to 0$ )."
+ This divergence occurs for all values of the source angle 3., This divergence occurs for all values of the source angle $\beta$.
+ Second. it is clear that δν20 as the lens approaches the observer (τιO and 11.D.).," Second, it is clear that $\delta_{\rm PM} \to 0$ as the lens approaches the observer $z_{\rm l} \to 0$ and $D_{\rm ls} \to D_{\rm s}$ )."
+ Third. when the source and lens redshifts and the lens mass are all lixed. orsi; increases with 2S.," Third, when the source and lens redshifts and the lens mass are all fixed, $\delta_{\rm PM}$ increases with $\beta$."
+ Fourth. when the source is well aligned with the lens (7κθε}. the angular beam separation scales with the lens mass as pstxA7.," Fourth, when the source is well aligned with the lens $\beta \ll \theta_E$ ), the angular beam separation scales with the lens mass as $\delta_{\rm PM} \propto M^{1/2}$."
+" Finally. in the opposite limit in which 43 Ae. dpa, becomes independent of the mass of the lens (although this particular case is perhaps less relevant than the others because when 26E the counter-image is [aint)."," Finally, in the opposite limit in which $\beta \gg \theta_E$ , $\delta_{\rm PM}$ becomes independent of the mass of the lens (although this particular case is perhaps less relevant than the others because when $\beta \gg \theta_E$ the counter-image is faint)."
+ shows the results for other power law profiles. including one that is extended but steeper than isothermal 1.5). the isothermal profile (4= 1). and one shallower prolile (= 1.5).," shows the results for other power law profiles, including one that is extended but steeper than isothermal$\gamma=0.5$ ), the isothermal profile $\gamma=1$ ), and one shallower profile $\gamma=1.5$ )."
+ For all cases we fix j—10! deg. which is small enough that we are always looking at situations with multiple images.," For all cases we fix $\beta=10^{-4}$ deg, which is small enough that we are always looking at situations with multiple images."
+ As we shall see. for shallow profiles the angular beam separation is not very sensitive to the choice of 3.," As we shall see, for shallow profiles the angular beam separation is not very sensitive to the choice of $\beta$."
+ Qualitatively. steep profiles (5< 1) behave in a similar wav as the point mass case. with ? increasing monotonically with zj and cdiverging as the lens nears the source.," Qualitatively, steep profiles $\gamma<1$ ) behave in a similar way as the point mass case, with $\delta$ increasing monotonically with $z_{\rm l}$ and diverging as the lens nears the source."
+ In these cases we find that 9 increases with the source position :? (the case >=0 is shown in2: the case 5=0.5 is not shown).The isothermal profile (7.= 1) represents a transition (ης 9 inercasing monotonically with the lens redshift and reaching a finite value as τι ze ," In these cases we find that $\delta$ increases with the source position $\beta$ (the case $\gamma=0$ is shown in; the case $\gamma=0.5$ is not shown).The isothermal profile $\gamma=1$ ) represents a transition case, with $\delta$ increasing monotonically with the lens redshift and reaching a finite value as $z_{\rm l} \to z_{\rm s}$ ."
+This is another case we can understand analytically., This is another case we can understand analytically.
+" Por 3<65 there are two images atpositions £6,»=Op c."," For $\beta < \theta_E$ there are two images atpositions $\theta_{1,2}=\theta_E\pm\beta$ ."
+ Combining these with the deflection angles àο=CD.1.18 vields the angular beam separation An SIS has Gpx Di. so when zz; we," Combining these with the deflection angles $\hat{\alpha}_{1,2}=(D_{\rm s}/D_{\rm ls})\,\theta_E$ yields the angular beam separation -1 .An SIS has $\theta_E\propto D_{\rm ls}$ , so when $z_{\rm l}\rightarrow z_{\rm s}$ we"
+reproduce their SEDs.,reproduce their SEDs.
+ Ou the other haud. up to 6 ERGs have properties that are formally consistent with the strict definition of beige dustless. old and passively evolved spheroidals at +>1 amc With tgonuarionD2.," On the other hand, up to 6 ERGs have properties that are formally consistent with the strict definition of being dustless, old and passively evolved spheroidals at $z \gtrsim 1$ and with $z_{{\rm formation}}>2$."
+ Although the observed sample is rather simall aud. still complete. it is tempting to speculate on how the above results may have Huplications ou the problem of the formation of elliptical ealaxies.," Although the observed sample is rather small and still incomplete, it is tempting to speculate on how the above results may have implications on the problem of the formation of elliptical galaxies."
+ Iu fact. the existence and the abundance of high-: ellipticalsis one of the most controversial issues of galaxy evolution.," In fact, the existence and the abundance of $z$ ellipticals is one of the most controversial issues of galaxy evolution."
+ Some works claimed that the nuuber of galaxies with the red colors expected for passively evolved sphleroidals is lower compared to the predictions of passive lIunimositv evolution(e.g. Wantfinaun .Clhnarlot White 1996: Zepf 1997: Franceschini et al.," Some works claimed that the number of galaxies with the red colors expected for $z$ passively evolved spheroidals is lower compared to the predictions of passive luminosity evolution (e.g. Kauffmann, Charlot White 1996; Zepf 1997; Franceschini et al."
+ 1998: Darger et al., 1998; Barger et al.
+ 1999)., 1999).
+ However. other works did not confirm the existence of such a deficit up to :x2 (eg Totani Yoshii 1997: Benitez et al.," However, other works did not confirm the existence of such a deficit up to $z\approx2$ (e.g. Totani Yoshii 1997; Benitez et al."
+ 1999: Droadhurst Bowens 1999: Schade et al., 1999; Broadhurst Bowens 1999; Schade et al.
+ 1999)., 1999).
+ This picture is complicated by the possibility hat such spleroidals are quer because of a low level of residual star formation. thus escaping the selection criteria sed on red colors (e.g. Jimenez et al.," This picture is complicated by the possibility that such spheroidals are bluer because of a low level of residual star formation, thus escaping the selection criteria based on red colors (e.g. Jimenez et al."
+ 1999)., 1999).
+ Thus. it is clear that a reliable comparison between he observed abundance of hieh-: ellipticals aud the oue expected from the various galaxy formation models eau be formed ouly if the fraction of hieh-: cllipticals in ERG saluples is firmly established.," Thus, it is clear that a reliable comparison between the observed abundance of $z$ ellipticals and the one expected from the various galaxy formation models can be performed only if the fraction of $z$ ellipticals in ERG samples is firmly established."
+ If taken at face value. our analvsis sugeestsOO that a substantial fraction of ERGs are consistent withbeiug hüeh-: spheroidals.," If taken at face value, our analysis suggests that a substantial fraction of ERGs are consistent with being $z$ spheroidals."
+ Should spectroscopy of complete aud lareer samples σολσα such result. this would allow a reliable comparison ofthe observed aud predicted nuubers of splicroidals at +> 1.," Should spectroscopy of complete and larger samples confirm such result, this would allow a reliable comparison ofthe observed and predicted numbers of spheroidals at $z \gtrsim 1$ ."
+any particular galaxy. and are often larger in size.,"any particular galaxy, and are often larger in size."
+ Radio haloes are usually projected towards the cluster centre. while relies are scen towards the periphery (e.g.Giovannini&Peretti2004).," Radio haloes are usually projected towards the cluster centre, while relics are seen towards the periphery \citep[e.g.][]{Giovannini04}."
+. Phere are 730 radio haloes in nearby (z«20.4) clusters of galaxies (e.g.Cio- 2009).. and there are ~30 clusters of ealaxics with at least one radio relic (Ciovannini& 2004).," There are $\sim$ 30 radio haloes in nearby $<$ 0.4) clusters of galaxies \citep[e.g.][]{Giovannini09}, and there are $\sim$ 30 clusters of galaxies with at least one radio relic \citep{Giovannini04}."
+. While mocels for haloes range from re- of particles by turbulence to production of relativistic electrons by haclronic collisons. relics are believed to arise due to cluster mergers and/or matter accretion (Sarazin1999:Ryuetal.2003:Plrommerllunstead&Corbett2008:BrownIuudnick 2009).," While models for haloes range from re-acceleration of particles by turbulence to production of relativistic electrons by hadronic collisons, relics are believed to arise due to cluster mergers and/or matter accretion \citep{Sarazin99, Ryu03, Pfrommer06, Giacintucci08, mjh08,Brown09}."
+. Recent work suggests that halos are found. in massive. unrelaxed clusters. with the radio and X-ray uminositv being strongly. correlated. consistent. with he re-acceleration scenario (Brunettietal.2007:Ven-urietal.2008:Cassano 2009).," Recent work suggests that halos are found in massive, unrelaxed clusters, with the radio and X-ray luminosity being strongly correlated, consistent with the re-acceleration scenario \citep{Brunetti07, Venturi08, Cassano09}."
+. Llowever. the presen studies have been based on X-ray. selected: clusters of galaxies. ancl possible biases arising from it should be »»rne in mind.," However, the present studies have been based on X-ray selected clusters of galaxies, and possible biases arising from it should be borne in mind."
+ For example. the limited sensitivity of he radio observations would make it easier to cdetec ios in only the more X-rav luminous clusters. of ealaxies.," For example, the limited sensitivity of the radio observations would make it easier to detect halos in only the more X-ray luminous clusters of galaxies."
+ Raclio relics on the other hand. are. believes Oo arise due to mergers accompanied by shocks and/or matter accretion (e.g. Bagchi et al., Radio relics on the other hand are believed to arise due to mergers accompanied by shocks and/or matter accretion (e.g. Bagchi et al.
+ 2006. ane references therein).," 2006, and references therein)."
+ These shocks are capable of accelerating particles to high energies. giving rise to the observed. svuchrotron radio emission.," These shocks are capable of accelerating particles to high energies, giving rise to the observed synchrotron radio emission."
+ Harris.Ixa-pahi&Ekers(1980). ancl Tribble(1993) were amongs the carly ones to suggest and. explore the possibility of acceleration. of particles due to shock [ronts on a large scale caused. by mergers.," \citet{Harris80} and \citet{Tribble93}
+ were amongst the early ones to suggest and explore the possibility of acceleration of particles due to shock fronts on a large scale caused by mergers."
+ “These ideas were expanded upon by EnBlinetal.(1998)... Itoettiger.jurns&Stone(1900) EnBlin&|Ciopal-Ixrishna and Ricker&Sarazin(2001).. producing more sophisticated. models.," These ideas were expanded upon by \citet{Ensslin98}, \citet{Roettiger99}
+ \citet{Ensslin01} and \citet{Ricker01}, producing more sophisticated models."
+ The ATLAS radio image at 1.4. Cllz (Fig. 9)), The ATLAS radio image at 1.4 GHz (Fig. \ref{disturbed}) )
+ shows two more extended sources within ~20 arcmin of the WAT source. one of which (91051) appears to be a radio relic (Middelbergetal.2008)... while the other (1110) is an FRI radio galaxy.," shows two more extended sources within $\sim$ 20 arcmin of the WAT source, one of which (S1081) appears to be a radio relic \citep{Middelberg08}, while the other (S1110) is an FRI radio galaxy."
+ Superpositions of the radio image of the relic on an optical DSS red image as well as an infrared (im image are shown in Fig. 10.., Superpositions of the radio image of the relic on an optical DSS red image as well as an infrared $-\mu$ m image are shown in Fig. \ref{relic}.
+ While no optical object is visible within the radio contours. there is an infrared object towards the central region of the source.," While no optical object is visible within the radio contours, there is an infrared object towards the central region of the source."
+ This object has been classified. as an She galaxy (optical template type 5) using a total of 6 photometric bands by. Rowan-Robinson et al. (, This object has been classified as an Sbc galaxy (optical template type 5) using a total of 6 photometric bands by Rowan-Robinson et al. (
+2008).,2008).
+ Its photometric redshift’ has been estimated to be 1.18., Its photometric redshift has been estimated to be 1.18.
+ Given the properties of the object. it is likely to be unrelated.," Given the properties of the object, it is likely to be unrelated."
+ The radio properties of the WAT and these two sources are summarisec in Table ει., The radio properties of the WAT and these two sources are summarised in Table \ref{disturbedtable}.
+ At a redshift of 0.22. the relic woule have a physical size of ~274 kpe and a luminosity. of 107* W which would make it similar to the relies found. in the periphery of clusters of ealaxies in the local Universe (Ferrarictal.2008).," At a redshift of 0.22, the relic would have a physical size of $\sim$ 274 kpc and a luminosity of $\times$ $^{23}$ W $^{-1}$, which would make it similar to the relics found in the periphery of clusters of galaxies in the local Universe \citep{Ferrari08}."
+. The relie is at a projected. distance of ~2 Alpe from the ¢D galaxy., The relic is at a projected distance of $\sim$ 2 Mpc from the cD galaxy.
+ Typically relies have been observec at distances of about a Alpe from the cluster centre. although some systems are known to have relies up," Typically relics have been observed at distances of about a Mpc from the cluster centre, although some systems are known to have relics up"
+has a poloidal cuerey fraction Γρ=0.052.,has a poloidal energy fraction $E_{\rm p}/E=0.052$.
+" The maegneticto-thermal euergv ratio £/Uzz1/100. which should mean that for stability we require E,/E20.0L (Draithwaite2009)."," The magnetic-to-thermal energy ratio $E/U\approx1/400$, which should mean that for stability we require $E_{\rm p}/E \gtrsim 0.04$ \citep{Braithwaite:2009}."
+. We see then that this value of EU. is near the upper liuüt for stability im other words. we are near the boundary of validity of the weak-field approximation used in Section 2..," We see then that this value of $E/U$ is near the upper limit for stability – in other words, we are near the boundary of validity of the weak-field approximation used in Section \ref{analytic}."
+ Using seni-analvtie methods an axisviunietrie non force-free maeuctostatic equilibrium which could exist in any non-convective stellar region: the radiative core of solar-type stars. the external envelope of inassive stars. and compact objects.," Using semi-analytic methods an axisymmetric non force-free magnetostatic equilibrium which could exist in any non-convective stellar region: the radiative core of solar-type stars, the external envelope of massive stars, and compact objects."
+ Using umuerical siuulations. we then find stability of the mixed configuration under all imaginable perturbations.," Using numerical simulations, we then find stability of the mixed configuration under all imaginable perturbations."
+ This is the first time the stability of an analvtically-derived stellar magnetic equilibrimm has (enconfrined nuuercallv., This is the first time the stability of an analytically-derived stellar magnetic equilibrium has beenconfirmed numerically.
+The time duration per event is small.,The time duration per event is small.
+ Note that. if deviations from baseline of 2% can be detected. the duration of an event can be 3 times longer-ie.. the time taken to cross 3 Einstein diameters.," Note that, if deviations from baseline of $2\%$ can be detected, the duration of an event can be $3$ times longer–i.e., the time taken to cross $3$ Einstein diameters."
+ Given the fact that there ts a distribution of velocities. and also a distribution of orientations. producing an even broader distribution of transverse velocities. we expect both shorter and longer events. Because the Einstein ring is small. we are in the low-density regime. and events should primarily be single-source events.," Given the fact that there is a distribution of velocities, and also a distribution of orientations, producing an even broader distribution of transverse velocities, we expect both shorter and longer events, Because the Einstein ring is small, we are in the low-density regime, and events should primarily be single-source events."
+ Nevertheless. as (7) indicates. sequences of events are expected because the angular speed is large.," Nevertheless, as (7) indicates, sequences of events are expected because the angular speed is large."
+ The total rate per square degree is also high., The total rate per square degree is also high.
+" Within 100 pe. there would be ~[00 masses of 10°"" em per square degree."," Within $100$ pc, there would be $\sim 100$ masses of $10^{30}$ gm per square degree."
+ The duty cycle for events associated with these nearby lenses could be high enough to allow monitoring specifically designed to discover short-duration events to either discover or place limits on planet-mass. free-streaming dark matter.," The duty cycle for events associated with these nearby lenses could be high enough to allow monitoring specifically designed to discover short-duration events to either discover or place limits on planet-mass, free-streaming dark matter."
+ If short-duration events are observed. and if the monitoring occurs regularly over the period of several years. several tracks of events should be observed.," If short-duration events are observed, and if the monitoring occurs regularly over the period of several years, several tracks of events should be observed."
+ The number of tracks actually associated with lenses effect can be estimated through statistical simulations., The number of tracks actually associated with lenses effect can be estimated through statistical simulations.
+ Individual cases in which the tracks are reliably established as being caused by lensing will be important. because the distance to the lens and the lens may may be measured in such cases.," Individual cases in which the tracks are reliably established as being caused by lensing will be important, because the distance to the lens and the lens may may be measured in such cases."
+ If no such events are discovered. limits on disk dark matter in the form of Jupiter-mass objects can be derived.," If no such events are discovered, limits on disk dark matter in the form of Jupiter-mass objects can be derived."
+ Such limits on Halo dark matter are already strong. based on the EROS and MACHO data Alcock et 1998.," Such limits on Halo dark matter are already strong, based on the EROS and MACHO data Alcock et 1998."
+ Mesolensing observations. however. can place limits on disk dark matter.," Mesolensing observations, however, can place limits on disk dark matter."
+ In addition. the limits can be placed on lower mass values.," In addition, the limits can be placed on lower mass values."
+ For lower-mass objects. the Einstein angle is smaller.," For lower-mass objects, the Einstein angle is smaller."
+ This does not cause problems with finite source size. but it does decrease the time of events. making it more difficult to discover them.," This does not cause problems with finite source size, but it does decrease the time of events, making it more difficult to discover them."
+ If. however. all of the dark matter is in the form of such low-mass objects. there will be more of them.," If, however, all of the dark matter is in the form of such low-mass objects, there will be more of them."
+ This means that à larger number will be closer. so that the distribution of values of 0p will still include some larger values.," This means that a larger number will be closer, so that the distribution of values of $\theta_E$ will still include some larger values."
+ It also means that there will be a larger number of lenses with velocities directed along more radial paths. producing fewer events per year. but with each event having a longer duration than if the motion were perpendicular to our line of sight.," It also means that there will be a larger number of lenses with velocities directed along more radial paths, producing fewer events per year, but with each event having a longer duration than if the motion were perpendicular to our line of sight."
+ We don't know whether the outer regions of our solar system. harbor planet-mass objects.," We don't know whether the outer regions of our solar system, harbor planet-mass objects."
+ At distances of ~1000 AU a planet would have to be far more massive than Jupiter. in order for its dynamical influence on the known outer planets to be discernible (Hogg et 1991).," At distances of $\sim 1000$ AU a planet would have to be far more massive than Jupiter, in order for its dynamical influence on the known outer planets to be discernible (Hogg et 1991)."
+ Mesolensing provides an independent way to derive limits., Mesolensing provides an independent way to derive limits.
+ It is possible. if a planet-mass object is close enough to us. for its angular size to be larger than its Einstein ring.," It is possible, if a planet-mass object is close enough to us, for its angular size to be larger than its Einstein ring."
+" The requirement that it fit inside its Einstein ring ts: where M, and AR, are the mass and radius of the planet. respectively."," The requirement that it fit inside its Einstein ring is: where $M_p$ and $R_p$ are the mass and radius of the planet, respectively."
+" harbor planets comparable in mass For a Jupiter-mass planet located 2000 AU from us. the size of the Einstein ring is roughly 0.02"". For less massive planets. however. the Einstein angle can be in the milliaresecond range (e.g.. the Earth at 2000 AU) or even smaller."," harbor planets comparable in mass For a Jupiter-mass planet located $2000$ AU from us, the size of the Einstein ring is roughly $0.02''.$ For less massive planets, however, the Einstein angle can be in the milliarcsecond range (e.g., the Earth at $2000$ AU) or even smaller."
+ The event rate 1s high nevertheless. because the angular speed on the sky ts high.," The event rate is high nevertheless, because the angular speed on the sky is high."
+" A mass 2000 AU from us traverses an angle of approximately 200"" every 6 months.", A mass $2000$ AU from us traverses an angle of approximately $200''$ every $6$ months.
+" If the Einstein angle is 0.02"". the lens will move through ~5000 Einstein diameters in 6 months. and the Einstein crossing time will be approximately an hour."," If the Einstein angle is $0.02'',$ the lens will move through $\sim 5000$ Einstein diameters in $6$ months, and the Einstein crossing time will be approximately an hour."
+" The area we will perceive the lens to cover across the source field per year is large: 4L1"" every 6 months."," The area we will perceive the lens to cover across the source field per year is large: $4\, \Box''$ every $6$ months."
+ The background field could be out of the ecliptic. because the distribution of outer solar system masses may be close to spherical.," The background field could be out of the ecliptic, because the distribution of outer solar system masses may be close to spherical."
+ If such a planet. or planet exists. lensing by it is therefore most likely to be detected by wide-field monitoring programs like Pan-STARRS and LSST.," If such a planet, or planet exists, lensing by it is therefore most likely to be detected by wide-field monitoring programs like Pan-STARRS and LSST."
+ The smaller values of fy possible for these programs will also merease the rate at which any such planets cause events., The smaller values of $f_T$ possible for these programs will also increase the rate at which any such planets cause events.
+ Although the frequency of detected events and their durations will depend on the direction. there are important common elements across directions.," Although the frequency of detected events and their durations will depend on the direction, there are important common elements across directions."
+ First. events will be of short duration. typically hours or less.," First, events will be of short duration, typically hours or less."
+ Second. individual lenses should each give rise to sequences of events.," Second, individual lenses should each give rise to sequences of events."
+" These sequences will trace the motion of the Earth. exhibiting a ""forward"" and ""backward"" motion every year."," These sequences will trace the motion of the Earth, exhibiting a “forward"" and “backward"" motion every year."
+ This back and forth swing will repeat on a yearly basis. with slight changes due to the relatively slow motions of the lens and background stars.," This back and forth swing will repeat on a yearly basis, with slight changes due to the relatively slow motions of the lens and background stars."
+ It is Important to note that astrometric effects can be significant as well. and may be studied with. e.g.. Gata (see Gaudi and Bloom 2005).," It is important to note that astrometric effects can be significant as well, and may be studied with, e.g., Gaia (see Gaudi and Bloom 2005)."
+" Discussions of ""dark matter"" typically focus on matter whose nature is yet to be understood."," Discussions of “dark matter"" typically focus on matter whose nature is yet to be understood."
+ The local neighborhood is. however. filled with dark objects whose nature we think we understand (BHs. NSs. WDs. and low-mass dwarfs). but of which we have relatively few nearby examples.," The local neighborhood is, however, filled with dark objects whose nature we think we understand (BHs, NSs, WDs, and low-mass dwarfs), but of which we have relatively few nearby examples."
+ Mesolensing provides a way to conduct a census of such objects. providing mass estimates and distances for many. and opening the door to detailed study of some individual systems.," Mesolensing provides a way to conduct a census of such objects, providing mass estimates and distances for many, and opening the door to detailed study of some individual systems."
+ Mesolensing observations are guaranteed to identify BHs. NSs. WDs. and low-mass stars within ~1 kpe of Earth.," Mesolensing observations are guaranteed to identify BHs, NSs, WDs, and low-mass stars within $\sim 1$ kpc of Earth."
+ The calculations of $3 indicate that events due to nearby lenses (1) should be part of virtually every microlensing data set collected so far. and that (2) wide-field monitoring programs will regularly discover signals due to nearby lenses; this provides additional motivation for Pan-STARRS and LSST. (," The calculations of 3 indicate that events due to nearby lenses (1) should be part of virtually every microlensing data set collected so far, and that (2) wide-field monitoring programs will regularly discover signals due to nearby lenses; this provides additional motivation for Pan-STARRS and LSST. ("
+3) A combination of programs will regularly discover BHs within a few kpe. as well as NSs and other objects of lower-mass.,"3) A combination of programs will regularly discover BHs within a few kpc, as well as NSs and other objects of lower-mass."
+by the Leverhulme Trust. and. STI by PPARC: PAT is a PPARC Lecturer Fellow.,by the Leverhulme Trust and STK by PPARC; PAT is a PPARC Lecturer Fellow.
+The first generation of stars. Population OUT stars. are thought to form from almost metal-free eas with jietallicity of ~10.MZ... (see eig. Carr. Boud. Arnett 1981: Carr 1991).,"The first generation of stars, Population III stars, are thought to form from almost metal-free gas with metallicity of $\sim 10^{-10} Z_\odot$ (see e.g., Carr, Bond, Arnett 1984; Carr 1994)."
+ Such Population ΤΗ stars iav have produced heavy elements at redshifts 2= 10., Such Population III stars may have produced heavy elements at redshifts $z \gtrsim 10$ .
+ Recent observatious of quasar absorption spectra show that the metal abundance in intergalactic space ise10?Z. (eg. Cowie Songaila 1998).," Recent observations of quasar absorption spectra show that the metal abundance in intergalactic space is $\approx 10^{-3} Z_\odot$ (e.g., Cowie Songaila 1998)."
+ Also. it is widely accepted that a siguicant portion of old halo stars in our Galaxy lave metallicity lower than 10?Z. (Deers. Preston. Shectinan 1992).," Also, it is widely accepted that a significant portion of old halo stars in our Galaxy have metallicity lower than $10^{-3}Z_\odot$ (Beers, Preston, Shectman 1992)."
+ In addition. some blue conact dwarf galaxies are known to be extremely mictal-voor (m10?Z.. Kunth Sargent 1986: Pustiluik et al.," In addition, some blue compact dwarf galaxies are known to be extremely metal-poor $\approx 10^{-2} Z_\odot$, Kunth Sargent 1986; Pustilnik et al."
+ 2001)., 2001).
+ Thus. star formation in protogalaxies must lave xoceeded in very iietaledeficieut. environments.," Thus, star formation in protogalaxies must have proceeded in very metal-deficient environments."
+ When the metallicity is lower than ZX10?Z.. the cooling bv heavy elemieuts is less effective than cooling w olvdrogen aud helimu (c.e.. Yoshi Sabauo 1980: Boliirinecr Teusler 1989: Onimkai 2000: Nishi Tashiro 2000).," When the metallicity is lower than $Z \lesssim 10^{-2}Z_\odot$, the cooling by heavy elements is less effective than cooling by hydrogen and helium (e.g., Yoshii Sabano 1980; Böhhringer Hensler 1989; Omukai 2000; Nishi Tashiro 2000)."
+ For the formation of Population III stars from such uctal-deficicut eas. the cooling by primordial hydrogen uolecules (10) plavs an esseutial role (Alatsuda. Sato. Takeda 1969: Yonevama 1972: Wutchins 1976: Silk 1977: Yoshii Sabano 1980: Carlberg 1981: Lepp Shull 1981: Palla. Salpeter. Staller 1983: Yoshii Saio 1986: Shapiro Itaug 19857: Uchara et al.," For the formation of Population III stars from such metal-deficient gas, the cooling by primordial hydrogen molecules $_2$ ) plays an essential role (Matsuda, Sato, Takeda 1969; Yoneyama 1972; Hutchins 1976; Silk 1977; Yoshii Sabano 1980; Carlberg 1981; Lepp Shull 1984; Palla, Salpeter, Stahler 1983; Yoshii Saio 1986; Shapiro Kang 1987; Uehara et al."
+ 1996: Waiman. Thoul Loeb 1996: Nishi et al.," 1996; Haiman, Thoul Loeb 1996; Nishi et al."
+ 1998: Abel et al., 1998; Abel et al.
+ 1998: Dronuu. Coppi. Larson 1999: Abel. Bryan. Norman 2000: Nalkuuura Uinmenmra 1999a. 2001. hereafter Papers I and ID.," 1998; Bromm, Coppi, Larson 1999; Abel, Bryan, Norman 2000; Nakamura Umemura 1999a, 2001, hereafter Papers I and II)."
+ Beceut studies have revealed that star formation in priuordial eas is considerablv different four prescut-day star formation. duaplving that the stellar initial mass function (AIF) night be different from the Salpeter-like IMFE and the IME might be time-varvine in the course of galaxy evolution from the carly collapsing stages to the present day (see also Larson 1998: Zept Silk 1996: Chabrier 1999).," Recent studies have revealed that star formation in primordial gas is considerably different from present-day star formation, implying that the stellar initial mass function (IMF) might be different from the Salpeter-like IMF and the IMF might be time-varying in the course of galaxy evolution from the early collapsing stages to the present day (see also Larson 1998; Zepf Silk 1996; Chabrier 1999)."
+ Besides hydrogen iolecules. deuterated lydrogeu molecules (IID) cau be a significant coolant (e... Galli Palla 1998).," Besides hydrogen molecules, deuterated hydrogen molecules (HD) can be a significant coolant (e.g., Galli Palla 1998)."
+" Although. ΠΟ is less abundaut than I> (UID/IL,| ~10%10. 4). TED has a finite dipole moment and thus higher radiative transition probabilities than IL.."," Although HD is less abundant than $_2$ $_2$ ] $\sim 10^{-3}- 10^{-4}$ ), HD has a finite dipole moment and thus higher radiative transition probabilities than $_2$ ."
+ For example. the lowest rotational transitions of Is and WD have radiative transition probabilities of oy= 3 aud dy95«10P s f. respectively.," For example, the lowest rotational transitions of $_2$ and HD have radiative transition probabilities of $A_{20}=3\times 10^{-11}$ $^{-1}$ and $A_{10}=5\times 10^{-8}$ $^{-1}$ , respectively."
+ Also.the corresponding excitation cucreics of IT; and TDare," Also,the corresponding excitation energies of $_2$ and HDare"
+"From the inferred lobe cuthalpy £lebe the jet kinetic power cun be estimated as where #4.&72fes is the characteristic time required ο inflate a lobe with width /mae at the sound speed e;22),","From the inferred lobe enthalpy $E_\ditto{lobe}$, the jet kinetic power can be estimated as where $t_{_{\rm lobe}}\simeq l_{_{\rm lobe}}/c_s$ is the characteristic time required to inflate a lobe with width $l_{_{\rm lobe}}$ at the sound speed $c_s$."
+" At the powerful end. the lobes of ER ΤΠ galaxies dossess average kinetic inputs of —101aSeresi correspouding to the Thomson EddingtonL, limit for black voles with mass 10?Lot!ALF The fact that the upper lait for L, is close to the Eddington liuit for the largest black holes informs us hat. at least during periods of powerful activity. a value of Ay,c Lomay be appropriate for these sources. and that X,l may serve as a reasonable upper vou."," At the powerful end, the lobes of FR II galaxies possess average kinetic inputs of $L_{_{\rm J}}\simeq 10^{47}-10^{48}\,{\rm erg\,s^{-1}}$, corresponding to the Thomson Eddington limit for black holes with mass $10^9-10^{10}\,M_{\odot}$ The fact that the upper limit for $L_{_{\rm J}}$ is close to the Eddington limit for the largest black holes informs us that, at least during periods of powerful activity, a value of $\Lambda_{_{\rm Edd}}\simeq 
+1$ may be appropriate for these sources, and that $\Lambda_{_{\rm Edd}}\sim 1$ may serve as a reasonable upper bound."
+" For powerful FSRQs. thought to be the ou- counterpart of massive FR II galaxies. inter comparable values for £, from modcling the blazar spectral energy distribution."," For powerful FSRQs, thought to be the on-axis counterpart of massive FR II galaxies, infer comparable values for $L_\ditto{J}$ from modeling the blazar spectral energy distribution."
+ Tuformation on particle composition at the jet dissipation scale is best extracted from the spectra of FSROs. where both jet and disk cinission are present.," Information on particle composition at the jet dissipation scale is best extracted from the spectra of FSRQs, where both jet and disk emission are present."
+ The ecueral procedure is outlined by?., The general procedure is outlined by.
+. The jet is bathed by optical/UV photons from the BLR aud IR photons from the presumed obscuring torus., The jet is bathed by optical/UV photons from the BLR and IR photons from the presumed obscuring torus.
+" Owing to the large bulk Loreutz factor D~10 of the jet flow. ""coll electrons iu the jet have the capacity to Compton up-seatter these relatively soft photons (a process known as ""bulk-C'omptonization or first-order kinetic Sunvaev - Zeldovich effect) up to characteristic enorgles where hpc1006V is the typical seed plioton energy."," Owing to the large bulk Lorentz factor $\Gamma\sim10$ of the jet flow, “cold” electrons in the jet have the capacity to Compton up-scatter these relatively soft photons (a process known as “bulk-Comptonization” or first-order kinetic Sunyaev - Zel'dovich effect) up to characteristic energies where $h\nu_{_{\rm UV}}\simeq10\unit{eV}$ is the typical seed photon energy."
+" The expected bulk-Compton huninositv £,,. frou cold jet electrons may be written as a volume integral extending from the jet base ~ει to the dissipation scale hes below the region where electrons are shock-heated:240. where dE,/dt is the rate at which a cold electrou loses its orderly kinetic energv via the bulk-C'oniptonizatiou process. is the Thomson cross section πα”.OL.f(rRee)94 is the enerey deusity at radius R resulting from the fraction © of accretion luuünositv ος re-processed and isotropized by the DER."," The expected bulk-Compton luminosity $L_\ditto{BC}$ from cold jet electrons may be written as a volume integral extending from the jet base $\sim R_\ditto{G}/\epsilon_\ditto{rad}$ to the dissipation scale $R_\ditto{diss}$, i.e., below the region where electrons are shock-heated: where $\ud E_e/\ud t$ is the rate at which a cold electron loses its orderly kinetic energy via the bulk-Comptonization process, $\sigma_\ditto{T}$ is the Thomson cross section and $U_\ditto{BLR}=\xi\,L_{_{\rm acc}}/(4\pi R^2c)$ is the energy density at radius $R$ resulting from the fraction $\xi$ of accretion luminosity $L_\ditto{acc}$ re-processed and isotropized by the BLR."
+ Eq. (20)), Eq. \ref{e: bulk_Comp}) )
+" leads to au expression for the Thomson optical depth z,,, from the jet base up to the dissipation region: where we have durposedo that the expected bulk-Compton Iuuinositv £,,.. should not exceed the observed integrated power ως at soft N-rav energies Gvlcre the bulk-Conptonu emission should peak. see (19)))."," leads to an expression for the Thomson optical depth $\tau_\ditto{cold}$ from the jet base up to the dissipation region: where we have imposed that the expected bulk-Compton luminosity $L_\ditto{BC}$ should not exceed the observed integrated power $L_{_{\rm SX}}$ at soft X-ray energies (where the bulk-Compton emission should peak, see )."
+" Asstuning that the flux of cold electrons is conserved along the jet (1.c.. pair injection occurs primarily at the jet base). it follows that the Thomson depth 7. down to the dissipation scale satisfies 7,00S7,4. since the optical depth is greater at naller radii."," Assuming that the flux of cold electrons is conserved along the jet (i.e., pair injection occurs primarily at the jet base), it follows that the Thomson depth $\tau_\ditto{diss}$ down to the dissipation scale satisfies $\tau_{_{\rm diss}}\lesssim\tau_{_{\rm cold}}$, since the optical depth is greater at smaller radii."
+ Thus. from eqs. (12))," Thus, from eqs. \ref{eq:tau_diss}) )"
+ and (21)) the electron to proton ratio of the jet flow is constrained to be which informs us that electrons aud positrous. though donunant by uuniber. only advect with them au insignificant fraction of the jet power.," and \ref{eq:bc_lum}) ) the electron to proton ratio of the jet flow is constrained to be which informs us that electrons and positrons, though dominant by number, only advect with them an insignificant fraction of the jet power."
+ That is. the kinetic power of the jet is almost entirely carried by ious. as asstumed in refsecijetinodel.," That is, the kinetic power of the jet is almost entirely carried by ions, as assumed in \\ref{sec:jetmodel}."
+" Also. frou with ~10 it follows that τος20.2. aud the flow 3n then/n, dissipation reeion is optically thin."," Also, from with $n_e/n_p\sim10$ it follows that $\tau_\ditto{diss}\simeq0.2$, and the flow in the dissipation region is optically thin."
+" As in the case of estimating the jet composition. are the best sources for measuring the relative streneth ofthe magnetic field at the dissipation scale. parametrized by e,=UL/UT."," As in the case of estimating the jet composition, are the best sources for measuring the relative strength ofthe magnetic field at the dissipation scale, parametrized by $\epsilon_{_{\rm B}}=
+U'_{_{\rm B}}/U'_\ditto{th}$."
+" In the event that all of the low-euecrev svuchrotron (with integrated hIWmuuinositv L.) aud high-energy inverse Compton (£,,.) cussion is powered by clectrous accelerated in the dissipation region. then where Ut is the comoving euergv density of seed photous for the inverse Couptou process."," In the event that all of the low-energy synchrotron (with integrated luminosity $L_\ditto{S}$ ) and high-energy inverse Compton $L_\ditto{IC}$ ) emission is powered by electrons accelerated in the dissipation region, then where $U'_{\ditto{soft}}$ is the comoving energy density of seed photons for the inverse Compton process."
+" For FSRQs. Ut is inostlv contributed by external photous from the DLR and obscuring torus (as opposed to svuchrotrou seed photons. that dominate in BL Lacs). so that PU)... where U,,,, was defined iu refsec:coutent.."," For FSRQs, $U'_{\ditto{soft}}$ is mostly contributed by external photons from the BLR and obscuring torus (as opposed to synchrotron seed photons, that dominate in BL Lacs), so that $U'_\ditto{soft}\sim\Gamma^2\,U_\ditto{BLR}$ , where $U_\ditto{BLR}$ was defined in \\ref{sec:content}. ."
+ With help from eq. (11)), With help from eq. \ref{e: U_cr}) )
+" for the form of U we have where we iacde liberal use of the fact that £,~Lo for our jet model in FSROs.", for the form of $U'_\ditto{th}$ we have where we made liberal use of the fact that $ L_{_{\rm J}}\sim L_{_{\rm acc}}$ for our jet model in FSRQs.
+" It follows that which is consistent with the value e,~0.1 chosen iu 83.2..", It follows that which is consistent with the value $\epsilon_\ditto{B}\sim0.1$ chosen in \ref{sec:jetmodel}.
+ Comparable values result from detailed modeling of blazar spectral cnerey distributions(2)., Comparable values result from detailed modeling of blazar spectral energy distributions.
+. This coufiruis that the electromagnetic contribution to the jet energv flux at the dissipation scale is ucelieible compared to the proton kinetic flux., This confirms that the electromagnetic contribution to the jet energy flux at the dissipation scale is negligible compared to the proton kinetic flux.
+" The maguctic energy fraction €,, computed ii is indepeudent from the black hole mass AA.", The magnetic energy fraction $\epsilon_\ditto{B}$ computed in is independent from the black hole mass $M_{\bullet}$ .
+" The apparent lack of depondeuce of ou AL, may be musleading.", The apparent lack of dependence of $\epsilon_{_{\rm B}}$ on $M_{\bullet}$ may be misleading.
+" Iu fact. the fraction © of e,accretion Iuninosity that is re-processed aud isotropized in the BLR will"," In fact, the fraction $\xi$ of accretion luminosity that is re-processed and isotropized in the BLR will"
+of the OGLE-IL data. their age estimates are only reliable [or ages Z1 Gyr.,"of the OGLE-II data, their age estimates are only reliable for ages $\lesssim 1$ Gyr."
+ Consequently. the comparison between the ΟΙΤΗ CMD-based ages and our age determinations in this paper are only valid for ages S1 Cyr. as is also evident from the lack of data points for logCXge/vr)z9.0 in Fig.," Consequently, the comparison between the OGLE-II CMD-based ages and our age determinations in this paper are only valid for ages $\lesssim
+1$ Gyr, as is also evident from the lack of data points for $\log({\rm
+Age/yr}) \gtrsim 9.0$ in Fig."
+ Ebb. For the older sample clusters. for which age estimates are available based on spectral features or CMD analvsis. the comparison is extended on a one-by-one basis in Section 2.3..," \ref{lmccf.fig}b b. For the older sample clusters, for which age estimates are available based on spectral features or CMD analysis, the comparison is extended on a one-by-one basis in Section \ref{other.sec}."
+ The systematic cllect shown by the 1102 cata seems to have disappered when we compare our results to those of the OGLIZ-HE team: the resulting slope is 1.05+ 0.05. with an x-axis intercept at log(Age/vr)=—0.11+0.39.," The systematic effect shown by the H03 data seems to have disappered when we compare our results to those of the OGLE-II team; the resulting slope is $1.05 \pm 0.05$ , with an x-axis intercept at $\log(
+\mbox{Age/yr} ) = -0.11 \pm 0.39$."
+ This excellent agreement indicates that our assumption that the overall extinction is well represented by 0|)=0.10 mag (assuming the Calzetti attenuation law) is well justified: a change in extinction of |AL(B1|=0.05 would correspond to an olfset [rom the line of equality of lAlogCAge/vr)|20.10.," This excellent agreement indicates that our assumption that the overall extinction is well represented by $E(B-V) = 0.10$ mag (assuming the Calzetti attenuation law) is well justified; a change in extinction of $|\Delta E(B-V)|=0.05$ would correspond to an offset from the line of equality of $|\Delta\log(
+\mbox{Age/yr} )| \simeq 0.10$."
+ This elect is a direct. result. of the age-extinction. and. possibly (to a lesser extent) also of the age-metallicity degeneracy (see Section 2.2)).," This effect is a direct result of the age-extinction, and possibly (to a lesser extent) also of the age-metallicity degeneracy (see Section \ref{degeneracies.sec}) )."
+ The systematic effect. with respect to the LO3 results is still visible. although the scatter in the results is perhaps not surprisingly because we used the same data set much smaller than when we compare our results to those of OGLE-LL," The systematic effect with respect to the H03 results is still visible, although the scatter in the results is – perhaps not surprisingly because we used the same data set – much smaller than when we compare our results to those of OGLE-II."
+ Fable 1. contains the full numerical details for a quantitative comparison among the age determinations based. on the samples. discussed. here., Table \ref{regression.tab} contains the full numerical details for a quantitative comparison among the age determinations based on the samples discussed here.
+ We note that. any remaining offset between our results and those of the OCGLIZ-IL team. despite having used the same assumptions for the total extinction. towards the clusters. may have been introduced. by the exact value for the distance. moculus assumed by Pietrzviásski Ucdalski (2000). rnAl=18.23 mae. which is an essential ingredient for age determinations based on C€MD fitting.," We note that any remaining offset between our results and those of the OGLE-II team, despite having used the same assumptions for the total extinction towards the clusters, may have been introduced by the exact value for the distance modulus assumed by Pietrzyńsski Udalski (2000), $m-M
+= 18.23$ mag, which is an essential ingredient for age determinations based on CMD fitting."
+ Ifthey had. assumed a nominal mAL=18.50 mag. their resulting age estimates would," Ifthey had assumed a nominal $m-M = 18.50$ mag, their resulting age estimates would"
+been the subject of detailed weak-lensing analyses.,been the subject of detailed weak-lensing analyses.
+ Studies of weak lensing by clusters of galaxies probe the projected matter distributions on spatial scales 1 Mpc. significantly larger than the olfsets between the X-ray and lensing centres determined for the non-cooling Low clusters (Section 4.4).," Studies of weak lensing by clusters of galaxies probe the projected matter distributions on spatial scales $\sim 1$ Mpc, significantly larger than the offsets between the X-ray and lensing centres determined for the non-cooling flow clusters (Section 4.4)."
+ From their study of Abell 2218. Squires (1996) determine a lower bound to the mass within an SOO kpc (3.5 arcmin) radius of the cluster centre of 7.8E1.4.10HAL...," From their study of Abell 2218, Squires (1996) determine a lower bound to the mass within an 800 kpc (3.5 arcmin) radius of the cluster centre of $7.8 \pm 1.4 \times 
+10^{14}$."
+ This compares to an A-ray-detoermined mass within same region of: 4.31 ⋅⋅495.⊳↖lotAL.., This compares to an X-ray-determined mass within the same region of $4.31^{+0.32}_{-0.40} \times 10^{14}$.
+.Phe DNESquires efhe weak-lensing mass thus exceeds the X-ray determined value for the central SOO kpe by a factor ~1.8., The Squires (1996) weak-lensing mass thus exceeds the X-ray determined value for the central 800 kpc by a factor $\sim 1.8$.
+ Within a smaller 400 kpe (radius) region. however. Squires (1996) determine a mass of 24+£06.Lot (this value ias been estimated [rom their bie.," Within a smaller 400 kpc (radius) region, however, Squires (1996) determine a mass of $2.4 \pm 0.6 \times 10^{14}$ (this value has been estimated from their Fig."
+ 16) in good agreement with the X-ray measurement of 2.290151071 from he analvsis presented here., 16) in good agreement with the X-ray measurement of $2.29^{+0.17}_{-0.22} \times 10^{14}$ from the analysis presented here.
+ Smail (1997) also present results from a weak-lensing study of LIST images of Abell 2218 from which they measure a mass within the central 400 kpe (radius) region of the cluster of 2.10280.38-1074. iin excellent agreement with the Squires {1996) and X-ray results.," Smail (1997) also present results from a weak-lensing study of HST images of Abell 2218 from which they measure a mass within the central 400 kpc (radius) region of the cluster of $2.10 \pm 0.38 \times 10^{14}$, in excellent agreement with the Squires (1996) and X-ray results."
+ The weak-lensing anc X-ray results for Abell 2218 hus suggest an unusual(non-isothermal) projected: mass xolile between radii of 400 and SOO kpe with a density &racient [latter than r, The weak-lensing and X-ray results for Abell 2218 thus suggest an unusual (non-isothermal) projected mass profile between radii of 400 and 800 kpc [with a density gradient flatter than $r^{-1}$.
+ We note that the weak-lensing mass for Abell 2218 within 500 kpe (Squires. 1996) appears high. given that this mass is comparable to the value determined. for Abell2890 ibawithin a similar aperture (see below). despite the Lact the 20 keV. X-ray luminosity of Abell 2390 is ~4 times higher than that of Abell 2218.," We note that the weak-lensing mass for Abell 2218 within 800 kpc (Squires 1996) appears high given that this mass is comparable to the value determined for Abell 2390 within a similar aperture (see below), despite the fact that the $2-10$ keV X-ray luminosity of Abell 2390 is $\sim 4$ times higher than that of Abell 2218."
+ Such conclusions are not significantly allected by the remodelling of the X-ray mass profiles with the smaller core radii. discussed in Section 4.6.]," Such conclusions are not significantly affected by the remodelling of the X-ray mass profiles with the smaller core radii, discussed in Section 4.6.]"
+ From their weak-lensing study of Abell 2163. Squires al. (LO9Ta)determine (from their Fig.," From their weak-lensing study of Abell 2163, Squires (1997a) determine (from their Fig."
+ 5 )a mass within a 200 aresee (STO kpe) radius aperture of ~510H (with a factor 2 uncertainty). in good agreement with the N-rav determined. value. Mx=760%motOM... reported here.," 5) a mass within a 200 arcsec (870 kpc) radius aperture of $\sim 5 \times 10^{14}$ (with a factor $\sim 2$ uncertainty), in good agreement with the X-ray determined value, $M_{\rm X, weak} = 
+7.6^{+0.8}_{-0.7} \times 10^{14}$, reported here."
+" In addition. [rom ir study of the cooling-Iow cluster Abell 2390. Squires a£.n(1997b) e""measure (from their Fig."," In addition, from their study of the cooling-flow cluster Abell 2390, Squires (1997b) measure (from their Fig."
+ 3) a mass within r=200 arcsec kpe) of 10410HAL... in excellent agreement with the X-ray determined. value of S2TOSU10HAL...," 3) a mass within $r = 200$ arcsec (940 kpc) of $10 \pm 4 \times 10^{14}$, in excellent agreement with the X-ray determined value of $8.2^{+8.7}_{-3.0} \times 10^{14}$."
+ From their weak lensing analvsis of LIST images. Smail (1997) measure a mass within the central 400 kpc (raclius) region of Abell2744 CACIHIS) of3.7040.645104AL...," From their weak lensing analysis of HST images, Smail (1997) measure a mass within the central 400 kpc (radius) region of Abell 2744 (AC118) of $3.70 \pm 0.64 \times 10^{14}$."
+ Phis compares to the value determined from the X-ray analvsis. presented here of 2.330.51μοι10m (implying a wealk-lensing/X-rav mass ratio of 1.59ΟΙ}ο τμ), This compares to the value determined from the X-ray analysis presented here of $2.33^{+0.31}_{-0.24} \times 10^{14}$ (implying a weak-lensing/X-ray mass ratio of $1.59^{+0.49}_{-0.43}$ ).
+ From their analysis of weak lensing in the most, From their analysis of weak lensing in the most
+dack hole within the jet producing ACN.,black hole within the jet producing AGN.
+ Since the mass of a super massive black role will oulv iuerease with cosmic time. there should be pleuty of objects in the local universe with black holes in their ceutres which are of comparable or even greater mass han those producing stroug radio sources at high redslüft.," Since the mass of a super massive black hole will only increase with cosmic time, there should be plenty of objects in the local universe with black holes in their centres which are of comparable or even greater mass than those producing strong radio sources at high redshift."
+ Towever. at +~0 they only eive rise to weaker jets creating FRItype sources or they are quiesceut altogether.," However, at $z \sim 0$ they only give rise to weaker jets creating FRI-type sources or they are quiescent altogether."
+ The uost powerful FRILtype objects at low z are always found in poor eroups of galaxies., The most powerful FRII-type objects at low $z$ are always found in poor groups of galaxies.
+ The faster virialisation of the eas in laree objects like ealaxy clusters as opposed to sxinaller groups could prevent material from reaching the ceutre of potential radio source josts within these rich euvironnmients at low redshift and thereby depriving the black roles iu these objects of fuel (e.g. Ellingson. Creen Yee 1991).," The faster virialisation of the gas in large objects like galaxy clusters as opposed to smaller groups could prevent material from reaching the centre of potential radio source hosts within these rich environments at low redshift and thereby depriving the black holes in these objects of fuel (e.g. Ellingson, Green Yee 1991)."
+ If this scenario is correct. then im anv attempt to model the cosmological evolution of the radio galaxy vopulation from hierarchical structure formation it is necessary to take into account nof only the presence of a massive black hole but also the availability of fuel for the ACN in potential progenitors of powerful radio galaxies.," If this scenario is correct, then in any attempt to model the cosmological evolution of the radio galaxy population from hierarchical structure formation it is necessary to take into account not only the presence of a massive black hole but also the availability of fuel for the AGN in potential progenitors of powerful radio galaxies."
+ To decide which of the two scenarios for the cosnological evolution of the FRII source population is the more appropriate. particularly at lower radio huuinosities thaw those covered by the sample of Laing et al. (," To decide which of the two scenarios for the cosmological evolution of the FRII source population is the more appropriate, particularly at lower radio luminosities than those covered by the sample of Laing et al. ("
+1983). it is necessary to use fainter complete samples in the model comparison.,"1983), it is necessary to use fainter complete samples in the model comparison."
+ These will be available in the near future aud should help to identify the progenitor population of powerful radio galaxies aud to coustrain the cosiuological evolution of their environments., These will be available in the near future and should help to identify the progenitor population of powerful radio galaxies and to constrain the cosmological evolution of their environments.
+ We have poiuted out iu section 2 that the expansion of the cocoon of ΕΠΗ sources drives a strong bow shock iuto the surroundiue ICAL, We have pointed out in section 2 that the expansion of the cocoon of FRII sources drives a strong bow shock into the surrounding IGM.
+ The ICAL will be compressed aud heated by the shock and its properties will change significantly., The IGM will be compressed and heated by the shock and its properties will change significantly.
+ The problems of a strong shock expaucing iuto a gaseous atmosphere was first solved for the spherical case in a uniform atmosphere by Sedov (1959)., The problem of a strong shock expanding into a gaseous atmosphere was first solved for the spherical case in a uniform atmosphere by Sedov (1959).
+ A similar situation but with a constant enerev input to a cavity expanding behind the bow shock was investigated by Dyson. Falle Perry (1980).," A similar situation but with a constant energy input to a cavity expanding behind the bow shock was investigated by Dyson, Falle Perry (1980)."
+ Both solutious are self-similar aud since the expansion of the cocoon aud bow shock of FRIT radio sources should be self-similar as well (Ixaiser Alexander 1997. Section 2). we expect that these solutions can be extended to this case.," Both solutions are self-similar and since the expansion of the cocoon and bow shock of FRII radio sources should be self-similar as well (Kaiser Alexander 1997, Section 2), we expect that these solutions can be extended to this case."
+" The main differcuce is the clongated shape of the cocoons in ERIT sources,", The main difference is the elongated shape of the cocoons in FRII sources.
+ By assiuniue rotational sviuinetry about the jet axis we can reduce the umber of spatial dimension in the problem by oue., By assuming rotational symmetry about the jet axis we can reduce the number of spatial dimension in the problem by one.
+ Transferring the usual equations goveruiug the gas flow between bow shock aud cocoon to a selfsimularly expanding coordinate system it is possible to transform these equations to a set of partial differential equatious oe1 two indepeudent (spatial) coordinates., Transferring the usual equations governing the gas flow between bow shock and cocoon to a self-similarly expanding coordinate system it is possible to transform these equations to a set of partial differential equations in two independent (spatial) coordinates.
+ These can then be solved umucrically., These can then be solved numerically.
+ To proceed we have to asstme the ecometrical shape of the bow shock., To proceed we have to assume the geometrical shape of the bow shock.
+ The bow shocks in FRIT sources can not be observed directly., The bow shocks in FRII sources can not be observed directly.
+ IToxcever. we do not expect the laver of shocked gas between bow shock aud cocoon to be very thick and so it secs reasonable to assume a prolate ellipsoid shape for the bow shock which resembles the observed shapes of cocoons.," However, we do not expect the layer of shocked gas between bow shock and cocoon to be very thick and so it seems reasonable to assume a prolate ellipsoid shape for the bow shock which resembles the observed shapes of cocoons."
+ At the bow shock surface we assume strong shock conditions which. together with the assumed power law profile for the uuperturbed density distribution of the ICM. defines the initial conditions for the integration.," At the bow shock surface we assume strong shock conditions which, together with the assumed power law profile for the unperturbed density distribution of the IGM, defines the initial conditions for the integration."
+ The solution is then propagated numerically iuwards from the bow shock and stopped at the contact discontinuity., The solution is then propagated numerically inwards from the bow shock and stopped at the contact discontinuity.
+ The shape of this surface is uot known a priori but can be found using the condition that the eas at the contact discontinuity is not moving im the direction perpendicular to this surface., The shape of this surface is not known a priori but can be found using the condition that the gas at the contact discontinuity is not moving in the direction perpendicular to this surface.
+ Figure 5 shows the result of the integration for various values of the exponent of the external deusitv distribution. ον," Figure \ref{fig:dengray} shows the result of the integration for various values of the exponent of the external density distribution, $\beta$."
+ The shape of the external deusitv distribution is reflected. iu the flow region., The shape of the external density distribution is reflected in the flow region.
+ For steep exterual deusitv gradients. 9~2. the deusity distribution within the flow region is even steeper than that of the uuperturbed gas," For steep external density gradients, $\beta \sim
+2$, the density distribution within the flow region is even steeper than that of the unperturbed gas"
+In the previous sections we have shown that a number of features are detected in the individual sources spectra including the ΠΑΟ feature as well as the wwater ice absorption feature.,In the previous sections we have shown that a number of features are detected in the individual sources spectra including the HAC feature as well as the water ice absorption feature.
+ There are hints of some other features in (he sources such as marked in rellig pecs... bultheyarelypicallyo flowsigni ficance.," There are hints of some other features in the sources such as marked in \\ref{fig_specs}, but they are typically of low significance."
+" Pheseinceludeinparticular he aabsorplion complex (likely including the wwaler ice feature). the 6.85 and ΠΠΑΟ features and the CCO [eature discussed in re[sec,.o.."," These include in particular the absorption complex (likely including the water ice feature), the 6.85 and HAC features and the CO feature discussed in \\ref{sec_co}."
+F'oaddressthequestionofwhetherornollhesearepresentinoursample. westackedthespectraasshou [ig.lack.," To address the question of whether or not these are present in our sample, we stacked the spectra as shown in \\ref{fig_stack}."
+.Decausetherearetoofewsourcestolestlherobustnessoflhestackbglalingrandomlkyhal fo fil.," Because there are too few sources to test the robustness of the stack by taking randomly half of it, we used a different technique which involves removing each existing pair of sources, or 28 possibilities."
+ t reffi, This implies that features that appear in each of the 28 stacks need to be present in the majority of sources in order not to disappear in at least some of the combinations.
+gLachzoomshowsthatinthestackedspeetra. weseethatthe3.140n IHLLAC: feature is quite prominent. unsurprising given that we detect it in four of the individual spectra.," \\ref{fig_stackzoom} shows that in the stacked spectra, we see that the HAC feature is quite prominent, unsurprising given that we detect it in four of the individual spectra."
+ Potentially there is also wwaler ice absorption and as well as potentially the ΡΡΑΙΙ emission given some eexcess., Potentially there is also water ice absorption and as well as potentially the PAH emission given some excess.
+ The [feature is badly fit al the lower wavelength end lor most. but not all Ρονations.," The feature is badly fit at the lower wavelength end for most, but not all permutations."
+ MIPS16059 is probably responsible for some fraction of these (see .44.4). the rest is probably due to edge effects in sources without anv clear features. ," MIPS16059 is probably responsible for some fraction of these (see 4.4), the rest is probably due to edge effects in sources without any clear features. \\ref{fig_stack},"
+however. suggeststhatthe6 aabsorption complex is there in most sources.," however, suggests that the absorption complex is there in most sources."
+ The most prominent such feature is seen in MIDS22633 where the aabsorption is clear in (he individual spectrum as well., The most prominent such feature is seen in MIPS22633 where the absorption is clear in the individual spectrum as well.
+ One thing to keep in mind is (hat in a number of the sources. the ~ features in particular are not individually detected due to their Falling inside the unreliable IRS modules overlap region.," One thing to keep in mind is that in a number of the sources, the $\sim$ features in particular are not individually detected due to their falling inside the unreliable IRS modules overlap region."
+ This uncertainty. as well as the likely PPAIT make it difficult to estimate the average wwaler ice and IHLAC: feature depths.," This uncertainty, as well as the likely PAH make it difficult to estimate the average water ice and HAC feature depths."
+ rellig Jackzoomsuggeststhallheseare(72.9) ~ 00.22 and (731) ~ 00.30. but it should be born in mind that these are fairly uncertain.," \\ref{fig_stackzoom} suggests that these are $\langle\tau_{3.0}\rangle$ $\sim$ 0.22 and $\langle\tau_{3.4}\rangle$ $\sim$ 0.30, but it should be born in mind that these are fairly uncertain."
+Our method however is not to directly compare the peculiar velocity and acceleration of the Local Group: insteacl. we use the observed growth of the dipole to obtain these constraints.,"Our method however is not to directly compare the peculiar velocity and acceleration of the Local Group; instead, we use the observed growth of the dipole to obtain these constraints."
+" The Two Micron. All Sky Survey (2A\LASS. ?)) is the first. near-intrarecl survey of the whole skv (covering of the celestial sphere). and was performed in the period L9972001 in the J(1.25 yon). H(1.65jm) and Av,(2.16jm) bands. with the use of twin 1.3-m ground-based telescopes."," The Two Micron All Sky Survey (2MASS, \citealt{Skr}) ) is the first near-infrared survey of the whole sky (covering of the celestial sphere), and was performed in the period 1997–2001 in the $J\, (1.25\, \mu\mathrm{m})$ , $H\, (1.65\, \mu\mathrm{m})$ and $K_s\, (2.16\, \mu\mathrm{m})$ bands, with the use of twin 1.3-m ground-based telescopes."
+ All the data from the survey are available through the NASA/IPAC Infrared Science The main outcome of (his project are wo photometric catalogs: of point sources (PSC). containing about 471 million objects. ancl of extended ones (ASC). with more than 1.6million objects. mainiv galaxies (> 98%) and some diffuse Galactic sources (2)..," All the data from the survey are available through the NASA/IPAC Infrared Science The main outcome of this project are two photometric catalogs: of point sources (PSC), containing about 471 million objects, and of extended ones (XSC), with more than 1.6million objects, mainly galaxies $>98\%$ ) and some diffuse Galactic sources \citep{Jar04}."
+" The NSC. which was used for the purpose of our analvsis. is complete lor sources brighter (han A.13.5 mag (2.7mJy) ancl resolved diameters larger than ~10 15""."," The XSC, which was used for the purpose of our analysis, is complete for sources brighter than $K_s\simeq13.5$ mag $\sim2.7\,\mathrm{mJy}$ ) and resolved diameters larger than $\sim10$ – $15''$."
+ The near-infrared {hix is particularly useful for the purpose of large-scale structure studies as it samples the old stellar population. and hence the bulk of stellar mass. ancl it is minimally allectecl by dust in the Galactie plane (?)..," The near-infrared flux is particularly useful for the purpose of large-scale structure studies as it samples the old stellar population, and hence the bulk of stellar mass, and it is minimally affected by dust in the Galactic plane \citep{Jar04}."
+ An additional advantage of using 2MASS data. especially in the context of calculating the flux dipole. for which magnitudes are used. is the elobal photometric uniformity of the catalog. which was enforced by nightly photometric calibration (o an extensive set of standard star fields.," An additional advantage of using 2MASS data, especially in the context of calculating the flux dipole, for which magnitudes are used, is the global photometric uniformity of the catalog, which was enforced by nightly photometric calibration to an extensive set of standard star fields."
+ On the other hand. as any survey. 2MAÀSS is not perfect.," On the other hand, as any survey, 2MASS is not perfect."
+ It is biased against optically blue and low surface brightness galaxies. such as clwarls. but sensitive to the early tvpe. bulge-dominated ones.," It is biased against optically blue and low surface brightness galaxies, such as dwarfs, but sensitive to the early type, bulge-dominated ones."
+ As the former have very small luminosities and masses. (heir possible underrepresentation in (he catalog should not influence significantly our results.," As the former have very small luminosities and masses, their possible underrepresentation in the catalog should not influence significantly our results."
+ The 2MASS photometry olfers several types of magnitudes! [or extended objects. depending on the tvpe of aperture tiised ete.," The 2MASS photometry offers several types of `magnitudes' for extended objects, depending on the type of aperture used etc."
+" Throughout the whole analvsis we use the 20mag/sq."" isophotal fiducial elliptical aperture magnitudes. which are defined as magnitudes inside the elliptical isophote corresponding. to a surface. brightness. of. paga=i20Hmag/sq..."," Throughout the whole analysis we use the $20\,\mathrm{mag\slash{}sq.''}$ isophotal fiducial elliptical aperture magnitudes, which are defined as magnitudes inside the elliptical isophote corresponding to a surface brightness of $\mu_\mathrm{band}=20\,\mathrm{mag\slash{}sq.}''$."
+ We- prefer. those to the Kron- ones as the, We prefer those to the Kron ones as the
+ , 
+at the Effelsberg radio telescope.,at the Effelsberg radio telescope.
+ The receiver output was split ou two chanuels., The receiver output was split on two channels.
+ The signals οι one channel were processed by a FPGA processor (Altera Stratix S80) and then seut to the total power detector., The signals from one channel were processed by a FPGA processor (Altera Stratix S80) and then sent to the total power detector.
+ The signals frou the second chiaunel were applied straight to the total power detector., The signals from the second channel were applied straight to the total power detector.
+ The anewidth of the signals applied o the total power detector was equal to 20 MIIz., The bandwidth of the signals applied to the total power detector was equal to 20 MHz.
+ Tjose Channels provied two radio telescope outputs: oue output with REI nitfisation and another without., These channels provided two radio telescope outputs: one output with RFI mitigation and another without.
+ Pairs of radio source scans from puts were made siuultaueouslv., Pairs of radio source scans from these outputs were made simultaneously.
+" Because power detecor Was an iuteeral part of the radio telescope backend eqiipineut. the aleorithiu Hupleimeated nu FPGA oilv processed the IF (intoruediate frequc10} signal with the aim of ""cleaning it of RFI."," Because the total power detector was an integral part of the radio telescope backend equipment, the algorithm implemented in FPGA only processed the IF (intermediate frequency) signal with the aim of “cleaning"" it of RFI."
+" The analogue input signal was (lelized in 12-bit ADC wih. LOATsamples speed. then processec in FPGA, transformed back iuto analogue form aud applied to the total power detector."," The analogue input signal was digitized in 12-bit ADC with 40Msamples speed, then processed in FPGA, transformed back into analogue form and applied to the total power detector."
+ Winsorization of the sigτα]. in temporal domain (5= 0.05) was mipleimoeuted., Winsorization of the signal in temporal domain $\gamma=0.05$ ) was implemented.
+ Fie., Fig.
+ 5 shows eight scans of the radio source. cach scan represeuted by two panes: the top panel - scan with RFT iitigation. the lower panel - without RFI iitieation.," 8 shows eight scans of the radio source, each scan represented by two panels: the top panel - scan with RFI mitigation, the lower panel - without RFI mitigation."
+ Fig., Fig.
+ 9 displavs radio nuages of the source built using scans simular to those in Fig., 9 displays radio images of the source built using scans similar to those in Fig.
+ Nm: the left panel - without RFI mitigation. the right panel - with REI nütisation.," 8: the left panel - without RFI mitigation, the right panel - with RFI mitigation."
+ A new pulsar machine PUMA-2 has been lustalled at the Westerbork svuthesis racio clescope (WSRT)., A new pulsar machine PUMA-2 has been installed at the Westerbork synthesis radio telescope (WSRT).
+ The radio telescope works in tic(Larray mode in which all Ll signals from autenuas are added in phase. nee there is one outut (iu reality with two polarizations) as for a sineole dish.," The radio telescope works in tied-array mode in which all 14 signals from antennas are added in phase, i.e., there is one output (in reality with two polarizations) as for a single dish."
+ The 20 MIIz-baseband signals from each of the cight frequency channels of WSRT are dieiized (8 bit) and stored in the mass storage svsteu1 which has sufficient lard disk capacity to supyort at least 21 hours of contiuuous observations., The 20 MHz-baseband signals from each of the eight frequency channels of WSRT are digitized (8 bit) and stored in the mass storage system which has sufficient hard disk capacity to support at least 24 hours of continuous observations.
+ Signal processi1ο can therefore be undertakenoff-linc., Signal processing can therefore be undertaken.
+ Iu our experiment a lock of data recorded during 10 sec. LO«109 eiebt-bit suuples/sec. was used.," In our experiment a block of data recorded during 10 sec, $40
+\times 10^6$ eight-bit samples/sec, was used."
+ All REI mitigation xocessiue aud the total power cetector (TPD) were realized entirelv in software duringoff-line processing., All RFI mitigation processing and the total power detector (TPD) were realized entirely in software during processing.
+ The estimate of variance with exponenial weieliting was used., The estimate of variance with exponential weighting was used.
+ Fig., Fig.
+ 10 displays the resits of processing., 10 displays the results of processing.
+ Upper row. loft panel: TPD outputs for two polarizatioIs calculate‘dl from the raw data with RFI. right panel: TPD outputs. RET removed.," Upper row, left panel: TPD outputs for two polarizations calculated from the raw data with RFI, right panel: TPD outputs, RFI removed."
+ Middle row. left pamcl: exaunuple of a time fragment of the power sctim with RET: right uel: the same time fraemeut. REI removed.," Middle row, left panel: example of a time fragment of the power spectrum with RFI; right panel: the same time fragment, RFI removed."
+ Lower row shows pulsar ]xofiles after de-dispersion and olding at the pulsar period. left panel: pulsar profile averaged over 10 sec frou both polarizatious of raw data. middle pancl: similar pt]sar profic. RFI removed. right panel: pulsar profile restored with observational data obtained at 1120 MITz without any RFI which is put here for compariso1 with the profile iu the nidle panel.," Lower row shows pulsar profiles after de-dispersion and folding at the pulsar period, left panel: pulsar profile averaged over 10 sec from both polarizations of raw data, middle panel: similar pulsar profile, RFI removed, right panel: pulsar profile restored with observational data obtained at 1420 MHz without any RFI which is put here for comparison with the profile in the middle panel."
+ adio source DA?0 was οserved at WSRT at a ceutral frequency of 357 ΑΠ with the baudwekth equal to 20 MIIz iu the preseuce of stroug RET., Radio source DA240 was observed at WSRT at a central frequency of 357 MHz with the bandwidth equal to 20 MHz in the presence of strong RFI.
+ The RFI uitigatio1 svsteiin (REIAIS) Was σοςL for processing (Baan.Fridiuau&Millewar 2001)., The RFI mitigation system (RFIMS) was used for processing \citep{baan04}.
+. Analogue baseband signals were digitized (12bit. ADC. 0. Myauples/sec]. processed iu FPGA (Altera. StratixSsd) aud ΠΑΝΤΟΥned back to analogue form for subsequeut processiic in the WSRT correlator.," Analogue baseband signals were digitized (12bit ADC, 40 Msamples/sec), processed in FPGA (Altera StratixS80) and transformed back to analogue form for subsequent processing in the WSRT correlator."
+ The aleorithiu used for the removal of RET was similar to the algorithu with exponential weighting. except that the varieuice was not used as an output. instead “cleanedo. exponentially weighted signals were applied ο the correlator.," The algorithm used for the removal of RFI was similar to the algorithm with exponential weighting, except that the variance was not used as an output, instead “cleaned"", exponentially weighted signals were applied to the correlator."
+ Radio images of the source DÀ210 are s1OW]. in Fie., Radio images of the source DA240 are shown in Fig.
+ 1., 11.
+ Upper row. left panel: nage without RFI nuiligation: right panel: image with RFI nitieatk13.," Upper row, left panel: image without RFI mitigation; right panel: image with RFI mitigation."
+ Lower row: central parts of the 1nage preseute din the same order., Lower row: central parts of the image presented in the same order.
+ The streched form of the svthesized images is explained by the fact that the observations lasted & hours. iustead of a full 12 liour aperture svuthesis evcle.," The stretched form of the synthesized images is explained by the fact that the observations lasted 8 hours, instead of a full 12 hour aperture synthesis cycle."
+ RFI uitigation was implementec on cach of Ll x‘aco telescopes at WSRT., RFI mitigation was implemented on each of 14 radio telescopes at WSRT.
+ This ay give rise fo sone distortions Decause of the dierence of equipinout characteristics at the dierent autennas., This may give rise to some distortions because of the difference of equipment characteristics at the different antennas.
+" Other observations were lace specifically to judge the ""toxicitv of the RFI nitieatkon procedure."," Other observations were made specifically to judge the “toxicity"" of the RFI mitigation procedure."
+ Two pairs of radio images of radio source [C35[.Ur were svuthesized: oue was observed at 14120 MIIz with the RFI uütieatiou πλΤΟ suxl without it aud another at 315 MITz. also with the RFT nütisatiou svstem aud witlout," Two pairs of radio images of radio source 4C34.47 were synthesized: one was observed at 1420 MHz with the RFI mitigation system and without it and another at 345 MHz, also with the RFI mitigation system and without"
+In radio bright samples. GPS quasars are founc to have core-jet or complex structure. while GPS galaxies are found to have larger sizes with jets and lobes on both sides of à putative center of activity (Stanghellini et al.,"In radio bright samples, GPS quasars are found to have core-jet or complex structure, while GPS galaxies are found to have larger sizes with jets and lobes on both sides of a putative center of activity (Stanghellini et al."
+ LOOT)., 1997).
+ Although observations at another frequency. are needed to confirm. their. classification. allmost all raclio-bright GPS galaxies from Stanghellini et al (1997) can be classified. as CSOs.," Although observations at another frequency are needed to confirm their classification, allmost all radio-bright GPS galaxies from Stanghellini et al (1997) can be classified as CSOs."
+ The morphological cichotomy of GPS galaxies and quasars. and their very dillerent redshift distributions make it likely that GPS galaxies ancl quasars are not related to each other and just happen to have similar radio spectra.," The morphological dichotomy of GPS galaxies and quasars, and their very different redshift distributions make it likely that GPS galaxies and quasars are not related to each other and just happen to have similar radio spectra."
+ ]t has been speculated that GPS quasars are a subset. of Hat spectrum quasars in general (eg., It has been speculated that GPS quasars are a subset of flat spectrum quasars in general (eg.
+ Suellen et al., Snellen et al.
+ 1999a)., 1999a).
+ In addition. if galaxies and quasars were to be unified. by orientation. due to changes in its observed. radio spectrum. (Snellen et al.," In addition, if galaxies and quasars were to be unified by orientation, due to changes in its observed radio spectrum (Snellen et al."
+ 1998c). it is not expected that a GPS galaxy observed at a small viewing angle would be seen as a GPS quasar.," 1998c), it is not expected that a GPS galaxy observed at a small viewing angle would be seen as a GPS quasar."
+ Not all CSOs are GPS sources., Not all CSOs are GPS sources.
+ The contribution of the (possibly variable) Dat spectrum core can be significant and outshine the convex spectral shape produced. by the mini-lobes., The contribution of the (possibly variable) flat spectrum core can be significant and outshine the convex spectral shape produced by the mini-lobes.
+ This can be due to a small viewing angle towards the object. causing the Doppler boosted core and. fast. moving jet. which feeds the approaching mini-lobe. to be important (Snellen et al.," This can be due to a small viewing angle towards the object, causing the Doppler boosted core and fast moving jet, which feeds the approaching mini-lobe, to be important (Snellen et al."
+ 1998ce)., 1998c).
+ An example of such a CSO. possibly observed at a small viewing angle. is 1413|135 (Perlman οἱ al.," An example of such a CSO, possibly observed at a small viewing angle, is 1413+135 (Perlman et al."
+ 1994)., 1994).
+ In addition. the jets feeding the mini-Iobes can be significantlyIn curved. for example in 2352|495 by precession teacdhead et al.," In addition, the jets feeding the mini-lobes can be significantly curved, for example in 2352+495 by precession (Readhead et al."
+ 1996)., 1996).
+ This can cause parts of the jet to move at an angle close to the line of sight. with significant Doppler boosting as a result.," This can cause parts of the jet to move at an angle close to the line of sight, with significant Doppler boosting as a result."
+ In both cases the large contrast between the approaching and receding parts of the racio source makes it also increasinely dillicult to identily the object asa CSO., In both cases the large contrast between the approaching and receding parts of the radio source makes it also increasingly difficult to identify the object as a CSO.
+ Figure 13. shows the number of galaxies and. quasars. in ow faint GPS sample. classified. as CJ. CSO. CX and those not. possible to classify.," Figure \ref{class} shows the number of galaxies and quasars, in our faint GPS sample, classified as CJ, CSO, CX and those not possible to classify."
+ ALL three objects classified as CSOs are optically identified with galaxies., All three objects classified as CSOs are optically identified with galaxies.
+ Although this is in agreement with the findings of Stanghellini ct al (1997) for the raclio-bright sample. it should be noted that for only 4 quasars was it possible to make a classification.," Although this is in agreement with the findings of Stanghellini et al (1997) for the radio-bright sample, it should be noted that for only 4 quasars was it possible to make a classification."
+ This is mainly due to the fact that the angular sizes of the quasars are significantly smaller than the angular sizes of the ealaxies., This is mainly due to the fact that the angular sizes of the quasars are significantly smaller than the angular sizes of the galaxies.
+ Six out of 18 classifiable GPS galaxies are found to have €J or CN structures. anc 9 of the classifiable GPS ealaxies are found to have CD structures.," Six out of 18 classifiable GPS galaxies are found to have CJ or CX structures, and 9 of the classifiable GPS galaxies are found to have CD structures."
+ We conclude that the strong morphological cichotomy between GPS galaxies and quasars found. by Stanghellini (1997) in the bright GPS sample. is not as strong in this faint sample.," We conclude that the strong morphological dichotomy between GPS galaxies and quasars found by Stanghellini (1997) in the bright GPS sample, is not as strong in this faint sample."
+ Note. however. that the classification for the majority of the C.J and CD sources is based. on two components and. their relative spectral indices only.," Note, however, that the classification for the majority of the CJ and CD sources is based on two components and their relative spectral indices only."
+ This makes their classification rather tentative., This makes their classification rather tentative.
+ Firstly. a CD source could be erroneously classified as a €. source due to a dillerence in the observed age between the approaching and. receding lobe. causing a difference in observed radio spectrum of the two lobes.," Firstly, a CD source could be erroneously classified as a CJ source due to a difference in the observed age between the approaching and receding lobe, causing a difference in observed radio spectrum of the two lobes."
+ For a separation velocity of 0.4c. as observed for radio bright CLIPS ealaxies (Owsianik and Conway 1908: Owsianik. Conway and. Polatidis 1998). such an age cdilference can be as large as304.," For a separation velocity of 0.4c, as observed for radio bright GPS galaxies (Owsianik and Conway 1998; Owsianik, Conway and Polatidis 1998), such an age difference can be as large as."
+. Secondly. dillerenees in the local environments of the two lobes can also influence the spectra of the two lobes. resulting in an erroneous classification as core-Jet.," Secondly, differences in the local environments of the two lobes can also influence the spectra of the two lobes, resulting in an erroneous classification as core-jet."
+ For example. if only the two lobes had been visible. BISLO|6707 (fig 8)) could have been mistaken for a core-jet source. since the spectral index of the eastern lobe is Hatter than that of the western lobe.," For example, if only the two lobes had been visible, B1819+6707 (fig \ref{fig2}) ) could have been mistaken for a core-jet source, since the spectral index of the eastern lobe is flatter than that of the western lobe."
+ Alulti-frequeney VLBI observations have been presented of a faint sample of GPS sources., Multi-frequency VLBI observations have been presented of a faint sample of GPS sources.
+ ALL 47 sources in the sample were successfully. observed. at 5 Cillz. 26 sources were observed at 15 Gllz. and 20 sources were observed ab 1.6 Cllz.," All 47 sources in the sample were successfully observed at 5 GHz, 26 sources were observed at 15 GHz, and 20 sources were observed at 1.6 GHz."
+ In this wav of the sources have been mapped above and below their spectral peak., In this way of the sources have been mapped above and below their spectral peak.
+ The spectral decomposition allowed: us to classify 3. GPS galaxies as compact svnunetric objects (CSO). 1 galaxy ancl 1 quasar as complex (CX) sources. 2 quasars ancl 5 galaxies as core-jet (€) sources. and 9 galaxies ancl 2 quasars as compact doubles (CD).," The spectral decomposition allowed us to classify 3 GPS galaxies as compact symmetric objects (CSO), 1 galaxy and 1 quasar as complex (CX) sources, 2 quasars and 5 galaxies as core-jet (CJ) sources, and 9 galaxies and 2 quasars as compact doubles (CD)."
+ Twenty-five of the sources. could. not. be classified. 20 because they were too compact.," Twenty-five of the sources could not be classified, 20 because they were too compact."
+ The strong morphological dichotomy of GPS galaxies and quasars found by Stanghellini et al. (, The strong morphological dichotomy of GPS galaxies and quasars found by Stanghellini et al. (
+1997) in their radio bright GPS sample is not so clear in this sample.,1997) in their radio bright GPS sample is not so clear in this sample.
+ However. many of the sources classified as CD and CJ. have a two-component structure. making their classification only tentative.," However, many of the sources classified as CD and CJ have a two-component structure, making their classification only tentative."
+ The authors are ereatful to the stall of the EVN and VLBA for support) of the observing projects., The authors are greatful to the staff of the EVN and VLBA for support of the observing projects.
+ The VLBA is an instrument of the National Raclio Astronomy Observatory. which is operated: by Associated Universities. Inc. under a Cooperative Agreement with the National Science. Foundation.," The VLBA is an instrument of the National Radio Astronomy Observatory, which is operated by Associated Universities, Inc. under a Cooperative Agreement with the National Science Foundation."
+" ""This research was supported by the European Commission. TMIU Access to Large-scale Facilities programme under contract No."," This research was supported by the European Commission, TMR Access to Large-scale Facilities programme under contract No."
+ ERBFAIGEC950012. anc “PAIR Programme. Research Network Contract ERBEAIRAC'P96-0034 “CERE," ERBFMGECT950012, and TMR Programme, Research Network Contract ERBFMRXCT96-0034 “CERES”."
+which is based on the host galaxy extinction models of ?.. 2.. and ?. (see Paper 1. 833.1).,"which is based on the host galaxy extinction models of \citet{Hatano:1998}, \citet{Commins:2004}, , and \citet{Riello:2005} (see Paper I, 3.1)."
+ For the redshift prior we consider two alternative situations., For the redshift prior we consider two alternative situations.
+ Iu the first case. we assume a traditional SN survey strategv. where all SNe have spectroscopic nieasurenieuts to define their redshifts aud to aid in classification.," In the first case, we assume a traditional SN survey strategy, where all SNe have spectroscopic measurements to define their redshifts and to aid in classification."
+" The redshift prior for this scenario is based on the spectroscopic imcasurciuents reported bv ? and 7?/ for the SDSS and SNLS data sets. respectively,"," The redshift prior for this scenario is based on the spectroscopic measurements reported by \citet{Holtzman:2008} and \citet{Astier:2006} for the SDSS and SNLS data sets, respectively."
+ The prior is separately defined. for cach SN. as a Cassian centered on the ameasured τους with width defined by σ.=0.005., The prior is separately defined for each SN as a Gaussian centered on the measured redshift with width defined by $\sigma_z=0.005$.
+ The parameter estimates that SOFT returns in this case can be more or less directly compared to the SN cosmology results clsewhere in the literature., The parameter estimates that SOFT returns in this case can be more or less directly compared to the SN cosmology results elsewhere in the literature.
+ Qur second redshift prior is designed to evaluate how well SOFT can measure redshift aud distance without anv spectroscopic information., Our second redshift prior is designed to evaluate how well SOFT can measure redshift and distance without any spectroscopic information.
+ As a worst case scenario we asstune that there is no spectroscopic information available for anv of he SNe or their host galaxies. and there are iof even photometric redshift estinates frou he host galaxies.," As a worst case scenario we assume that there is no spectroscopic information available for any of the SNe or their host galaxies, and there are not even photometric redshift estimates from the host galaxies."
+ For this case woe want to use an unnforimative redshift prior. reflecting our otal ignorance of each object's redshift.," For this case we want to use an uninformative redshift prior, reflecting our total ignorance of each object's redshift."
+ One possibility would be a fiat distribution over :. iorinalized so that it iuteerates to unitv over the range of redshifts under consideration for cach SN candidate.," One possibility would be a flat distribution over $z$, normalized so that it integrates to unity over the range of redshifts under consideration for each SN candidate."
+ A slightly more realistic prior would ake iuto account the iucreasing volume of space as a function of redshift., A slightly more realistic prior would take into account the increasing volume of space as a function of redshift.
+" For sinplicitv we assume aa flat O4,0 universe so that the volue of a shell at redshift z (per muit solid angle) is This is the wninformative prior that we use throughout this paper. although we also tested the flat z prior aud found no significant changes in our conclusions."," For simplicity we assume a a flat $\Omega_M=0$ universe so that the volume of a shell at redshift z (per unit solid angle) is This is the uninformative prior that we use throughout this paper, although we also tested the flat z prior and found no significant changes in our conclusions."
+ As a first test. we want to evaluate Low well SOFT cau determine a redshift based on the SN light curve alone.," As a first test, we want to evaluate how well SOFT can determine a redshift based on the SN light curve alone."
+ For this task we use the Combined data set. containing 198 photometrically identified Type Ia SNe from subsets SDSS-D aud SNLS-A (sce Table 1)).," For this task we use the Combined data set, containing 198 photometrically identified Type Ia SNe from subsets SDSS-B and SNLS-A (see Table \ref{tab:subsets}) )."
+ Using the muinformative 2? redshift) prior. we determine a composite nembership function frou cach umbodel using Equation L.," Using the uninformative $z^2$ redshift prior, we determine a composite membership function from each model using Equation \ref{eqn:g(D|theta,Mj)}."
+ The membership πιοος are imareinalized over all nuisauce owanmeters tusue)) and combined using he fuzzv AND operator to eet a composite ucmbership fiction over redshift: g(:)., The membership functions are marginalized over all nuisance parameters ) and combined using the fuzzy AND operator to get a composite membership function over redshift: $g(z)$.
+ We locate he peak of this distribution to define :sorr. and assign error bars by measuring the width of the contour that coutains of the arca.," We locate the peak of this distribution to define $z_{SOFT}$, and assign error bars by measuring the width of the contour that contains of the area."
+ Figure 2. plots these SOFT redshift values against the spectroscopically measured redshifts from SDSS and SNLS., Figure \ref{fig:combozplot} plots these SOFT redshift values against the spectroscopically measured redshifts from SDSS and SNLS.
+ The residuals defined bv isoFTtaper Lave an ROMS scatter about zero of 0.015 for the SDSS-A sample. 0.059. for SNLS-B. and 0.051 for the combined sample.," The residuals defined by $z_{SOFT}-z_{spec}$ have an RMS scatter about zero of 0.045 for the SDSS-A sample, 0.059 for SNLS-B, and 0.051 for the combined sample."
+ The uncertainties for these data points are decidedly nou-Caussian aud asvuuuctric., The uncertainties for these data points are decidedly non-Gaussian and asymmetric.
+ Nevertheless. the reduced 42 feri givesz a rough. measure of the accuracy of our uncertaiutv estimates.andwe fiud," Nevertheless, the reduced $\chi^2$ term gives a rough measure of the accuracy of our uncertainty estimates,andwe find"
+is also detected if 5=0.34. but less clearly.,"is also detected if $\gamma=0.34$, but less clearly."
+" It should be noted that. i£ 4, = 0.300.24. Moe does not correlate with H. and hence with the environment. except for the fact that the density associated with the central galaxy is hieher."," It should be noted that, if $\gamma$ = 0.30–0.34, $N_{\rm IGC}$ does not correlate with $R$, and hence with the environment, except for the fact that the density associated with the central galaxy is higher."
+ Rather. the IGC population would be unilormlv spread all over Coma.," Rather, the IGC population would be uniformly spread all over Coma."
+ We will discuss this issue further in Section 4., We will discuss this issue further in Section 4.
+ For each value of 5. we have computed the average (Vice) in Coma using all but the central points.," For each value of $\gamma$, we have computed the average $\langle N_{\rm IGC}
+\rangle$ in Coma using all but the central points."
+ The results are listed in Table3. and have been plotted in Fig., The results are listed in Table\ref{meanIGC} and have been plotted in Fig.
+ 3. bv dashed lines., \ref{gamma} by dashed lines.
+ Dotted lines show the 1 opini. aitervals (σρωμις Is Lhe root mean square of the data point distribution)., Dotted lines show the 1 $\sigma_{\rm points}$ intervals $\sigma_{\rm points}$ is the root mean square of the data point distribution).
+ The estimated errors of the mean have also been computed using the usual formula and are the errors quoted in Table 3.., The estimated errors of the mean have also been computed using the usual formula and are the errors quoted in Table \ref{meanIGC}.
+" The extveme cases considered above lor the (ir) slope. ;=0.30 and 5,=0.39. characterize (wo distinct scenarios."," The extreme cases considered above for the $n(m)$ slope, $\gamma=0.30$ and $\gamma=0.39$, characterize two distinct scenarios."
+ We have seen (hat the latter produces negative Vice all over Coma.which indicates that ΠΠ) must be shallower.," We have seen that the latter produces negative $N_{\rm IGC}$ all over Coma,which indicates that $n(m)$ must be shallower."
+ Let us consider now in more detail (he possibility +=0.30., Let us consider now in more detail the possibility $\gamma=0.30$.
+ In this case. Fig.," In this case, Fig."
+ 3. shows that IGCs would exist ancl would be spread all over Coma., \ref{gamma} shows that IGCs would exist and would be spread all over Coma.
+ Furthermore. no relation with the distance to the galaxy. cluster center is apparent.," Furthermore, no relation with the distance to the galaxy cluster center is apparent."
+ ILowever. all the GC formation scenarios mentioned in 8??7 Πρίν (hat Nyee should follow the galaxy. cluster mass distribution.," However, all the GC formation scenarios mentioned in \ref{intro}
+ imply that $N_{\rm IGC}$ should follow the galaxy cluster mass distribution."
+ Asstuning this and in order to test our results. we will consider three different galaxy cluster simplistic mass distributions. p(t). and analyze the expected νου distribution for each model.," Assuming this and in order to test our results, we will consider three different galaxy cluster simplistic mass distributions, $\rho(R)$ , and analyze the expected $N_{\rm IGC}$ distribution for each model."
+" Makino(1994) has given for Coma a cut-off radius r,=3.5°. a core radius r.=0.24 (Le. rj/r.= 14.2) and. using //,=7248 km ! | (Freedmanetal.2001) for the IInbble parameter. a total mass 1.4xLOMAL.."," \citet{M94} has given for Coma a cut-off radius $r_t=3.5\arcdeg$, a core radius $r_c=0.24\arcdeg$ (i.e. $r_t/r_c=14.2$ ) and, using $H_0=72\pm8$ km $^{-1}$ $^{-1}$ \citep{F01} for the Hubble parameter, a total mass $1.4\times10^{15} M_{\odot}$."
+" This value for the IIubble parameter corresponds to a distance of 97.2 Mpc and tor,=5.9 Mpc and 7.=0.4 Alpe.", This value for the Hubble parameter corresponds to a distance of 97.2 Mpc and to $r_t=5.9$ Mpc and $r_c=0.4$ Mpc.
+" In this context. the three considered ewlaxy cluster mass distribution models correspond (o a Whine model. a homogeneous sphere of radius equal to the eut-olf radius of Coma (7,=2109 and a homogeneous sphere of radius equal to twice the eut-off radius of Coma."," In this context, the three considered galaxy cluster mass distribution models correspond to a King model, a homogeneous sphere of radius equal to the cut-off radius of Coma $r_t=210 \arcmin$ ) and a homogeneous sphere of radius equal to twice the cut-off radius of Coma."
+ All them are plotted in the upper panel of Figure 4.., All them are plotted in the upper panel of Figure \ref{massdis}. .
+ All the mass distributions are normalizedto a total galaxy cluster mass of 1.4xLOMAL. (Makino 1994)..," All the mass distributions are normalizedto a total galaxy cluster mass of $1.4\times10^{15}
+M_{\odot}$ \citep{M94}. ."
+ The vertical solid line represents the Coma, The vertical solid line represents the Coma
+the crystalline silicates observed in disks.,the crystalline silicates observed in disks.
+ An expanded model. which tracks specific forms of crystalline silicate. is required to establish in more detail the importance of this and other possible crystallisation mechanisms.," An expanded model, which tracks specific forms of crystalline silicate, is required to establish in more detail the importance of this and other possible crystallisation mechanisms."
+most of the error contour of 3EG J162148203.,most of the error contour of 3EG J1621+8203.
+ Additional partial coverage was also available from theROSAT High Resolution Imager (HRI) and the (Advanced Satellite for Cosmology and Astrophysics) Gas Imaging Spectrometer (GIS)., Additional partial coverage was also available from the High Resolution Imager (HRI) and the (Advanced Satellite for Cosmology and Astrophysics) Gas Imaging Spectrometer (GIS).
+ Historically. these fields have been studied in X-rays because they contained the radio-loud active galaxy NGC 6251 (e.g. Sambruna. Eracleous. Mushotzky 1999: Turner et al.," Historically, these fields have been studied in X-rays because they contained the radio-loud active galaxy NGC 6251 (e.g. Sambruna, Eracleous, Mushotzky 1999; Turner et al."
+ 1997: Mack. Kerp Klein 1997: Birkinshaw Worrall 1993). and the gamma-ray burst GRB 970815. which generated of-opportunity observations (Murakami et al.," 1997; Mack, Kerp Klein 1997; Birkinshaw Worrall 1993), and the gamma-ray burst GRB 970815, which generated target-of-opportunity observations (Murakami et al."
+ 1997)., 1997).
+ We have createdROSAT PSPC image of the field of 3EG J1621--8203 by co-adding exposure corrected sky maps of 14.7 ks of data taken during 1991 March 12-15. as shown in Fig.," We have created PSPC image of the field of 3EG J1621+8203 by co-adding exposure corrected sky maps of 14.7 ks of data taken during 1991 March 12-15, as shown in Fig."
+ I., 1.
+ TheROSAT HRI and GIS images of the partial field of 3EG J162148203 are shown in Fig., The HRI and GIS images of the partial field of 3EG J1621+8203 are shown in Fig.
+ 3., 3.
+ The figures also show the EGRET positions for the source. both as derived from the 3EG catalog. as well as that from the elliptical fits to the EGRET data.," The figures also show the EGRET positions for the source, both as derived from the 3EG catalog, as well as that from the elliptical fits to the EGRET data."
+ In the archivalROSAT and data overlapping the EGRET error circle ellipse of 3EG 11621-8203 we find several faint X-ray point sources., In the archival and data overlapping the EGRET error circle ellipse of 3EG J1621+8203 we find several faint X-ray point sources.
+ In addition. we have also searched theROSAT All Sky Survey (RASS) catalog for sources in the field of 3EG J16214+8203. particularly in the regions not covered by the pointedROSAT PSPC and HRI data.," In addition, we have also searched the All Sky Survey (RASS) catalog for sources in the field of 3EG J1621+8203, particularly in the regions not covered by the pointed PSPC and HRI data."
+ The source positions are marked in figures | and 3. and listed in Table 1.," The source positions are marked in figures 1 and 3, and listed in Table 1."
+ Count rates for the andROSAT sources were obtained following the method described in Gotthelf Kaspi (1998)., Count rates for the and sources were obtained following the method described in Gotthelf Kaspi (1998).
+ Photons were extracted using a 2’ radius aperture and the background contribution was estimated using a large annulus away from the souree., Photons were extracted using a $2^{\prime}$ radius aperture and the background contribution was estimated using a large annulus away from the source.
+ TheROSAT sources are listed in Table | along with their background subtracted count rates. detection significances. and hardness ratios.," The sources are listed in Table 1 along with their background subtracted count rates, detection significances, and hardness ratios."
+ Sources with count rates below 2.4«107 cts/s are not listed in the table., Sources with count rates below $2.4\times10^{-3}$ cts/s are not listed in the table.
+ The faintest source in Table 1 is a detection at the 3.00 level. with a background subtracted count rate of 0.0020+0.0006 cts/s in a 2.0’ diameter circular aperture.," The faintest source in Table 1 is a detection at the $3.0 \sigma$ level, with a background subtracted count rate of $0.0020\pm0.0006$ cts/s in a $^\prime$ diameter circular aperture."
+ For this analysis we report all sources down to the 3a detection level., For this analysis we report all sources down to the $3\sigma$ detection level.
+ The sources listed reached a minimum detectable intrinsic flux of 1.3«1077 erg em s! in the 0.1 — 2.4 keV band. assuming a power-law photon spectral index of 2.0. and no absorption (estimated Nj; is ~5«1079 cm. is expected to have a small effect).," The sources listed reached a minimum detectable intrinsic flux of $1.3\times
+10^{-14}$ erg $^{-2}$ $^{-1}$ in the 0.1 – 2.4 keV band, assuming a power-law photon spectral index of 2.0, and no absorption (estimated $N_H$ is $\sim 5\times 10^{20}$ $^{-2}$ is expected to have a small effect)."
+ We have searched for optical and radio counterparts to theROSAT and point sources shown in figures 1 and 3., We have searched for optical and radio counterparts to the and point sources shown in figures 1 and 3.
+ Most of the X-ray point sources have possible optical counterparts. as listed in Table 1. and some are coincident with faint radio sources.," Most of the X-ray point sources have possible optical counterparts, as listed in Table 1, and some are coincident with faint radio sources."
+ Individual sources denoted by their source numbers in Table | are described in the following section., Individual sources denoted by their source numbers in Table 1 are described in the following section.
+ I.6251:, 1.:
+ This is the bright FR I radio galaxy (Bicknell 1994: Urry Padovani 1995) at a redshift of 0.0234 (implied distance 91 Mpe for Hy=75 km s! ΜροΟ Ὀ., This is the bright FR I radio galaxy (Bicknell 1994; Urry Padovani 1995) at a redshift of 0.0234 (implied distance 91 Mpc for $H_0=75$ km $^{-1}$ $^{-1}$ ).
+ NGC 6251 is the parent galaxy of an exceptional radio jet from pe to Mpe scale which makes an angle of ~45° to our line of sight (Sudou Taniguchi 2000).," NGC 6251 is the parent galaxy of an exceptional radio jet from pc to Mpc scale which makes an angle of $\sim
+45^\circ$ to our line of sight (Sudou Taniguchi 2000)."
+ High dynamic range observations of NGC 6251 with the VLBI show the presence of a bright core. and an asymmetric jet. that implies relativistic beaming in this source (Jones et al.," High dynamic range observations of NGC 6251 with the VLBI show the presence of a bright core, and an asymmetric jet, that implies relativistic beaming in this source (Jones et al."
+ 1986)., 1986).
+ The VLBI observations are described further in $6., The VLBI observations are described further in 6.
+ NGC 6251 has been studied in X-rays in the past., NGC 6251 has been studied in X-rays in the past.
+ A detailed investigation of theROSAT PSPC data was carried out by Birkinshaw Worrall (1993). who found the ray emission to be consistent with a point source at the position," A detailed investigation of the PSPC data was carried out by Birkinshaw Worrall (1993), who found the X-ray emission to be consistent with a point source at the position"
+between the number of the neutron star binaries ancl (he back hole binaries.,between the number of the neutron star binaries and the back hole binaries.
+ Other models change the maximal masses of the black holes produced., Other models change the maximal masses of the black holes produced.
+ This is especially clear in (lie case of models G where the stellar winds are varied by a factor of (wo upwards (model G2) and downwards (model G1) We note that the shapes of these distributions do not depend on the sensitivity of a detector., This is especially clear in the case of models G where the stellar winds are varied by a factor of two upwards (model G2) and downwards (model G1) We note that the shapes of these distributions do not depend on the sensitivity of a detector.
+ In order to verily if the differences between the distributions are significant we (um {ο a simulation of a finite number of merger observations., In order to verify if the differences between the distributions are significant we turn to a simulation of a finite number of merger observations.
+ We assume that the (rue stellar evolution goes through one of the models of Table 1., We assume that the true stellar evolution goes through one of the models of Table 1.
+ We then simulate the observations of a given number of mergers (we use 20. LOO and 500 mergers) and for each such simulated observation we verify using the Kolmogorov 5nmirnov (XS) test if we can reject a hypothesis that the stellar evolution is described by model A. This allows to test (he sensitivity of the shape of the distribution of expected chirp mass observations to the underlving moclel parameters describing stellar evolution.," We then simulate the observations of a given number of mergers (we use 20, 100 and 500 mergers) and for each such simulated observation we verify using the Kolmogorov Smirnov (KS) test if we can reject a hypothesis that the stellar evolution is described by model A. This allows to test the sensitivity of the shape of the distribution of expected chirp mass observations to the underlying model parameters describing stellar evolution."
+ For each number of merger observations we repeal (his test 10000. (mes to obtain a distribution of INXS-test. probabilities and find the lowest probability that appeared in (the top one percentile of this distribution., For each number of merger observations we repeat this test 10000 times to obtain a distribution of KS-test probabilities and find the lowest probability that appeared in the top one percentile of this distribution.
+ We can now set a detection confidence level. sav al 10.? and compare each probability with this value: if it is higher we conclude that this particular model cannot be distinguished from model A with a eiven number of merger observations. while à smaller number means that this model can be distinguished. and (hat some constriants can be iposed on the particular parameter through which this model differes from model A. We present the results of the test in Figure 3..," We can now set a detection confidence level, say at $10^{-5}$ and compare each probability with this value: if it is higher we conclude that this particular model cannot be distinguished from model A with a given number of merger observations, while a smaller number means that this model can be distinguished, and that some constriants can be iposed on the particular parameter through which this model differes from model A. We present the results of the test in Figure \ref{res}. ."
+ Figure 3. presents a measure of sensitivitv of the expected observed distribution of chirp masses (o (he parameters describing stellar evolution., Figure \ref{res} presents a measure of sensitivity of the expected observed distribution of chirp masses to the parameters describing stellar evolution.
+ One can see [rom Figure 3. that even observations of a small number of mergers (open circles correspond to 20 mergers) vield highlv significant results for models El. G2 and ο. The reason for Chat can be is clear from Figure 2..," One can see from Figure \ref{res} that even observations of a small number of mergers (open circles correspond to 20 mergers) yield highly significant results for models E1, G2 and O. The reason for that can be is clear from Figure \ref{chidist}."
+ These are the models for which the maximal chirp mass in (he population is significantly lower than that for model A. Model G2 population (ie. with increased stellar winds) contains hardly aiv black holes., These are the models for which the maximal chirp mass in the population is significantly lower than that for model A. Model G2 population (i.e. with increased stellar winds) contains hardly any black holes.
+ In general we see that these observations are very sensitive to the value of maximum mass of stellar black holes in the population., In general we see that these observations are very sensitive to the value of maximum mass of stellar black holes in the population.
+ Model G1 (wilh decreased stellar winds) which allows for formation of black holes binaries with chirp nasses up to 16M. will stand ont wilh less (han a hundred merger observations.," Model G1 (with decreased stellar winds) which allows for formation of black holes binaries with chirp masses up to $16\,M_\odot$ will stand out with less than a hundred merger observations."
+ With a larger number of merger observation (stars in Figure 3)) correspond to 100 nereer detections) more parameters can be constrained., With a larger number of merger observation (stars in Figure \ref{res}) ) correspond to 100 merger detections) more parameters can be constrained.
+ Some constraints can be obtained or the value of common envelope elliciency oc4:À (models E: model E2 is similar to A)., Some constraints can be obtained for the value of common envelope efficiency $\alpha_{CE}\lambda$ (models E: model E2 is similar to A).
+ Other parameters describing mass transfer events like (he mass fraction accreted (models E). ancl the amount of angular momentum lost (models L) shall also be constrained.," Other parameters describing mass transfer events like the mass fraction accreted (models F), and the amount of angular momentum lost (models L) shall also be constrained."
+ Moreover also (he metallicity of (he progenitor stars may influence (he observed distributionat a significant, Moreover also the metallicity of the progenitor stars may influence the observed distributionat a significant
+described below. when the gamma-ray [Iuxes are obtained.,"described below, when the gamma-ray fluxes are obtained."
+ Among manv model spectra proposed [or the Galactic protons. we have chosen the power-law spectrum of index 2.0. the local interstellar spectrum. and a (ial spectrum [ου the Galactic ridge region.," Among many model spectra proposed for the Galactic protons, we have chosen the power-law spectrum of index 2.0, the local interstellar spectrum, and a trial spectrum for the Galactic ridge region."
+ We eive short description on them below., We give short description on them below.
+ The total gamuna-ray spectra for Ind2. LIS and TrialdGR are presented alter multiplving with E? for models A and B in Figs.6a. 6b and Ge.," The total gamma-ray spectra for Ind2, LIS and Trial4GR are presented after multiplying with $E_\gamma^2$ for models A and B in Figs.6a, 6b and 6c."
+" We note that model A gives harder eanmmnma-rayv spectra than model D for all 3 proton spectra for E,=>1 GeV (Fie.G): the asvinplolic power-law indices for model A/model B are 1.96/2.03 for Ind2. 2.65/2.71 for LIS and 2.47/2.53 for Trial4GR."," We note that model A gives harder gamma-ray spectra than model B for all 3 proton spectra for $E_\gamma>1$ GeV (Fig.6): the asymptotic power-law indices for model A/model B are $1.96/2.03$ for Ind2, $2.65/2.71$ for LIS and $2.47/2.53$ for Trial4GR."
+ It is to be noted that the model A 5gamma-ray spectra are harder than those of incident protons by index ο0.05., It is to be noted that the model A gamma-ray spectra are harder than those of incident protons by index $\sim 0.05$.
+ The EGRET count ancl exposure maps for (he observation period 1-4 have been downloaded [rom the EGRET archive (see footnote 3)., The EGRET count and exposure maps for the observation period 1-4 have been downloaded from the EGRET archive (see footnote 3).
+ All point sources listed in the EGRET 3rd Catalog (Ilartznan1999) are (hen removed using the point-spread function of EGRET for each enerev band and for the power-law index of each point source listed in the catalog., All point sources listed in the EGRET 3rd Catalog \citep{EGRET3rdCatalog} are then removed using the point-spread function of EGRET for each energy band and for the power-law index of each point source listed in the catalog.
+ The [Iux between LOO MeV and 10 GeV has been constrained (to that of each point source listed in the EGRET 2rd Catalog., The flux between 100 MeV and 10 GeV has been constrained to that of each point source listed in the EGRET 3rd Catalog.
+ The point-source-subtracted count map is then divided by the corresponding exposure map to make the intensity map., The point-source-subtracted count map is then divided by the corresponding exposure map to make the intensity map.
+ The intensity between the Galactic latitude 46.0 deg., The intensity between the Galactic latitude $\pm 6.0$ deg.
+ and Galactic longitude 430.0 deg has been summed ancl normalized to, and Galactic longitude $\pm 30.0$ deg has been summed and normalized to
+Note that the upper panelΑι.,Note that the upper panel.
+ This is so because the CLA search does not have the sensitivity to detect secondaries in (hat range., This is so because the CfA search does not have the sensitivity to detect secondaries in that range.
+" On (he other hand. the lower panel does include information on the bins below 100,."," On the other hand, the lower panel does include information on the bins below 100."
+. This panel presents the results of (he hieh-precision radial-velocity searches. and these searches could easily detect stars with secondaries in the range of. sav. 20100;.," This panel presents the results of the high-precision radial-velocity searches, and these searches could easily detect stars with secondaries in the range of, say, 20–100."
+. We assume that (hese binaries were not excluded from the various raclial-velocity searches at the first place. aud further assume Chat all. or at least most. findings of the various research eroups corresponding to this range of masses were already published.," We assume that these binaries were not excluded from the various radial-velocity searches at the first place, and further assume that all, or at least most, findings of the various research groups corresponding to this range of masses were already published."
+ If these two assumptions are correct. then the lower panel does represent the frequency of seconcdaries in the mass range of 20100.," If these two assumptions are correct, then the lower panel does represent the frequency of secondaries in the mass range of 20–100."
+"U,.. This panel shows that the frequency. of secondaries in (his range of masses is close to zero.", This panel shows that the frequency of secondaries in this range of masses is close to zero.
+" The present analvsis is not able to tell whether this ""brown-dwarf desert extends up to GO. 30 or LOO;."," The present analysis is not able to tell whether this “brown-dwarf desert” extends up to 60, 80 or 100."
+. The relative scaling of the planets and the stellar companions is not well known., The relative scaling of the planets and the stellar companions is not well known.
+ The spectroscopic binaries come [rom well defined samples = 577 stars for the lower-mass bin. and 312 stars for the higher-mass stars (Goldberg 2000).," The spectroscopic binaries come from well defined samples – 577 stars for the lower-mass bin, and 312 stars for the higher-mass stars (Goldberg 2000)."
+ ILowever. this is not the case lor the detected planets. specially because (he sample of published planets is not complete aud also because (he search samples of the different eroups are not well documented in the public domain.," However, this is not the case for the detected planets, specially because the sample of published planets is not complete and also because the search samples of the different groups are not well documented in the public domain."
+ We assumed. somewhatarbitrarily. that they come [rom a sample of 1000 stars. ancl scaled the stellar bins accordingly.," We assumed, somewhat, that they come from a sample of 1000 stars, and scaled the stellar bins accordingly."
+" We also rescaled the stellar bins to account [or the [act that their bins are larger. and (heir period range extends up to 3000 clavs. assuming, a [flat distribution in logP."," We also rescaled the stellar bins to account for the fact that their bins are larger, and their period range extends up to 3000 days, assuming a flat distribution in $\log P$."
+ Therefore the values of the stellar bins are our best estimate for the nunmber of binaries for 1000 stars within a mass range of 0.3 dex. aud up to a period of 1500 clavs.," Therefore the values of the stellar bins are our best estimate for the number of binaries for 1000 stars within a mass range of 0.3 dex, and up to a period of 1500 days."
+ Obviously the relative scaling of the (wo panels has a large uncertainty., Obviously the relative scaling of the two panels has a large uncertainty.
+ This scaling uncertainty is reflected in Cie error bars of the higher panel., This scaling uncertainty is reflected in the error bars of the higher panel.
+ Nevertheless we Chink (hat {he comparison is illuminating. as will be ciscussecl in the next section.," Nevertheless we think that the comparison is illuminating, as will be discussed in the next section."
+ The egrand picture that is emerginge from Figuree 3 stronglve. indicates that we have here two distinet populations., The grand picture that is emerging from Figure 3 strongly indicates that we have here two distinct populations.
+" The two populations are separated by à ""gap of about one decade of masses. in (he range between LO and LOOAZ;."," The two populations are separated by a “gap” of about one decade of masses, in the range between 10 and 100."
+. Such a gap was already noticed by many early studies (Dasri Marey 1997; Mavor. Queloz Udry 1998; Mavor. Uday Queloz 1998: Marev Butler 1998).," Such a gap was already noticed by many early studies (Basri Marcy 1997; Mayor, Queloz Udry 1998; Mayor, Udry Queloz 1998; Marcy Butler 1998)."
+ Those early papers binned (he mass distribution linearly., Those early papers binned the mass distribution linearly.
+" Here we follow our previous work (Mazeh et 11998) ancl use a logarithmic scale to study (he mass distribution. because of the large range of masses. 0.5.1000,.. involved."," Here we follow our previous work (Mazeh et 1998) and use a logarithmic scale to study the mass distribution, because of the large range of masses, 0.5–1000, involved."
+ The logarithic, The logarithmic
+MIPS data (Dickinsonetal..2003).,MIPS data \citep{dickinson03}.
+. In particular. they used deep intermediate-band imaging from the Subaru telescope to provide photometry with finer wavelength sampling than is possible from standard broad band data.," In particular, they used deep intermediate-band imaging from the Subaru telescope to provide photometry with finer wavelength sampling than is possible from standard broad band data."
+ This was done in part to enable more accurate photometric redshifts. but for our purposes here it provides a valuable means of tracing the detailed shape of the UV rest-frame spectrum and to look for the presence of a UV absorption We start with the GOODS-Herschel catalogue of 533 sources detected at 100 jum. and identify 325 sources with a SNR (flux over error) at 100 yam larger than 5.," This was done in part to enable more accurate photometric redshifts, but for our purposes here it provides a valuable means of tracing the detailed shape of the UV rest-frame spectrum and to look for the presence of a UV absorption We start with the $Herschel$ catalogue of 533 sources detected at 100 $\mu$ m, and identify 325 sources with a SNR (flux over error) at 100 $\mu$ m larger than 5."
+ We further restrict the sample to the 235 sources also detected at 160 pm with a SNR >3 and we cross-correlate them with the MUSYC catalogue of 84402 sources with photometric redshifts., We further restrict the sample to the 235 sources also detected at 160 $\mu$ m with a SNR $ > 3$ and we cross-correlate them with the MUSYC catalogue of 84402 sources with photometric redshifts.
+ The tolerance radius between IRAC and optical coordinates is taken to be 2 arcsec. it results in a sample of 219 sources cross-matched.," The tolerance radius between IRAC and optical coordinates is taken to be 2 arcsec, it results in a sample of 219 sources cross-matched."
+ Total fluxes in optical and bands are computed from aperture fluxes as described in(2010)., Total fluxes in optical and bands are computed from aperture fluxes as described in.
+ Before selecting sources according to their redshift we must choose the bands we will use: the redshift range will be selected to ensure a good sampling of the UV range., Before selecting sources according to their redshift we must choose the bands we will use: the redshift range will be selected to ensure a good sampling of the UV range.
+ All the optical broad bands of the MUSYC survey are considered (U. U38. B. V. R.Lz) and J.H.K bands are added when available.," All the optical broad bands of the MUSYC survey are considered (U, U38, B, V, R,I,z) and J,H,K bands are added when available."
+ We also add data from intermediate band filters whose So depth is fainter than 25 ABmag (table 2 of Cardamoneetal. (2010)))., We also add data from intermediate band filters whose $\sigma$ depth is fainter than 25 ABmag (table 2 of \citet{cardamone09}) ).
+ Practically we use [A-427.445.484.505.527.550.574.598.624.651.679 and 738.," Practically we use IA-427,445,484,505,527,550,574,598,624,651,679 and 738."
+ Figure | illustrates the rest-frame wavelengths corresponding to these filters as a function of z. In order to ensure a good photometric sampling around the UV bump (2175 A) we select galaxies with redshifts between 0.95 and 2.2., Figure 1 illustrates the rest-frame wavelengths corresponding to these filters as a function of z. In order to ensure a good photometric sampling around the UV bump (2175 $\AA$ ) we select galaxies with redshifts between 0.95 and 2.2.
+ In this redshift range we have more than 10 photometric bands available in the UV (1300-3000 A) and a good sampling around 2175 A (Fig.1))., In this redshift range we have more than 10 photometric bands available in the UV (1300-3000 $\AA$ ) and a good sampling around 2175 $\AA$ \ref{lambda}) ).
+ 76 galaxies fulfill the, 76 galaxies fulfill the
+required to predict the fluxes in the IRAC bands would be straightforward.,required to predict the fluxes in the IRAC bands would be straightforward.
+ Indeed. 41 of the 49 PG WDs already have SDSS photometry and 47 are detected in J band by 2MAÀSS.," Indeed, 41 of the 49 PG WDs already have SDSS photometry and 47 are detected in $J$ band by 2MASS."
+ OF course. an IR excessat 4.5jn could have causes other (han planets. in particular. circumstellar disks.," Of course, an IR excessat $4.5\,\mu$ m could have causes other than planets, in particular, circumstellar disks."
+ At present. these alternatives can be easily distinguished by observing ihe WDs at the otherSpi/zer TRAC bands: warm-planets would only show an excess ad 4.5jan. while cooldust emission would be detected in most IRAC bands.," At present, these alternatives can be easily distinguished by observing the WDs at the other IRAC bands: warm-planets would only show an excess at $4.5\mu$ m, while cool-dust emission would be detected in most IRAC bands."
+ In fact. all but one of the WDs with debris disks show excess emission in all 4 TRAC bands (von Hippel et al.," In fact, all but one of the WDs with debris disks show excess emission in all 4 IRAC bands (von Hippel et al."
+ 2007: Jura οἱ al., 2007; Jura et al.
+ 2007). ancl the majority of them also show excess in the IxX-band (Ixilie et al.," 2007), and the majority of them also show excess in the K-band (Kilic et al."
+ 2006)., 2006).
+ There is only one WD. G166-58. known to have a debris disk that shows up as an excess redward of 5j (Farihi et al.," There is only one WD, G166-58, known to have a debris disk that shows up as an excess redward of $5\mu$ m (Farihi et al."
+ 2007)., 2007).
+ After Spitzer's ervogen is exhausted. it will still be possible to cdillerentiate planets from debris disks since these disks would show excesses both in the 3.6jam and 4.550 bands.," After 's cryogen is exhausted, it will still be possible to differentiate planets from debris disks since these disks would show excesses both in the $3.6\mu$ m and $4.5\mu$ m bands."
+ However. the channel for discriminating between planets and cooler disks will disappear.," However, the channel for discriminating between planets and cooler disks will disappear."
+" The candidates found by ""Warm SpZzer [rom observations at 2.6yan and 4.5jan can still be distinguished [rom cool disks in some cases."," The candidates found by “Warm ” from observations at $3.6\,\mu$ m and $4.5\,\mu$ m can still be distinguished from cool disks in some cases."
+ The median distance of the PG WDs is GS pc. implying that planets Iving at projected separations r_2140AU would be separately resolved bySpizer.," The median distance of the PG WDs is 68 pc, implying that planets lying at projected separations $r_\perp\ga 140\,\au$ would be separately resolved by."
+ Such planets would be recognized bv their lack of optical counterpart (Gin SDSS [or the regions covered by that survev)., Such planets would be recognized by their lack of optical counterpart (in SDSS for the regions covered by that survey).
+" Because the planet orbit. expands bv ""nal/orig=M/Mwp. ie. à factor 57. these separations correspond to a4;Z20—25AU. which is larger than the orbits of Jovian planets in the Solar System."," Because the planet orbit expands by $a_{\rm final}/a_{\rm orig}=M/M_{\rm WD}$, i.e., a factor 5–7, these separations correspond to $a_{\rm orig}\ga 20-25\,\au$, which is larger than the orbits of Jovian planets in the Solar System."
+ Moreover. (here is a correlation between the ages and distances of the (argets.," Moreover, there is a correlation between the ages and distances of the targets."
+ Since the PG survey is hotter (vounger) objects can be detected at larger distances.," Since the PG survey is magnitude-limited, hotter (younger) objects can be detected at larger distances."
+ The median distances for PG WDs with ages <300 Myr ancl x1 Gyr are 106 pe and 89 pe. respectively.," The median distances for PG WDs with ages $\leq300$ Myr and $\leq1$ Gyr are 106 pc and 89 pc, respectively."
+ Hence. planets at orbital separations die235Hr—409AU would be resolved withSpilzer around the WDs vounger than 1 Gyr.," Hence, planets at orbital separations $a_{\rm orig}\ga 35-40\,\au$ would be resolved with around the WDs younger than 1 Gyr."
+" IXasperοἱal.(2007). [Failed to find anv planets of m>3Αμ, and a>30AU around 22 voung GINM stars. while Lafreniereetal.(2007). failed to find any ab a>40AU around 85 voung GINM stars."," \citet{kasper07} failed to find any planets of $m>2\,M_{\rm jup}$ and $a>30\,\au$ around 22 young GKM stars, while \citet{lafreniere07} failed to find any at $a>40\,\au$ around 85 young GKM stars."
+ However. the orbits of planets around massive stars could be larger.," However, the orbits of planets around massive stars could be larger."
+ The could in principle confirm most of the WD planets detected bySpilzer observations., The could in principle confirm most of the WD planets detected by observations.
+" TheJWST Near-Infrared Camera (NIRCAAD EWIIM at 4.5 yam is ~0.15""."," The Near-Infrared Camera (NIRCAM) FWHM at $4.5\,\mu$ m is $\sim0.15''$."
+ Since the planet/WD flix ratios at this wavelength must be greater than 2%.. resolution at 1. EWIIM is feasible.," Since the planet/WD flux ratios at this wavelength must be greater than , resolution at 1 FWHM is feasible."
+ At the median sample distance (D—68 pe). the minimum allowedorbit agua2710AU (or au;σωZ2 AU) corresponds to a," At the median sample distance $D=68$ pc), the minimum allowedorbit $a_{\rm final}>10\,\au$ (or $a_{\rm orig}\ga 2\,\au$ ) corresponds to a"
+"fu,=2N(H2)/(N(H)+2N(H3)), N(H$)/N(H5), and I(C*), where N(X) is the column density of species X over the whole cloud and H$ and H5 are ortho- and para-H2 respectively.","$f_{{\rm H}_2}=2N({\rm H}_2)/(N({\rm H})+2N({\rm H}_{2}))$ $N({\rm H}_2^{\rm o})/N({\rm H}_2^{\rm p})$ , and $I({\rm{C}^+})$, where $N({\rm X})$ is the column density of species ${\rm X}$ over the whole cloud and ${\rm H}_2^{\rm o}$ and ${\rm H}_2^{\rm p}$ are ortho- and ${\rm{H}_2}$ respectively."
+" From thesequantities, we define a x? error function by: where X9> is the quantity derived from observations, XP? the same quantity from a model, and σους the observational uncertainty (see Table 1))?.."," From thesequantities, we define a $\chi^{2}$ error function by: where $X_i^{\rm obs}$ is the quantity derived from observations, $X_i^{\rm mod}$ the same quantity from a model, and $\sigma_{\rm obs}$ the observational uncertainty (see Table \ref{tab:XmodXobs}) )."
+ We computed a grid of models for gas densities ranging from 30 to 200 cm-? and UV fields G from 0.25 to 1.0., We computed a grid of models for gas densities ranging from 30 to 200 $cm^{-3}$ and UV fields $G$ from 0.25 to 1.0.
+ Figure 2 shows the resulting x? iso-values., Figure \ref{chi2_o} shows the resulting $\chi^{2}$ iso-values.
+" Contours are smooth, and limit a well defined minimum where y?<1, i.e. where model results are,mean,, closer to observations than observational uncertainties."," Contours are smooth, and limit a well defined minimum where $\chi^{2}<1$, i.e. where model results are, closer to observations than observational uncertainties."
+" One can see that the best compromise is reached for a rather low radiation field (G=0.4, which is close to the standard Habing field) and a mean total proton density of ng=80cm-?."," One can see that the best compromise is reached for a rather low radiation field $G=0.4$, which is close to the standard Habing field) and a mean total proton density of $n_{\rm H}=80\,{\rm{{cm}}}^{-3}$."
+ In the following we will refer to those values as our reference model., In the following we will refer to those values as our reference model.
+ The reference model results are compared to the observables in Table 1.., The reference model results are compared to the observables in Table \ref{tab:XmodXobs}.
+" Abundance and temperature profiles are illustrated on Fig. 3,,"," Abundance and temperature profiles are illustrated on Fig. \ref{CoupeHH2},"
+ 4 and 5cloud.., \ref{CoupeCOCCH} and \ref{Temp}.
+" The HI to Πο transition occurs close to the cloud edge at an extinction of A,~ 107%.", The ${\rm H~{\sc I}}$ to ${\rm H}_2$ transition occurs close to the cloud edge at an extinction of $A_{v}\sim10^{-3}$ .
+ Most of the carbon is in the form of Ct at all depths., Most of the carbon is in the form of ${\rm C}^+$ at all depths.
+ The fraction of atomic carbon is constant throughout the cloud., The fraction of atomic carbon is constant throughout the cloud.
+ T'he density profiles of CO and CH molecules follow that of Ho., The density profiles of ${\rm CO}$ and ${\rm CH}$ molecules follow that of ${\rm H}_2$ .
+" Temperature varies from 60 to 80 K, and is lowest at the cloud edge."," Temperature varies from 60 to 80 K, and is lowest at the cloud edge."
+ The temperature bumps visible on Fig., The temperature bumps visible on Fig.
+ 5 are due to heating by Πο formation., \ref{Temp} are due to heating by ${\rm H}_2$ formation.
+ The gas heating rate is 1.8x10-?4ergcm-?s-! at the cloud center and comes mainly from photo- effect on grains.," The gas heating rate is $1.8\times 10^{-24}\,{\rm{{erg\, cm}}}^{-3}\,{\rm s}^{-1}$ at the cloud center and comes mainly from photo-electric effect on grains."
+" In the outer layers where the temperature peaks, photo-electric heating reaches 3.2x10-?ergcm-?s-! and there is a lower but significant heating due to Hy formation (2.2x107?ergcm?s- 1)."," In the outer layers where the temperature peaks, photo-electric heating reaches $3.2\times 10^{-24}\,{\rm{{erg\, cm}}}^{-3}\,{\rm s}^{-1}$ and there is a lower but significant heating due to ${\rm H}_{2}$ formation $2.2\times 10^{-24}\,{\rm{{erg\, cm}}}^{-3}\,{\rm s}^{-1}$ )."
+" Cooling is dominated by C*,withHa contributing up to around 10-? A,."," Cooling is dominated by ${\rm C}^{+}$ ,with${\rm H}_{2}$ contributing up to around $10^{-2}\, A_{v}$ ."
+" This leads to an integrated Ct emissivity of 2x1079ergcm""?s!sr-! in good agreement with theISO observation."," This leads to an integrated ${\rm{{C}^+}}$ emissivity of $2\times 10^{-6}\,{\rm{{erg\, cm}}}^{-2}\,{\rm s}^{-1}\,{\rm{{sr}}}^{-1}$ in good agreement with theISO observation."
+" Inthe next sections, we discuss each of the observables and their dependence on the two model parameters ng and G."," Inthe next sections, we discuss each of the observables and their dependence on the two model parameters $_{H}$ and $G$ ."
+Cepheids are powerful standard caudles because thev follow a PeriodLunuinositv (PL) relation.,Cepheids are powerful standard candles because they follow a Period–Luminosity (PL) relation.
+ This relation has been determined using Cepheids im the Large Magallanic Cloud (LMC) in optical and near infrared bauds (?).., This relation has been determined using Cepheids in the Large Magallanic Cloud (LMC) in optical and near infrared bands \citep{Laney1994}.
+ LAIC Cepheids also provide iusight iuto stellar astroplivsics as they have lower metallicity relative to Galactic Cepheids. which has an effect on the pulsation of Cepheids.," LMC Cepheids also provide insight into stellar astrophysics as they have lower metallicity relative to Galactic Cepheids, which has an effect on the pulsation of Cepheids."
+ Receutlv. infrared PL relations have been derived using Spitzer observations from the SAGE program (2) bv ? aud bv ον," Recently, infrared PL relations have been derived using Spitzer observations from the SAGE program \citep{Meixner2006} by \cite{Ngeow2008} and by \cite{Freedman2008}."
+ These infrared PL relatious are important for extragalactic studies. aud this will be even more so when the James Web Space Telescope begins operation.," These infrared PL relations are important for extragalactic studies, and this will be even more so when the James Web Space Telescope begins operation."
+ The infrared PL relations are powerful tools because imoetallicitv does not contribute senificautlv (7) and because the pulsation amplitude decreases in the infrared., The infrared PL relations are powerful tools because metallicity does not contribute significantly \citep{Freedman2008} and because the pulsation amplitude decreases in the infrared.
+ However. IRAS observations have found infrared excesses in Calactic Cepheids (?)..," However, IRAS observations have found infrared excesses in Galactic Cepheids \citep{Deasy1988}."
+ Tuterterometric observations have also detected the existence of circumstellar cuvelopes around a number of Calactic Cepheids (277) in the Ivbaud.," Interferometric observations have also detected the existence of circumstellar envelopes around a number of Galactic Cepheids \citep{Kervella2006, Merand2006, Merand2007} in the K–band."
+ The observations of infrared excess duuply that there may be an additional uncertainty du iufrared PL relatious. aud the excesses may. play a role iu LALIC Cepheids.," The observations of infrared excess imply that there may be an additional uncertainty in infrared PL relations, and the excesses may play a role in LMC Cepheids."
+ Oue proposed mncchanisin for generating circuiistellar shells aud infrared excesses about Cauactic Cepheids is a stellar wind similar to those generated from other evolved pulsating stars., One proposed mechanism for generating circumstellar shells and infrared excesses about Galactic Cepheids is a stellar wind similar to those generated from other evolved pulsating stars.
+ Mass loss is believed to geuerate shells about asviuptotie giaut branch stars [where the massloss rates are related to pulsation aud dust condensation in the atmospheres. iu both the Milkv Way aud the LMC (?j]].," Mass loss is believed to generate shells about asymptotic giant branch stars [where the mass--loss rates are related to pulsation and dust condensation in the atmospheres, in both the Milky Way and the LMC \citep{Mattsson2008}] ]."
+ One would not expect dust to form in the atinospheres of Cepheids because the temperatures are eyeater than 1500 A. aud dustdriving is not a plausible nechanisin to generate a stellar wind.," One would not expect dust to form in the atmospheres of Cepheids because the temperatures are greater than $1500$ $K$, and dust–driving is not a plausible mechanism to generate a stellar wind."
+ However. it has )oen argued that Calactic Cepheids use pulsation aud shocks generated by pulsation to eject mass (7).," However, it has been argued that Galactic Cepheids use pulsation and shocks generated by pulsation to eject mass \citep{Willson1989}."
+" 7? developed an analytic model to study the affect of oilsatiou aud shocks. iu combination with radiativeine driviug. on massloss rates in Calactic Cepheids. xedietiug vates of the order 10j|"" to 10TAL./yr."," \cite{Neilson2008} developed an analytic model to study the affect of pulsation and shocks, in combination with radiative--line driving, on mass–loss rates in Galactic Cepheids, predicting rates of the order $10^{-10}$ to $10^{-7} M_\odot /yr$."
+ It is argued that at some large distance from the surface of the Cepheid. the wind cools enough that a small raction of the eas condenses iuto a dust shell that xoduces an infrared excess.," It is argued that at some large distance from the surface of the Cepheid, the wind cools enough that a small fraction of the gas condenses into a dust shell that produces an infrared excess."
+ The predicted iufrared excesses from the theoretical model are consistent witli roth interferometric observations aud TRAS observatious, The predicted infrared excesses from the theoretical model are consistent with both interferometric observations and IRAS observations
+The most massive Galactic globular cluster w Cen. is observed to have unique physical properties. such as a very Hattened shape for a globular cluster (c.g. Alevlan 1987). broad metallicity: distribution. (e.g. Freeman. Rodgers 1975: Norris et al.,"The most massive Galactic globular cluster $\omega$ Cen is observed to have unique physical properties, such as a very flattened shape for a globular cluster (e.g., Meylan 1987), broad metallicity distribution (e.g., Freeman Rodgers 1975; Norris et al."
+ 1906). strong variations of nearly all clement abundances among its stars (e.g. Norris Da Costa 1995: Smith et al.," 1996), strong variations of nearly all element abundances among its stars (e.g., Norris Da Costa 1995; Smith et al."
+ 2000). kinematical difference between its metal-rich. ancl metal-poor stellar populations (e.g.. Norris et al.," 2000), kinematical difference between its metal-rich and metal-poor stellar populations (e.g., Norris et al."
+ 1997). multiple stellar. populations with cillerent spatial clistributions (e.g.. Pancino et al.," 1997), multiple stellar populations with different spatial distributions (e.g., Pancino et al."
+ 2000: Ferraro et al., 2000; Ferraro et al.
+ 2002). star formation history extending over a few Gyr (Lee et al.," 2002), star formation history extending over a few Gyr (Lee et al."
+ 1999: Smith ct al., 1999; Smith et al.
+ 2000). and. its very bound retrograde orbit with respect to the Galactic rotation ( Dinescuet al.," 2000), and its very bound retrograde orbit with respect to the Galactic rotation (Dinescu et al."
+ 1999)., 1999).
+ Phese unique eharacteristies have been considered to suggest that there are remarkable cdillerences in star formation histories. chemical enrichment processes. and structure formation between w Cen and other Galactic normal elobular clusters (e.g. Hilker Richtler 2000. 2002) The observed extraordinary nature of w Cen has attracted much. attention from theoretical and numerical works on chemical and dynamical evolution of w Cen (c.g. Icke Aleafnno 1988: Carraro Lia 2000: Cinedin ct al.," These unique characteristics have been considered to suggest that there are remarkable differences in star formation histories, chemical enrichment processes, and structure formation between $\omega$ Cen and other Galactic normal globular clusters (e.g., Hilker Richtler 2000, 2002) The observed extraordinary nature of $\omega$ Cen has attracted much attention from theoretical and numerical works on chemical and dynamical evolution of $\omega$ Cen (e.g., Icke Alcaínno 1988; Carraro Lia 2000; Gnedin et al."
+ 2002: Zhao 2002)., 2002; Zhao 2002).
+ One of the most extensively discussed scenario for w Cen formation is that w Cen is the surviving nucleus of an ancient nucleated chwarl galaxy with its outer stellar envelope almost entirely removed. by tidal stripping of the Galaxy (Zinnecker et al., One of the most extensively discussed scenario for $\omega$ Cen formation is that $\omega$ Cen is the surviving nucleus of an ancient nucleated dwarf galaxy with its outer stellar envelope almost entirely removed by tidal stripping of the Galaxy (Zinnecker et al.
+ 1988: Freeman 1993)., 1988; Freeman 1993).
+" -""he observed. atypical bimodal or multi-modcel metallicity distribution (e.g.. Norris et al."," The observed atypical bimodal or multi-model metallicity distribution (e.g., Norris et al."
+ 1996) and the metal-rich stellar. population 2 4 Gyr vounger than the metal-poor (Lee et al., 1996) and the metal-rich stellar population 2 $-$ 4 Gyr younger than the metal-poor (Lee et al.
+ 1999: LHülker Richtler 2000: Llughes Wallerstein 2000) have been suggested to support. this scenario., 1999; Hilker Richtler 2000; Hughes Wallerstein 2000) have been suggested to support this scenario.
+ However. of becausethe lack of extensive numerical studies on dynamical evolution ofquetealed dwarf galaxies interacting/mereging with the Galaxy. it remains unclear when and how an ancient. nucleated. dwarl galaxy. loseonly its stellar envelope without totally destroving its nucleus in its dynamical interaction with the Galaxy.," However, because of the lack of extensive numerical studies on dynamical evolution of dwarf galaxies interacting/merging with the Galaxy, it remains unclear when and how an ancient nucleated dwarf galaxy lose its stellar envelope without totally destroying its nucleus in its dynamical interaction with the Galaxy."
+ q.sPhe purpose of⋅ this. paper is. to demonstrate that w Cen can be formed from an ancient nucleated dwarl, The purpose of this paper is to demonstrate that $\omega$ Cen can be formed from an ancient nucleated dwarf
+One of the primary goals of future galaxy recshilt surveys is o determine the physics behind the accelerating expansion of the Universe by making an accurate measurement. of he growth rate. f. of large scale structure (2)..,"One of the primary goals of future galaxy redshift surveys is to determine the physics behind the accelerating expansion of the Universe by making an accurate measurement of the growth rate, $f$, of large scale structure \citep{2009ExA....23...39C}."
+ Measuring he growth rate with an error of less than 104 is one of he main science goals of Euclid. as this will allow us to clistineuish mocified gravity from dark energy models.," Measuring the growth rate with an error of less than $10\%$ is one of the main science goals of Euclid, as this will allow us to distinguish modified gravity from dark energy models."
+ With an independent. measurement of the expansion history. the oedieted growth rate for a dark energy model would agree with the observed. value of f if general relativity holds.," With an independent measurement of the expansion history, the predicted growth rate for a dark energy model would agree with the observed value of $f$ if general relativity holds."
+ We use simulations of three quintessence dark energy models which have cilferent expansion histories. linear erowth rates ancl power spectra compared to ACDAL," We use simulations of three quintessence dark energy models which have different expansion histories, linear growth rates and power spectra compared to $\Lambda$ CDM."
+" In à previous paper. ?.. we carried out the first fully consistent N-bocly simulations of quintessence dark energy.Nl taking into account different expansion histories. linear. theory power spectra ancl best. fitting. cosmological parameters Qu. 3, and fü). for cach model."," In a previous paper, \citet{2010MNRAS.401.2181J}, we carried out the first fully consistent N-body simulations of quintessence dark energy, taking into account different expansion histories, linear theory power spectra and best fitting cosmological parameters $\Omega_{\rm m}$, $\Omega_{\rm b}$ and $H_0$, for each model."
+ ln this paper we examine the redshift space distortions in the SUGRA. CNR and 2EXP quintessence models.," In this paper we examine the redshift space distortions in the SUGRA, CNR and 2EXP quintessence models."
+ These models are representative of a broader class of quintessence models which have different erowth histories and dark energy. densities at. carly times compared to ACDAL, These models are representative of a broader class of quintessence models which have different growth histories and dark energy densities at early times compared to $\Lambda$ CDM.
+ In particular the SUGRA model has a linear growth rate that dillers from CDM by ~20% at >—5 and the CNR. model has high levels of dark energy at carly times. pp0.03 at 2~200.," In particular the SUGRA model has a linear growth rate that differs from $\Lambda$ CDM by $\sim 20\%$ at $z=5$ and the CNR model has high levels of dark energy at early times, $\Omega_{\rm \tiny{DE}} \sim 0.03$ at $z \sim 200$."
+ The 2EXP moclel has à similar expansion history to ACDAL at low recshilts. οκ5. despite having a dynamical equation of state for the dark energy component.," The 2EXP model has a similar expansion history to $\Lambda$ CDM at low redshifts, $z<5$, despite having a dynamical equation of state for the dark energy component."
+ For more details on each of the dark enereyv models sec ?.., For more details on each of the dark energy models see \citet{2010MNRAS.401.2181J}.
+ Redshift space distortions observed in galaxy surveys are the result. of peculiar. velocities which are coherent on large scales. leacing to a boost in the observed redshift space power spectrum compared to the real space power spectrum. (?)..," Redshift space distortions observed in galaxy surveys are the result of peculiar velocities which are coherent on large scales, leading to a boost in the observed redshift space power spectrum compared to the real space power spectrum \citep{Kaiser:1987qv}."
+ On small scales these peculiar velocities are incoherent ancl give rise to a damping in the ratio of the redshift to real space power spectrum., On small scales these peculiar velocities are incoherent and give rise to a damping in the ratio of the redshift to real space power spectrum.
+ The Ixaiser formula is a prediction of the boost in this ratio on very laree scales. where the growth is assumed. to be linear. and can be expressed as a function of the linear growth rate and bias. neglecting all non-linear contributions.," The Kaiser formula is a prediction of the boost in this ratio on very large scales, where the growth is assumed to be linear, and can be expressed as a function of the linear growth rate and bias, neglecting all non-linear contributions."
+ In previous work. using N-body simulations in a periodic cube of 3005. Mpe on a side. 7. found that the measured value of 3=fb. where b is the linear bias. deviates from the Kaiser formula on wavelengths of 505.+ Alpe or more as a result of these non-linearities.," In previous work, using N-body simulations in a periodic cube of $300 h ^{-1}$ Mpc on a side, \citet{Cole:1993kh} found that the measured value of $\beta =f/b$, where $b$ is the linear bias, deviates from the Kaiser formula on wavelengths of $50 h^{-1}$ Mpc or more as a result of these non-linearities."
+ 7. extended this analvsis to slightly. larger scales using the Zeldovich approximation combined with a dispersion. model. where non-linear velocities are treated as random. perturbations lc the linear theory velocity., \citet{Hatton:1997xs} extended this analysis to slightly larger scales using the Zel'dovich approximation combined with a dispersion model where non-linear velocities are treated as random perturbations to the linear theory velocity.
+ These previous studies. co not provide an accurate cleseription of the non-lincaritics in the velocity field for two reasons., These previous studies do not provide an accurate description of the non-linearities in the velocity field for two reasons.
+ Firstly. the Zel'dovich approximation does not model the velocities correctly. as it only treats part of the bulk motions.," Firstly, the Zel'dovich approximation does not model the velocities correctly, as it only treats part of the bulk motions."
+ Secondly. in a computational box of length 300A Mpe. the power which determines the bulk [Lows has not converged.," Secondly, in a computational box of length $300 h ^{-1}$ Mpc, the power which determines the bulk flows has not converged."
+ In this work we use a large Computational box of side 15007. Mpe. which allows us to measure redshift space distortions on Large scales to far greater accuracy than in previous work.," In this work we use a large computational box of side $1500 h^{-1}$ Mpc, which allows us to measure redshift space distortions on large scales to far greater accuracy than in previous work."
+ In this paper we find that the ratio of the monopole of the redshift space power spectrum. to the real space power spectrum agrees with the linear theory Ixaiser formula only on extremely large scales &«0.035 in both ACDAL and the quintessence dark energy models., In this paper we find that the ratio of the monopole of the redshift space power spectrum to the real space power spectrum agrees with the linear theory Kaiser formula only on extremely large scales $k<0.03 h$ $^{-1}$ in both $\Lambda$ CDM and the quintessence dark energy models.
+ We still, We still
+In Fig.,In Fig.
+" 5. for a fixed n,=1and kp,=0.1. the frequeney is calculated in terms of j;."," 5, for a fixed $\eta_b=1$and $k_y \rho_s=0.1$, the frequency is calculated in terms of $\eta_i$."
+ The two negative real solutions Fig., The two negative real solutions Fig.
+" 2) merge for i;z1.78 vielding a pair of complex conjugate solutions with a very weakly decreasing real part (Q,=——L.67 at i;=1.78 and QO,=L7la yi=10).", 2) merge for $\eta_i\geq 1.78$ yielding a pair of complex conjugate solutions with a very weakly decreasing real part $\Omega_r=-1.67$ at $\eta_i=1.78$ and $\Omega_r=-1.71$ at $\eta_i=10$ ).
+ Note that in all these cases the used values of the ratios pp—LofLg. pm=LofLyg in fact imply a wide range of possible values for ωνby.Lo.," Note that in all these cases the used values of the ratios $\eta_i\equiv L_n/L_{\sss T}$, $\eta_b\equiv L_n/L_{\sss B}$ in fact imply a wide range of possible values for $L_n, \, L_{\sss T},
+L_{\sss B}$."
+ Since ων=Όρων~l/l. this also implies a wide range of possible frequencies.," Since $\omega_r= \Omega_r \omega_{*e}\sim 1/L_n$, this also implies a wide range of possible frequencies."
+ The electrostatic instability cliseussecd here. implies. an electric field varving in time and space. and having very different components in. the parallel. and. perpendicular directions.," The electrostatic instability discussed here implies an electric field varying in time and space, and having very different components in the parallel and perpendicular directions."
+ Such an electric. field. can accelerate: plasma xwticles in both perpencdieular ancl parallel directions with respect to the magnetic field. vector., Such an electric field can accelerate plasma particles in both perpendicular and parallel directions with respect to the magnetic field vector.
+ An acceleration always exists in the presence of electrostatic: perturbations. vet vpicallv it is sporadic and acts mainly on the small amount of particles from the far tail in the distribution function.," An acceleration always exists in the presence of electrostatic perturbations, yet typically it is sporadic and acts mainly on the small amount of particles from the far tail in the distribution function."
+" The critical value of such an electric field. above which the sulk electron. runaway effect. takes place. in a Lully ionized slasmais (Dreicer1959). E,οἱ(dieuA)."," The critical value of such an electric field, above which the bulk electron runaway effect takes place, in a fully ionized plasmais \citep{dr} $E_d=e
+L_{ei}/(4 \pi \varepsilon_0 \lambda_d^2)$."
+" llere. ολ) is the Coulomb logarithm. Ay=AqAÀ(AS,| is the plasma Debye radius. Aqj9erjfiepj: Crpep) are. respectively. the thermal velocity ancl the plasma requeney of the j species. and b=nοπου|11] is the impact. parameter for electron-ion collisions.αν."," Here, $L_{ei}=\log(\lambda_d/b)$ is the Coulomb logarithm, $\lambda_d=\lambda_{de} \lambda_{di}/(\lambda_{de}^2+
+\lambda_{di}^2)^{1/2}$ is the plasma Debye radius, $\lambda_{dj}=\vtj/\omega_{pj}$, $\vtj, \omega_{pj}$ are, respectively, the thermal velocity and the plasma frequency of the $j$ species, and $b=[e^2/[12 \pi \varepsilon_0 (T_e+ T_i)]$ is the impact parameter for electron-ion collisions."
+" For the xwanmeters used in the previous text we have £,;=19. Ay=0005 m. and the Dreicer field is 0.11 V/m. Assuming the xwallel wave-Iength of about 1006500) km. the amplitude of he electrostatic potential ó necessary to achieve the Dreicer value is about 1.8(9). WY."," For the parameters used in the previous text we have $L_{ei}=19$, $\lambda_d=0.0005\;$ m, and the Dreicer field is $0.11\;$ V/m. Assuming the parallel wave-length of about $100 (500)\;$ km, the amplitude of the electrostatic potential $\phi$ necessary to achieve the Dreicer value is about $1.8 (9)\;$ KV."
+" However. in the perpendicular direction this same potential gives the electric. field Ao hat is around £=5.7(29) KV/m. For the parallel wave-eneth of about 1000. km we would have ó=1S KY and consequently 4;=57 IxXV/m. These estimates are for the number density: ag—I0""Li 3 ""7."," However, in the perpendicular direction this same potential gives the electric field $k_y \phi$ that is around $E_\bot=5.7 (29)\;$ KV/m. For the parallel wave-length of about $1000\;$ km we would have $\phi=18\;$ KV and consequently $E_\bot=57\;$ KV/m. These estimates are for the number density $n_0=10^{16}\;$ $^{-3}$ ."
+ Taking the number density. one order of magnitude larger ny=10 m vields £j=1 Vim. The electric Ποιά corresponding to this value in the parallel direction would. (for the three given parallel wavelengths) in the perpendicular direction have the magnitude of 54. 270. and 540 KV/m. respectively.," Taking the number density one order of magnitude larger $n_0=10^{17}\;$ $^{-3}$ yields $E_d=1\;$ V/m. The electric field corresponding to this value in the parallel direction would (for the three given parallel wavelengths) in the perpendicular direction have the magnitude of $54$ , $270$, and $540\;$ kV/m, respectively."
+ The time needed for the perturbations to achieve such values can be estimated. from. the previously calculated erowth-rates., The time needed for the perturbations to achieve such values can be estimated from the previously calculated growth-rates.
+ Taking as an example Fig., Taking as an example Fig.
+ 3. the maximum erowth-rate is 5ωνc1.," 3, the maximum growth-rate is $\gamma/\omega_{*e}\simeq
+1$."
+" Taking L,=10 m. 3;=1;Ty=10"" We Buy=3-10? T. for A,=2 m we have we=9 Hz."," Taking $L_n=10^3\;$ m, $T_i=T_e=T_0=10^6\;$ K, $B_0=3\cdot 10^{-2}\;$ T, for $\lambda_y=2\;$ m we have $\omega_{*e}=9\;$ Hz."
+" Assuming some small starting value of the electrostatic potential V. the growth time till it gets some value o, is fy2 log(ó,/W)/5."," Assuming some small starting value of the electrostatic potential $\Psi$ , the growth time till it gets some value $\phi_1$ is $t_g\simeq
+\log(\phi_1/\Psi)/\gamma$ ."
+" Taking ο(κ)=0.01 this vields W=0.86. V. Hence. the value oO,=9 KY discussed above is achieved. within /,,271 s. Observethat for L,,=10! m. this growth timebecomes 10 seconds."," Taking $e \Psi/(\kappa T_i)=0.01$ this yields $\Psi=0.86\;$ V. Hence, the value $\phi_1=9\;$ KV discussed above is achieved within $t_g\simeq 1\;$ s. Observethat for $L_n=10^4\;$ m, this growth timebecomes $10$ seconds."
+ We conclude that the presented. instability can vield the extremely large values of the cleetric field: reported, We conclude that the presented instability can yield the extremely large values of the electric field reported
+"Ajj, is given by the familiar value A,;,=[807z178: for flat ACDM models. A,;,. can be approximated by Ajj,22(1837+82:6—3997)/O,,(2) with =Q,,(2)—1 (Bryan Norman 1993).","$\Delta_{vir}$ is given by the familiar value $\Delta_{vir} =
+18\pi^2\approx 178$; for flat $\LCDM$ models, $\Delta_{vir}$ can be approximated by $\Delta_{vir} \approx (18\pi^2+82x-39x^2)/\Omega_m(z)$ with $x=\Omega_m(z)-1$ (Bryan Norman 1998)."
+ Jenkins et al., Jenkins et al.
+" has specifically stated that their formula gives better fits to dn/dAl with A,;,=118 regardless of the cosmological model. so we followed this instruction."," has specifically stated that their formula gives better fits to $dn/dM$ with $\Delta_{vir}=178$ regardless of the cosmological model, so we followed this instruction."
+" The lensing cross section ty, in eq. (", The lensing cross section $\sigma_{\rm lens}$ in eq. (
+1) depends on the mass profile of the lenses.,1) depends on the mass profile of the lenses.
+ Since our interest is (o predict the lensing probability over a wide range of image separations. we consider both galaxy-size halos (hat are responsible for lens svstenis of a lew arcseconds. ancl c]Iuster-size halos for larger 0.," Since our interest is to predict the lensing probability over a wide range of image separations, we consider both galaxy-size halos that are responsible for lens systems of a few arcseconds, and cluster-size halos for larger $\theta$."
+ In the former case. we approximate (he galaxy mass distribution as an SIS with a density prolile p(r)=02/(2zGr?). where σι is the 1-d velocity dispersion.," In the former case, we approximate the galaxy mass distribution as an SIS with a density profile $\rho(r)=\sigma_v^2/(2\pi G r^2)$, where $\sigma_v$ is the 1-d velocity dispersion."
+" An SIS lens produces an image separation of 20,5. where 0,=4z(0,./0)?Dj./D, is the Einstein radius. and the cross section is ej,=z(05D;)?l6z(o./0)(IDDi/DL)? (e.g.. Schneider οἱ al."," An SIS lens produces an image separation of $2\theta_E$, where $\theta_E=4\pi (\sigma_v/c)^2
+D_{ls}/D_s$ is the Einstein radius, and the cross section is $\sigma_{\rm lens}=\pi (\theta_E D_l)^2 = 16\pi^3 (\sigma_v/c)^4
+(D_lD_{ls}/D_s)^2$ (e.g., Schneider et al."
+" 1992). where D,.Dj. and Dj; ave the angular οποίοι distances to the source. to the lens. and between the lens and the source. respectively."," 1992), where $D_s, D_l$, and $D_{ls}$ are the angular diameter distances to the source, to the lens, and between the lens and the source, respectively."
+" Using σι=(x67M?N4,p/G6). we can then relate ej4,, to the halo mass which is needed Lor eq. ("," Using $\sigma_v=(\pi G^3 M^2 \Delta_{vir}\bar\rho/6)^{1/6}$, we can then relate $\sigma_{\rm lens}$ to the halo mass which is needed for eq. ("
+1).,1).
+ For lensing by cluster-size halos. we approximate the cluster mass distribution bv the density profile determined from halos in numerical simulations (Navarro et al.," For lensing by cluster-size halos, we approximate the cluster mass distribution by the density profile determined from halos in numerical simulations (Navarro et al."
+" 1997): utr/r,). where utr)=1/Lr(1+.r)?]."," 1997): $\rho(r)=\bar\rho\,\bar\delta\,
+u(r/r_s)$ , where $u(x)=1/[x(1+x)^2]$."
+ This profile is shallower than the SIS for the inner paris of a virialized halo but is steeper at large radii., This profile is shallower than the SIS for the inner parts of a virialized halo but is steeper at large radii.
+" Other than (he virial mass. (his profile is described by a concentration parameter e=At4,/r;. where H,;, is the virial radius of the halo discussed earlier. and ry is a scale radius."," Other than the virial mass, this profile is described by a concentration parameter $c \equiv R_{vir}/r_s$, where $R_{vir}$ is the virial radius of the halo discussed earlier, and $r_s$ is a scale radius."
+ For e. we take into account both (he mass and redshilt dependence and use the relation of Bullock et al. (," For $c$ , we take into account both the mass and redshift dependence and use the relation of Bullock et al. ("
+"2001): ον}= P. where M,=15xLOMATAL.","2001): $c(M,z)=9
+(1+z)^{-1} (M/M_*)^{-0.13}$ , where $M_*=1.5\times 10^{13}
+h^{-1}M_\odot$."
+" The density amplitude 9 is related toe by à=A,c'[ln(l+e)—e/(1ο)1/3.", The density amplitude $\bar\delta$ is related to $c$ by $\bar\delta=\Delta_{vir} c^3 [\ln(1+c)-c/(1+c)]^{-1}/3$.
+" The lensing cross section for NEW halos is determined by the parameter 5j=p9r,/X,. where X,=(7?/daG)(D./D,D).) is the critical surface mass densitv."," The lensing cross section for NFW halos is determined by the parameter $\kappa_0=\bar\rho \bar\delta\,r_s/\Sigma_{cr}$, where $\Sigma_{cr}=(c^2/4\pi G)(D_s/D_l D_{ls}) $ is the critical surface mass density."
+" An NEW halo will have multiple images if the source is within (he radial caustic of angular size 2,4, [rom the halo center.", An NFW halo will have multiple images if the source is within the radial caustic of angular size $\beta_{rad}$ from the halo center.
+" We compute δρ as a Iunction of halo mass. lens redshilt. aud source redshilt bv solving; d32/00—0. where 2 ancl 9 are the positions of the source and (he image. respectively,"," We compute $\beta_{rad}$ as a function of halo mass, lens redshift, and source redshift by solving $d\beta/d\theta=0$, where $\beta$ and $\theta$ are the positions of the source and the image, respectively."
+ The cross section for lensing bv an NEW halo is then On.πω)”., The cross section for lensing by an NFW halo is then $\sigma_{\rm lens}=\pi (\beta_{rad} D_l)^2$.
+ We also need to determine the halo mass M; in eq. (, We also need to determine the halo mass $M_{min}$ in eq. (
+1).,1).
+ llere we use the [act that the angular separation of the outermost images is insensitive to the value of 2 (Schneider et al., Here we use the fact that the angular separation of the outermost images is insensitive to the value of $\beta$ (Schneider et al.
+ 1992) and simplify the calculations by using; ;?—0 for a perfectly aligned source-lens configuration., 1992) and simplify the calculations by using $\beta=0$ for a perfectly aligned source-lens configuration.
+" We assume a “cooling mass” of M,z1.5:1075TAL. above which the lenses are assigned (he NEW prolile and below which the lenses are SIs."," We assume a “cooling mass” of $M_c \approx 1.5\cdot
+10^{13} h^{-1} M_{\odot}$, above which the lenses are assigned the NFW profile and below which the lenses are SIS."
+ Uncertainties introduced by AM. on (he lensing probability are discussed in 3., Uncertainties introduced by $M_c$ on the lensing probability are discussed in 3.
+ To compare the predicted lensing probabilities withobservational results. we use the combined data from JVAS and CLASS. which olfer the largest uniformly selected sample of," To compare the predicted lensing probabilities withobservational results, we use the combined data from JVAS and CLASS, which offer the largest uniformly selected sample of"
+data.,data.
+ As cleseribecl in Paper IV. we find 4=0.54+0.21. for lines above a 0.24 (threshold. while Wevinann οἱ ((1993) measure y=0.15+0.23.," As described in Paper IV, we find $\gamma=0.54\pm0.21$, for lines above a 0.24 threshold, while Weymann et (1998) measure $\gamma=0.15\pm0.23$."
+ See Paper IV [or more discussion of the significance and underlying causes of this difference., See Paper IV for more discussion of the significance and underlying causes of this difference.
+ Wefind ΟΛte.)0.61.5.5x107.7.07) and (1.910.0.88.23.4x107.6.21) for (Q\).Q4)=(1..0.) and lines with rest equivalent widths above 0.24 and 0.32 respectively.," Wefind $(A_{\rm HM},B_{\rm HM},z_{c},S)=(3.0\times 10^{-12},0.61,5.5 \times 10^{-7}, 7.07)$ and $(1.9 \times 10^{-11}, 0.38, 3.4 \times 10^{-7}, 6.21)$ for $(\Omega_{\rm M},\Omega_{\Lambda})=(1.,0.)$ and lines with rest equivalent widths above 0.24 and 0.32 respectively."
+ These fitsto Equ., These fitsto Equ.
+ 19. are shown in Figure 13((a)., \ref{equ:dave} are shown in Figure \ref{fig:dndz}( (a).
+ In panel (b). we plot L(z). as expressed in Equ. 14..," In panel (b), we plot $\Gamma (z)$, as expressed in Equ. \ref{equ:hmgam},"
+ evaluated using the parameters found [rom the fit to Equ., evaluated using the parameters found from the fit to Equ.
+ 19 above., \ref{equ:dave} above.
+ The IIM96 solution and (he solution derived from the full FOS ancl MMT data sets are represented by the thick and thin solid lines respectively., The HM96 solution and the solution derived from the full FOS and MMT data sets are represented by the thick and thin solid lines respectively.
+" The small values of z, derived from dA/dz above translate into ionization rates that do not decrease dramatically with decreasing redshift and result. [rom the less pronounced flattening of 4/dz relative to the Kev Project.", The small values of $z_{c}$ derived from $d{\cal N}/dz$ above translate into ionization rates that do not decrease dramatically with decreasing redshift and result from the less pronounced flattening of $d{\cal N}/dz$ relative to the Key Project.
+ These fits are particularly insensitive to the normalization. μι. so the errors on this parameter are large.," These fits are particularly insensitive to the normalization, $A_{\rm HM}$, so the errors on this parameter are large."
+ These fits should therefore not be interpreted as measurements of Ες) as reliable as those found directly. from (he absorption line data., These fits should therefore not be interpreted as measurements of $\Gamma (z)$ as reliable as those found directly from the absorption line data.
+ But we find them instructive nonetheless., But we find them instructive nonetheless.
+ The observed Ες) Falls short of the ionization rate needed to fully account For the change in the Ly-a line density with redshift. indicating that if the value of > al low redshift is indeed slightly larger than that found by the Nev Project. A/dz maw still be consistent with a non-evolving population of Ly-a absorbers in the sense noted above. but the formation of structure in the low redshift universe must play a signilicant role in determining the character of the Ly-a forest line density.," The observed $\Gamma (z)$ falls short of the ionization rate needed to fully account for the change in the $\alpha$ line density with redshift, indicating that if the value of $\gamma$ at low redshift is indeed slightly larger than that found by the Key Project, $d{\cal N}/dz$ may still be consistent with a non-evolving population of $\alpha$ absorbers in the sense noted above, but the formation of structure in the low redshift universe must play a significant role in determining the character of the $\alpha$ forest line density."
+ KF93 performed a similar measurement with a small subsample of this total sample- ihe IST Quasar Absorption Line Kev Project data of Baheall et ((1993)., KF93 performed a similar measurement with a small subsample of this total sample- the HST Quasar Absorption Line Key Project data of Bahcall et (1993).
+ We compare our result to that from Sample 2 of IKF93. which was constructed from the Baheall et ((1993) data excluding one BAL quasar and all heavy. element absorption svstems.," We compare our result to that from Sample 2 of KF93, which was constructed from the Bahcall et (1993) data excluding one BAL quasar and all heavy element absorption systems."
+ The Ixev Project sample has since been supplemented (Bahcall et 11996. Jannuzi οἱ al.," The Key Project sample has since been supplemented (Bahcall et 1996, Jannuzi et al."
+ 1993) and those data have been included when appropriate in the complete archival sample of FOS spectra presented in Paper LI., 1998) and those data have been included when appropriate in the complete archival sample of FOS spectra presented in Paper III.
+" The mean intensity IKF93 derive from their Sample 2 (b=35 kms |. 3=148. 5—0.21) aU), ⋅↽↽ ⋅↽ ↕⋟∖⊽⇀↱≻⋅∩⋮⋡⋝⋡↓⋗⋖⊥∩−⋟↕≼↲↕⋅≸↔↴⋟∖⋃∖⊽↓≺∢∐↓−⋟∐∠↓⋟∖⊽↕⋅↓⋅⊺↥∐⋟∖⇁↕⋅≼↲⋟∖⊽∏"," The mean intensity KF93 derive from their Sample 2 $b=35$ km $^{-1}$, $\beta$ =1.48, $\gamma$ =0.21) is $5.0^{+20.}_{-3.4} \times 10^{-24}$ ergs $^{-1}$ $^{-2}$ $^{-1}$ $^{-1}$."
+∐↕⋟∖⊽↥∪∖∖⇁≼↲↕⋅⊔↥≀↧↴∐∪⋯⋅⋟∖⊽↓∪↕⋅⊳∶⊥∣≻⋡∖↽≀↕↴ Th; ∕↽↾ factor of ~ 13. though the errors are large on both results are large enough that they are consistent.," This result is lower than ours for $z < 1$ by a factor of $\sim 13$ , though the errors are large on both results are large enough that they are consistent."
+ We use 162 lines in our low redshift solution for (να). 65 more than KF93.," We use 162 lines in our low redshift solution for $J(\nu_{0})$ , 65 more than KF93."
+dvnamical characterises of their small bodies in more detail.,dynamical characteristics of their small bodies in more detail.
+ We acknowledge the use of the computational facilities at the Department of Terrestrial Magnetism al the Carnegie Institution of Washington., We acknowledge the use of the computational facilities at the Department of Terrestrial Magnetism at the Carnegie Institution of Washington.
+ This work has been supported by the NASA Astrobiology Institute uncer Cooperative Agreement NNAOLICCOSA at the Institute for Astronomy at the University of Hawaii for NI., This work has been supported by the NASA Astrobiology Institute under Cooperative Agreement NNA04CC08A at the Institute for Astronomy at the University of Hawaii for NH.
+velocity). and total space velocity 27.=C7|VELAE.,"velocity), and total space velocity $T^2=U^2+V^2+W^2$."
+ Space velocities were corrected to the Local Standard of Rest (LSR: Dehnen&Binney 1998)). with CU positive toward the Galactic anticenter. V positive in the direction of Galactic rotation. and M positive toward the north Galactic pole.," Space velocities were corrected to the Local Standard of Rest (LSR; \citealt{deh98}) ), with $U$ positive toward the Galactic anticenter, $V$ positive in the direction of Galactic rotation, and $W$ positive toward the north Galactic pole."
+ In the absence of correction to the Εμ]. the assumption of zero racial velocity is equivalent to 2=(i.," In the absence of correction to the LSR, the assumption of zero radial velocity is equivalent to $T={v_{\rm tan}}$."
+ Tangential speeds were calculated [rom The Sun does not have zero velocity relative to the LSR: depending on viewing direction. the Solar motion contributes Cua=LO15 (Dehnen&Binney1998:Mihalas&Binney 1981).," Tangential speeds were calculated from The Sun does not have zero velocity relative to the LSR; depending on viewing direction, the Solar motion contributes $v_{\rm tan}=10-15$ $^{-1}$ \citep{deh98,mih81}."
+. Table 2. lists the calculated parameter means for the SDSS DZ and DC white dwarfs., Table \ref{tbl2} lists the calculated parameter means for the SDSS DZ and DC white dwarfs.
+ As a first attempt to establish the relationship. if any. between the DZ stars ancl the LSAL one can search for spatial or kinematical correlations with calcium abundance (Zuckermanetal.2003:Aannestad1993).," As a first attempt to establish the relationship, if any, between the DZ stars and the ISM, one can search for spatial or kinematical correlations with calcium abundance \citep{zuc03,
+aan93}."
+.. Figure l plots the calcium. abundances (throughout. this paper AN/Y] = logn(CN)/n(Y)D. relative to helium. of the 146 DZ stars as a function of cllective temperature. height above the Galactic mid-plane. ancl tangential speed.," Figure \ref{fig1} plots the calcium abundances (throughout this paper [X/Y] = $\log
+\,[n({\rm X})/n({\rm Y})]$ ), relative to helium, of the 146 DZ stars as a function of effective temperature, height above the Galactic mid-plane, and tangential speed."
+ A ow abundance eutolf at higher cllective temperatures. and arecr distances is apparent in the upper ancl miele xiuiels: these are observational biases arising [rom higher atmospheric opacities ancl lower spectroscopic sensitivity. respectively Dufouretal.2007:Ixoester20052a).," A low abundance cutoff at higher effective temperatures and larger distances is apparent in the upper and middle panels; these are observational biases arising from higher atmospheric opacities and lower spectroscopic sensitivity, respectively \citep{duf07,koe05a}."
+. The classical conundrum of the DZ white dwarls presents itself well in the upper panel: the majority of stars have no detected hyvelrowen. including those with the highest calcium abundances.," The classical conundrum of the DZ white dwarfs presents itself well in the upper panel; the majority of stars have no detected hydrogen, including those with the highest calcium abundances."
+ The micelle and lower panels of Figure 1 are noteworthy because they fail to reveal a physical correlation between acereted calcium and. height above the Calactic disk or tangential speed., The middle and lower panels of Figure \ref{fig1} are noteworthy because they fail to reveal a physical correlation between accreted calcium and height above the Galactic disk or tangential speed.
+ The Galactic gas ane dust laver extends roughly ppc above the Galactic mid-plane. and vet some of most highly polluted DZ white dwarls lie well above this region.," The Galactic gas and dust layer extends roughly pc above the Galactic mid-plane, and yet some of most highly polluted DZ white dwarfs lie well above this region."
+ The lower panel is relevant in the context of gravitationally driven aceretion. expected. if the DZ stars obtain their heavy elements. while moving through the ISM.," The lower panel is relevant in the context of gravitationally driven accretion, expected if the DZ stars obtain their heavy elements while moving through the ISM."
+ “Phe lack of correlation between speed and the metal abundances argues. qualitatively but strongly. against ILuid. accretion of ISM.," The lack of correlation between speed and the metal abundances argues, qualitatively but strongly, against fluid accretion of ISM."
+ Figure 2. plots the masses of calcium contained in the convection zones of these metal-enriched stars using their abundances together with the convective envelope. masses [or logg=8 heliumerich white chvarls (οσο2009).," Figure \ref{fig2} plots the masses of calcium contained in the convection zones of these metal-enriched stars using their abundances together with the convective envelope masses for $\log\,g=8$ helium-rich white dwarfs \citep{koe09}."
+ The figure shows that the most highly pollutecl DZ stars currently harbor up to 1077 gg of caleium in their convection zones., The figure shows that the most highly polluted DZ stars currently harbor up to $^{22}$ g of calcium in their convection zones.
+ Phis is a truly remarkable amount of a single heavy element. a mass in calcium alone that is roughly equivalent to the total mass of a 2O0kkm diameter. asteroicl(!).," This is a truly remarkable amount of a single heavy element, a mass in calcium alone that is roughly equivalent to the total mass of a km diameter asteroid(!)."
+ A more tvpical: DZ. star appears to contain: around 20 ee of calcium. still a prodigious amount. anc corresponds to( a time-averaged metal aceretion rate of 2.10 eess+ (Paribietal. 2009).. and a total accreted. mass of metals of just under 1077 ee. ligure 3. plots the hydrogen abundances in 37 of the 146 DZ stars where hydrogen is detected or inferred (Dufouretal. 2007)... versus height above the Galactic mid-plane and tangential speed.," A more typical DZ star appears to contain around $^{20}$ g of calcium, still a prodigious amount, and corresponds to a time-averaged metal accretion rate of $2\times10^8$ $^{-1}$ \citep{far09a}, and a total accreted mass of metals of just under $10^{22}$ g. Figure \ref{fig3} plots the hydrogen abundances in 37 of the 146 DZ stars where hydrogen is detected or inferred \citep{duf07}, versus height above the Galactic mid-plane and tangential speed."
+ No obvious pattern is seen. although there may be a higher density of DZA stars near the Galactic disk and perhaps also toward more modest speeds. but the former may be an observational bias due to. diminished spectroscopic sensitivity.," No obvious pattern is seen, although there may be a higher density of DZA stars near the Galactic disk and perhaps also toward more modest speeds, but the former may be an observational bias due to diminished spectroscopic sensitivity."
+ In any case. one should expect a correlation between these quantities if the hydrogen. were accreted [rom ISM.," In any case, one should expect a correlation between these quantities if the hydrogen were accreted from ISM."
+ As discussed by. Dupuisetal.(1993a)... white divarfs accreting Solar abundance matter within the ISM at 310 3 for 10 vvr can accumulate the heavy. element mass fractions observed in the convective envelopes of DZ stars. such as plotted in Figure 2 (however. this cannot account for their hydrogen abundances. discussed in detail below).," As discussed by \citet{dup93a}, white dwarfs accreting Solar abundance matter within the ISM at $3\times10^{11}$ $^{-1}$ for $^6$ yr can accumulate the heavy element mass fractions observed in the convective envelopes of DZ stars, such as plotted in Figure \ref{fig2}
+ (however, this cannot account for their hydrogen abundances, discussed in detail below)."
+ Yet some care must be taken not to invoke such high accretion rates without skepticism: while these rates are sullicient to account for the calcium data (that is the nature of the hypothesis). the physical plausibility is of equal importance.," Yet some care must be taken not to invoke such high accretion rates without skepticism; while these rates are sufficient to account for the calcium data (that is the nature of the hypothesis), the physical plausibility is of equal importance."
+ Below. the problem. of interstellar accretion," Below, the problem of interstellar accretion"
+"For each of the 60 sources, we were able to investigate the relation between the y-ray photon flux and 37 GHz flux densities during the 11-month 1FGL period.","For each of the 60 sources, we were able to investigate the relation between the $\gamma$ -ray photon flux and 37 GHz flux densities during the 11-month 1FGL period."
+" However, only for a subsample of 45 sources did we achieve an adequate decomposition of the total flux density curves, allowing us to categorize in detail the radio flare state during the y-ray maxima."," However, only for a subsample of 45 sources did we achieve an adequate decomposition of the total flux density curves, allowing us to categorize in detail the radio flare state during the $\gamma$ -ray maxima."
+" Using the finely sampled light curves from the Metsáhhovi quasar monitoring program, we compare simultaneous y-ray fluxes and 37 GHz flux densities."," Using the finely sampled light curves from the Metsähhovi quasar monitoring program, we compare simultaneous $\gamma$ -ray fluxes and 37 GHz flux densities."
+ Figure 1 shows that simultaneous measurements at the two bands appear to be positively correlated., Figure 1 shows that simultaneous measurements at the two bands appear to be positively correlated.
+" However, we find significant differences between quasars and BL Lacs, which we describe below."," However, we find significant differences between quasars and BL Lacs, which we describe below."
+" By applying the Spearman's rank correlation test, two very clear results emerge."," By applying the Spearman's rank correlation test, two very clear results emerge."
+" First, there is a significant positive correlation between the y-ray photon flux and the 37 GHz flux density for quasars, while the BL Lac fluxes are not correlated."," First, there is a significant positive correlation between the $\gamma$ -ray photon flux and the 37 GHz flux density for quasars, while the BL Lac fluxes are not correlated."
+" Second, the strength and the significance of the correlations is different for each type of quasar."," Second, the strength and the significance of the correlations is different for each type of quasar."
+" The photon flux - flux density correlation for quasars is absent for QSOs, significant (p=0.47,P,,j 99.9%) for LPQs, and very significant (o=0.50,P,,;> 99.9%) for HPQs."," The photon flux - flux density correlation for quasars is absent for QSOs, significant $\rho= 0.47, P_{null}= 99.9\%$ ) for LPQs, and very significant $\rho= 0.50, P_{null} > 99.9\%$ ) for HPQs."
+" Such a dependence on the degree of optical polarization may arise naturally if the polarization indicates the viewing angle of the jet, with sources with high optical polarization having their jets oriented closest to our line of sight ?).."," Such a dependence on the degree of optical polarization may arise naturally if the polarization indicates the viewing angle of the jet, with sources with high optical polarization having their jets oriented closest to our line of sight \citep[e.g.][]{hovatta_2009}."
+" The dependence of the flux - flux relation on optical polarization agrees with previous results, where it has been shown that the brightest y-ray emitters have preferentially smaller viewing angles (??) and consequently higher Doppler boosting factors (???).."," The dependence of the flux - flux relation on optical polarization agrees with previous results, where it has been shown that the brightest $\gamma$ -ray emitters have preferentially smaller viewing angles \citep{valtaoja_1995,anne_2003} and consequently higher Doppler boosting factors \citep{lister_2009,savolainen_2010,tornikoski_2010}."
+" Since y-ray fluxes and radio flux densities are significantly correlated for sources where the relativistic jet is aligned close to our line of sight, this implies that there is a strong coupling between the radio and the y-ray emission mechanisms."," Since $\gamma$ -ray fluxes and radio flux densities are significantly correlated for sources where the relativistic jet is aligned close to our line of sight, this implies that there is a strong coupling between the radio and the $\gamma$ -ray emission mechanisms."
+ That the correlation is seen on monthly, That the correlation is seen on monthly
+a recurrent nova.,a recurrent nova.
+ However. only one outburst has been recorded for V2672 Oph. but it is highly. probable that several others have been missed due to the following reasons: (1) V2672 Oph lies less than 4° away from the ecliptic. and thus sulfers from. periods of long seasonal. invisibility due to the conjunction. with the Sun.," However, only one outburst has been recorded for V2672 Oph, but it is highly probable that several others have been missed due to the following reasons: (1) V2672 Oph lies less than $^\circ$ away from the ecliptic, and thus suffers from periods of long seasonal invisibility due to the conjunction with the Sun."
+ Furthermore. at every lunation. it also sulfers from. the proximity to the Moon: (2) its 6 = 26°44 southern declination means that it is only observable for a brief fraction. of the vear bv northern latitudes where most. observers have been historically concentrated: (3) the faint peak brightness attained by V2672 Oph (V 11.35. D—13.1) means that it would have remained outside the range of detection bv patrol surveys until very recently. when the introduction of large format CCDs allowed. coverage of signilicant. areas of the skv with short focal length. telescopes.," Furthermore, at every lunation it also suffers from the proximity to the Moon; (2) its $\delta$ = $-$ $^\circ$ $^\prime$ southern declination means that it is only observable for a brief fraction of the year by northern latitudes where most observers have been historically concentrated; (3) the faint peak brightness attained by V2672 Oph $V$$\sim$ 11.35, $B$$\sim$ 13.1) means that it would have remained outside the range of detection by patrol surveys until very recently, when the introduction of large format CCDs allowed coverage of significant areas of the sky with short focal length telescopes."
+ In fact. V2672 Oph was discovered by Ix. Iagaki using a 0.21-m [/3 rellector. a light-collector far more sensitive that the usual digital SLR cameras used by Japanese amateurs to discover. Galactic novae: and (4) a decline time /2(1 22.3 days means that the nova would have returned below the threshold for discovery (V 14) in just three days.," In fact, V2672 Oph was discovered by K. Itagaki using a 0.21-m f/3 reflector, a light-collector far more sensitive that the usual digital SLR cameras used by Japanese amateurs to discover Galactic novae; and (4) a decline time $t_2(V)$ =2.3 days means that the nova would have returned below the threshold for discovery $V$$\sim$ 14) in just three days."
+ Given all these restrictions. it is indeed surprising that even this outburst of V2672 Oph has been discovered.," Given all these restrictions, it is indeed surprising that even this outburst of V2672 Oph has been discovered."
+ The central region of the Galaxy has been imaged many times especially by amateurs looking for impressive pictures., The central region of the Galaxy has been imaged many times especially by amateurs looking for impressive pictures.
+ Lt is quite possible that other outbursts of V2672 Oph lic uncletected on such archive images. especially those imaging into red wavelengths.," It is quite possible that other outbursts of V2672 Oph lie undetected on such archive images, especially those imaging into red wavelengths."
+ A devoted search is highly encouraged., A devoted search is highly encouraged.
+ Lt is less probable that an outburst could be discovered. by inspection of historical plate archives., It is less probable that an outburst could be discovered by inspection of historical plate archives.
+ In fact. the B~13.1 mag attained at maximum. places V2672 Oph below the limiting magnitude of most patrol plates collected worldwide in the past.," In fact, the $B$$\sim$ 13.1 mag attained at maximum, places V2672 Oph below the limiting magnitude of most patrol plates collected worldwide in the past."
+ Using the morpho-kinematical code (Version3.56.stellen&Lopez2006:Stellenetal.2010) we have analysecl ancl disentangled the threc-dimensional geometry and kinematic structure of the early. outburst spectra of V2672 Oph.," Using the morpho-kinematical code \citep[Version 3.56,][]{SL06,SKW10}
+ we have analysed and disentangled the three-dimensional geometry and kinematic structure of the early outburst spectra of V2672 Oph."
+ was originally developed. to model the complex structures of planetary nebulae. ane is based on computationally efficient. mathematical. representations of the visual world which allows for the construction of objects placed at any orientation in a cubic volume., was originally developed to model the complex structures of planetary nebulae and is based on computationally efficient mathematical representations of the visual world which allows for the construction of objects placed at any orientation in a cubic volume.
+ has been developed so far for modelling optically thin environments. therefore. one must make the assumption that for what is observed the optical depth is low and that absorption has not altered greatly the shape of the profile.," has been developed so far for modelling optically thin environments, therefore, one must make the assumption that for what is observed the optical depth is low and that absorption has not altered greatly the shape of the profile."
+" The adopted: model is based on previous studies of classical novae. which explored. the structures of classical novae from resolved. optical imaging ane hyverodyvnamical modelling (e.g.Slavin,O'Brien.&Dunlop1995:Lloyd.O'Brien.&Bocle 1997)."," The adopted model is based on previous studies of classical novae, which explored the structures of classical novae from resolved optical imaging and hydrodynamical modelling \citep[e.g.][]{SOD95,LOB97}. ."
+. We performed Gaussian fitting using the LRAL taskSPECFIT on days. | 2.34 through [8.33 after outburst., We performed Gaussian fitting using the IRAF task on days $+$ 2.34 through $+$ 8.33 after outburst.
+ These allow us to decompose dillerent Gaussian components of the La lino and retrieve information such as the ENIM of likely. components and their radial velocity. displacements. (Lable 4))., These allow us to decompose different Gaussian components of the $\alpha$ line and retrieve information such as the FWHM of likely components and their radial velocity displacements (Table \ref{tb:gauss}) ).
+ We note the presence of the DID at 6614 A(see Figure S and sect., We note the presence of the DIB at 6614 (see Figure \ref{fig:early} and sect.
+ 3.2). which is not considered. curing detailed modelling given its small equivalent width.," 3.2), which is not considered during detailed modelling given its small equivalent width."
+ The values derived in Table 4 are used. to find the displacement of the system from the rest wavelength of the La line and to determine the size of the remnant using the values forthe ENIM ancl racial velocity displacement. of component 1., The values derived in Table \ref{tb:gauss} are used to find the displacement of the system from the rest wavelength of the $\alpha$ line and to determine the size of the remnant using the values forthe FWHM and radial velocity displacement of component 1.
+"The optical afterglow light curve of GRB 050904 is well fitted by a smoothed broken power-law model between 0.1 and 10 days after the trigger. with best-fit decay indexes «e,=0.724 and ao=2.4x0.4 and a break time at 4%=2.6+1.0 days (2). ","The optical afterglow light curve of GRB 050904 is well fitted by a smoothed broken power-law model between 0.1 and 10 days after the trigger, with best-fit decay indexes $\alpha_1=0.72\pm0.10$ and $\alpha_2=2.4\pm0.4$ and a break time at $t_b=2.6\pm1.0$ days \citep{Tagliaferri2005}. ."
+"2. quote a;=0.85+0.08 and similar a> and ty, fitting optical data from 7+0.3 days.", \cite{Kann2007} quote $\alpha_1=0.85\pm0.08$ and similar $\alpha_2$ and $t_b$ fitting optical data from $T+0.3$ days.
+ The temporal break is achromatic. and therefore it has been interpreted as evidence for a jet with opening angle @~Ε/Γ). where Γή) is the ejecta Lorentz factor at the epoch of the temporal break (e.g. ?)..," The temporal break is achromatic, and therefore it has been interpreted as evidence for a jet with opening angle $\theta\sim 1/\Gamma(t_b)$, where $\Gamma(t_b)$ is the ejecta Lorentz factor at the epoch of the temporal break \citep[e.g.][]{Sari1999}."
+ We thus simply normalize the best-fit optical light curve model (from Taghaferrt et al., We thus simply normalize the best-fit optical light curve model (from Tagliaferri et al.
+ 2007) at the 7+3.25 day unabsorbed flux value (see Fig.l)., 2007) at the $T+3.25$ day unabsorbed flux value (see Fig.1).
+ The X-ray flux extrapolations at 7+0.47. T+1.25 days. and 7+3.4 days. are 0.03. 0.07. and 0.001 pJy.," The X-ray flux extrapolations at $T+0.47$, $T+1.25$ days, and $T+3.4$ days, are 0.03, 0.07, and 0.001 $\mu$ Jy."
+ We associated an uncertainty of 25% to the former two values and of about 50% to the last value. encompassing the uncertainty of the late X-ray flux measures.," We associated an uncertainty of $25\%$ to the former two values and of about $50\%$ to the last value, encompassing the uncertainty of the late X-ray flux measures."
+ Starting from the results obtained in the previous section. we take from the literature the optical/near-IR data corrected for Galactic absorption at T+0.47. 1.25. and 3.4 days. where the photometric data in the Z band is taken from the accurate re- recently published by ?.. and we compare them with the simultaneous X-ray afterglow unabsorbed fluxes obtained in the previous section.," Starting from the results obtained in the previous section, we take from the literature the optical/near-IR data corrected for Galactic absorption at T+0.47, 1.25, and 3.4 days, where the photometric data in the Z band is taken from the accurate re-analysis recently published by \cite{Zafar2010}, and we compare them with the simultaneous X-ray afterglow unabsorbed fluxes obtained in the previous section."
+ We then fit a power-law spectral model to the data. setting the spectral index and normalization to the values obtained from X-ray afterglow.," We then fit a power-law spectral model to the data, setting the spectral index and normalization to the values obtained from X-ray afterglow."
+" No spectral variation 1s expected during each epoch since v.«vy,,,, at these times (see 82.2)."," No spectral variation is expected during each epoch since $\nu_c<\nu_{X,opt}$ at these times (see 2.2)."
+" Given the large uncertainty affecting the X-ray photon index we have estimated at T+3.254 days. we conservatively performed our analysis using both the expected (Ες=2.15€ 0.10. BO""=1.15x 0.10) and the measured (Ty=2.4+0.5. px= 1440.5) X-ray photon indexes."," Given the large uncertainty affecting the X-ray photon index we have estimated at $T+3.254$ days, we conservatively performed our analysis using both the expected $\Gamma_X^{exp}=2.15\pm0.10$ , $\beta_X^{exp}=1.15\pm0.10$ ) and the measured $\Gamma_X=2.4\pm0.5$, $\beta_X=1.4\pm0.5$ ) X-ray photon indexes."
+ Figure 2. shows the obtained broad band SEDs at the three epochs., Figure \ref{fig:f2} shows the obtained broad band SEDs at the three epochs.
+ Squares give the X-ray fluxes estimated by us in the previous section. while diamonds are the values reported in ?..," Squares give the X-ray fluxes estimated by us in the previous section, while diamonds are the values reported in \cite{Zafar2010}."
+ We note that our X-ray fluxes deviate significantly from those reported in ?. at T+0.47 and T+1.25 days., We note that our X-ray fluxes deviate significantly from those reported in \cite{Zafar2010} at $T+0.47$ and $T+1.25$ days.
+ The latter discrepancy may be due to different light curve modeling and flux extrapolation., The latter discrepancy may be due to different light curve modeling and flux extrapolation.
+ In ?. the X-ray light curve is modeled with a smoothed broken power-law. however parameter values are not reported. hence preventing further comparative studies.," In \cite{Zafar2010} the X-ray light curve is modeled with a smoothed broken power-law, however parameter values are not reported, hence preventing further comparative studies."
+ The hatched areas (delimited by dashed lines) indicate the range of intrinsic power-laws consistent with the £2c7 uncertainty on By., The hatched areas (delimited by dashed lines) indicate the range of intrinsic power-laws consistent with the $\pm2\sigma$ uncertainty on $\beta_X$.
+ The dashed line in the middle shows the power-law resulting from the best-fit slope to the X-ray data at late epochs 1.4)., The dashed line in the middle shows the power-law resulting from the best-fit slope to the X-ray data at late epochs $\beta_X=1.4$ ).
+ The gray shaded areas show the range of intrinsic power-laws consistent with the +lo uncertainty on By., The gray shaded areas show the range of intrinsic power-laws consistent with the $\pm1\sigma$ uncertainty on $\beta_X$.
+ The cyan shaded areas show the power-laws expected from synchrotron emission (for an electron spectral index p.~2.]-2.5 and VeXYyoy. see 32.2).," The cyan shaded areas show the power-laws expected from synchrotron emission (for an electron spectral index $p\sim2.1-2.5$ and $\nu_c<\nu_{X,opt}$, see 2.2)."
+ The colored areas show the optical to X-ray SED best-fit obtained assuming a power-law model with index free to vary within the range By+2c (Tab., The colored areas show the optical to X-ray SED best-fit obtained assuming a power-law model with index free to vary within the range $\beta_X\pm2\sigma$ (Tab.
+ 3) and the extinction/attenuation curve associated with each panel., 3) and the extinction/attenuation curve associated with each panel.
+ To quantify the need of dust extinction and reddening we have fitted the data at each epoch with different extinction and attenuation curves., To quantify the need of dust extinction and reddening we have fitted the data at each epoch with different extinction and attenuation curves.
+ The question of whether dust reddening is better described by an extinction curve or by an attenuation law depends on the geometry of the system., The question of whether dust reddening is better described by an extinction curve or by an attenuation law depends on the geometry of the system.
+ For quasars. the simple dusty “screen” geometry applies. but for galaxies one has to consider that generally dust i5 mixed with the emitting sources (either stars or tonized gas).," For quasars, the simple dusty “screen” geometry applies, but for galaxies one has to consider that generally dust is mixed with the emitting sources (either stars or ionized gas)."
+ As a consequence. attenuation curves Le. the ratio of the observed spectrum to the intrinsic total light emitted by the whole system before dust absorption. are likely more appropriate for galaxies.," As a consequence, attenuation curves i.e. the ratio of the observed spectrum to the intrinsic total light emitted by the whole system before dust absorption, are likely more appropriate for galaxies."
+ In the case of GRBs it is not clear which geometry is more appropriate., In the case of GRBs it is not clear which geometry is more appropriate.
+ The screen case is certainly more appropriate when the size of the emitting optical afterglow is smaller than the distribution of the dusty medium., The screen case is certainly more appropriate when the size of the emitting optical afterglow is smaller than the distribution of the dusty medium.
+ However. the dust may well be the same as produced by the progenitor before the explosion.," However, the dust may well be the same as produced by the progenitor before the explosion."
+ In this case the size of the dusty medium can be comparable to the one of the expanding afterglow. and an attenuation curve may be more appropriate.," In this case the size of the dusty medium can be comparable to the one of the expanding afterglow, and an attenuation curve may be more appropriate."
+ As a consequence. to describe dust reddening. we consider the SMC extinetion curve. the Calzetti attenuationlaw.. the mean extinction curve (MEC) resulting from the analysis of 33 quasars at z>4 done by ?.. and the corresponding attenuation curve (ΜΕ).," As a consequence, to describe dust reddening, we consider the SMC extinction curve, the Calzetti attenuation, the mean extinction curve (MEC) resulting from the analysis of 33 quasars at $\geq 4$ done by \cite{Gallerani2010}, and the corresponding attenuation curve $_{att}$ )."
+ We also test the SN-type extinction curve proposed by ?.. which reproduces the dust extinctior observedina BAL QSO at z=6.2 (??)..," We also test the SN-type extinction curve proposed by \cite{Todini2001}, which reproduces the dust extinction observedin a BAL QSO at z=6.2 \citep{Maiolino2004,Gallerani2010}. ."
+" In Figure 3.. we plot the extinction and attenuation curves adopted in our analysis. normalized to Aaooo. Le. to the extinction value at the rest frame wavelength of 1,,,,=3000A."," In Figure \ref{fig:f3}, we plot the extinction and attenuation curves adopted in our analysis, normalized to $A_{3000}$, i.e. to the extinction value at the rest frame wavelength of $\lambda_{rest}=3000$."
+ We report the best-fit results for each extinction/attenuation curve and at each epoch in Tables 2.. 3.. and 4.. where we give the best-fit By and Asooo. as well as the resulting best-fit reduced $ and the associated probability.," We report the best-fit results for each extinction/attenuation curve and at each epoch in Tables \ref{tab:t4}, , \ref{tab:t5}, and \ref{tab:t6}, where we give the best-fit $\beta _X$ and $A_{3000}$, as well as the resulting best-fit reduced $\tilde{\chi}^2$ and the associated probability."
+" Within each table we leave By free to vary in different intervals: By€ [0.9-1.9] in Table 2 and Bye [0.4-2.4] in Table 3.. which correspond to the Io and 20 errors on the measured X-ray photon index at late times. respectively (see $22.1): By€ [1.05-1.25] in Table 4+ is the range expected from synchrotron emission with v.€vy, and p~2.]—2.5 (see 822.2)."," Within each table we leave $\beta_X$ free to vary in different intervals: $\beta _X~\epsilon~$ [0.9-1.9] in Table \ref{tab:t4} and $\beta _X~\epsilon~$ [0.4-2.4] in Table \ref{tab:t5}, which correspond to the $\sigma$ and $\sigma$ errors on the measured X-ray photon index at late times, respectively (see 2.1); $\beta _X~\epsilon~$ [1.05-1.25] in Table \ref{tab:t6} is the range expected from synchrotron emission with $\nu_c<\nu_{X,opt}$ and $p\sim2.1-2.5$ (see 2.2)."
+ The data in bold highlight the extinction/attenuation curve giving a good fit to the data (PU?<y)« 68%)) at all epochs., The data in bold highlight the extinction/attenuation curve giving a good fit to the data $P(\tilde{\chi}^2<\tilde{\chi}^2_{BF})<68$ ) at all epochs.
+ In some cases the By best-fit value coincides with the lowerextreme of the interval considered., In some cases the $\beta_X$ best-fit value coincides with the lowerextreme of the interval considered.
+ In these cases. we cannot compute the Ic error on Agony and we can only provide an interval within. which this parameter is allowed tovary.," In these cases, we cannot compute the $\sigma$ error on $A_{3000}$ and we can only provide an interval within which this parameter is allowed tovary."
+ Moreover. for some combinations of Byintervals and extinction curves. the best-fit values are characterized by a P«y 68%:; In these cases. we only provide the best-fit values.," Moreover, for some combinations of $\beta_X$intervals and extinction curves, the best-fit values are characterized by a $P(\tilde{\chi}^2<\tilde{\chi}^2_{BF})>68$ ; in these cases, we only provide the best-fit values."
+ Tables 2 to 4. show that. regardless of the interval adopted for fx and taking only the best-fit with Ρο«y)68% into," Tables \ref{tab:t4} to \ref{tab:t6} show that, regardless of the interval adopted for $\beta_X$ and taking only the best-fit with $P(\tilde{\chi}^2<\tilde{\chi}^2_{BF})<68$ into"
+There has been ai growiug database οἳ direct supermassive black hole (SMDIT) mass CAM) estimates from the ceuters of nearby ealactic bulges (e.g...Caham 2008b)..,"There has been a growing database of direct supermassive black hole (SMBH) mass ) estimates from the centers of nearby galactic bulges \citep[e.g.,][]{2008PASA...25..167G}."
+ While the Πες of our current abilities to sieuificautlv expaud this database may have been reached (Batcheldor&Isockemoer2009).. the last decade has seen a wealth of cestimates that have increased the SMDIT catalog from 13 or 26 (Ferrarese&Meritt2000:Gebhardtetal.2000) to ~70 (Caahaim2008b:Πα2008:Caitekinetal.2009)..," While the limits of our current abilities to significantly expand this database may have been reached \citep{2009PASP..121.1245B}, the last decade has seen a wealth of estimates that have increased the SMBH catalog from 13 or 26 \citep{2000ApJ...539L...9F,2000ApJ...539L..13G} to $\sim$ 70 \citep{2008PASA...25..167G,2008MNRAS.386.2242H,2009ApJ...698..198G}."
+ Au intense interest in populating the SMDII database was sparked by observed correlations between aand fundamental properties of their host bulges (6.9..al.2003:Tavineg&Ris200:Pizzellaet 2005).. ," An intense interest in populating the SMBH database was sparked by observed correlations between and fundamental properties of their host bulges \citep[e.g.,][]{1995ARA&A..33..581K,1998AJ....115.2285M,2000ApJ...539L...9F,2000ApJ...539L..13G,2001ApJ...563L..11G,
+2002ApJ...578...90F,2003ApJ...589L..21M,2003MNRAS.341L..44B,2004ApJ...604L..89H,2005ApJ...631..785P}. ."
+These sscaling relations have generated numerous theoretical investigations (e.g...Ciotti&vanAlbada2001:Adamsetal.2003:Cattaneoct2005:Robertsonet2006) and have possibly added valuable limits to evolutionary models (e.g..Heckinanetal.2001:Writhe&Loeb2005:Treuetal.2007:Ciotti 2008)..," These scaling relations have generated numerous theoretical investigations \citep[e.g.,][]{2001ApJ...552L..13C,2003ApJ...591..125A,2005MNRAS.364..407C,
+2006ApJ...641...90R} and have possibly added valuable limits to evolutionary models \citep[e.g.,][]{2004ApJ...613..109H,
+2005ApJ...634..910W,2007ApJ...667..117T,2008arXiv0808.1349C}."
+ The degree to which a SAIBIVs sphere of influence. rà;. is resolved has been used as a quality measure for cestimates (Ferrarese2002:Marconi&IInut2003:Valhuietal.200 D)...," The degree to which a SMBH's sphere of influence, $r_i$, is resolved has been used as a quality measure for estimates \citep{2002ApJ...578...90F,2003ApJ...589L..21M,2004ApJ...602...66V}."
+The only method available toa priori determune ifa cestimate can be made is to assmne r;=GAL/o? (Peebles 1972)... , The only method available toa priori determine if a estimate can be made is to assume $r_i = GM_\bullet/\sigma_\ast^2$ \citep{1972ApJ...178..371P}. .
+"However. to calculate +;in galaxies with known oc... lis estimated using the rrelation given by logAM,=a|Jlog(c./2001ans4), where à—8&1.82 and 9=X719."," However, to calculate $r_i$in galaxies with known , is estimated using the relation given by $\log{M_\bullet} = \alpha + 
+\beta\log{(\sigma_\ast/200\kms)}$, where $\alpha=8.1-8.2$ and $\beta=3.7-4.9$ ."
+" The observed scatter. e, is OL dex (Novalsctal.2006:Caalam&Li2009)..."," The observed scatter, $\epsilon$, is 0.4 dex \citep{2006ApJ...637...96N,2009ApJ...698..812G}."
+ Following this. Figure 1 demonstrates where r; Will be resolved. given a spatial resolution of 071.," Following this, Figure \ref{fig:1} demonstrates where $r_i$ will be resolved, given a spatial resolution of $\Re=0\farcs1$ ."
+ A value of R=(071 is used here as that is the typical FWIIM of the ZZ9T PSF., A value of $\Re=0\farcs1$ is used here as that is the typical FWHM of the PSF.
+ To date ZZST has been responsible for most eestimates., To date has been responsible for most estimates.
+ Ferrarese&Ford(2005.EF05). discussed the linüted abilities ofMST to resolve r;. and CuültelàuetC09) found the irclation to be biased when applying the r; aremuent.," \citet[][FF05]{2005SSRv..116..523F} discussed the limited abilities of to resolve $r_i$, and \citet[][G09]{2009ApJ...698..198G} found the relation to be biased when applying the $r_i$ argument."
+" The iufluence of +; cuts ou the ielation is continued in thisLetter with two siguificaut advances,", The influence of $r_i$ cuts on the relation is continued in this with two significant advances.
+ First. a sample of all galaxies with a kuown (C100 Mpce) is used.," First, a sample of all galaxies with a known $<$ 100 Mpc) is used."
+ It is inportaut to only use these ealaxies as there are no lealaxies at 1l Mpc. for example.," It is important to only use these galaxies as there are no galaxies at 1 Mpc, for example."
+ Second. raucous cestimates are applied to cach galaxy. be.. no galaxy is assed to intriusicallv fall on the ielation.," Second, random estimates are applied to each galaxy, i.e., no galaxy is assumed to intrinsically fall on the relation."
+ This ensures no pre-selection of galaxies that le on the ielation., This ensures no pre-selection of galaxies that lie on the relation.
+ Throughout. a distinction is made between theobserved ielation (published values) aud the irclation (the relation that can be fitted using the data simulated here).," Throughout, a distinction is made between the relation (published values) and the relation (the relation that can be fitted using the data simulated here)."
+ Tn short. these simulations take a galaxy with a kuowu aand distance. assign a Af... then calculate r;;.," In short, these simulations take a galaxy with a known and distance, assign a , then calculate$r_i$ ."
+ The r; argunent isthenapplied: if +; isunresolved the ealaxy is removed fromthe sample (the low mass cut)., The $r_i$ argument isthenapplied; if $r_i$ isunresolved the galaxy is removed fromthe sample (the low mass cut).
+ If theassigned aoeeecnerates a galaxy that liesabove the, If theassigned generates a galaxy that liesabove the
+ If theassigned aoeeecnerates a galaxy that liesabove the , If theassigned generates a galaxy that liesabove the
+ If theassigned aoeeecnerates a galaxy that liesabove the o, If theassigned generates a galaxy that liesabove the
+ If theassigned aoeeecnerates a galaxy that liesabove the ob, If theassigned generates a galaxy that liesabove the
+ If theassigned aoeeecnerates a galaxy that liesabove the obs, If theassigned generates a galaxy that liesabove the
+ If theassigned aoeeecnerates a galaxy that liesabove the obse, If theassigned generates a galaxy that liesabove the
+ If theassigned aoeeecnerates a galaxy that liesabove the obser, If theassigned generates a galaxy that liesabove the
+ If theassigned aoeeecnerates a galaxy that liesabove the obsere, If theassigned generates a galaxy that liesabove the
+ If theassigned aoeeecnerates a galaxy that liesabove the obseree, If theassigned generates a galaxy that liesabove the
+ If theassigned aoeeecnerates a galaxy that liesabove the obsereed, If theassigned generates a galaxy that liesabove the
+clear that strong absorption Is required.,clear that strong absorption is required.
+ These three svstenis also require the absorption to be ionised (or clumpsy. or time variable).," These three systems also require the absorption to be ionised (or clumpy, or time variable)."
+ In. all other systems there is no requirement for the absorbers to be strongly ionisced., In all other systems there is no requirement for the absorbers to be strongly ionised.
+ We find that there is a tendency. for the absorption column density to play oll against the temperature distribution., We find that there is a tendency for the absorption column density to play off against the temperature distribution.
+ As an exampe. we plot confidence contours for the fitec values of ng and à from the spectrum of V603 Acq! in reflig: confidence..," As an example, we plot confidence contours for the fitted values of $n_{\rm H}$ and $\alpha$ from the spectrum of V603 Aql in \\ref{fig:confidence}."
+ Lt can o* seen that the parameters. are COLLOated., It can be seen that the parameters are correlated.
+ Therefore. if t1e contribution [from cool gas is overestimated in our models. this might lead to a systematic overestimate of the absorption column density.," Therefore, if the contribution from cool gas is overestimated in our models, this might lead to a systematic overestimate of the absorption column density."
+ We believe this ds quite likely. since. our. power law temperature distribution does indeed have a Larger contribution [roni cool gas than do cooling Low models (in which the emission measure is weighted by the inverse of the emissivity at that teniperature).," We believe this is quite likely, since our power law temperature distribution does indeed have a larger contribution from cool gas than do cooling flow models (in which the emission measure is weighted by the inverse of the emissivity at that temperature)."
+ For tils reason we advise caution in interpreting our fitted column densities., For this reason we advise caution in interpreting our fitted column densities.
+ In order to asses whether intrinsic absorption is required at all (in systems other than V426 Oph. LS Peg and EL UMa) we turn to the subset of svstems for which we have an independent estimate of the interstellar absorption column densities (indicated in reftable:done)).," In order to asses whether intrinsic absorption is required at all (in systems other than V426 Oph, LS Peg and EI UMa) we turn to the subset of systems for which we have an independent estimate of the interstellar absorption column densities (indicated in \\ref{table:done}) )."
+ Using the F-test statistic. we [find that inclusion of absorption in excess of the interstellar value does not make a significant contribution to our fits for the quiescent cwarf novae (SS Cve and. VW Livi)., Using the F-test statistic we find that inclusion of absorption in excess of the interstellar value does not make a significant contribution to our fits for the quiescent dwarf novae (SS Cyg and VW Hyi).
+ However. the improvement in the fit statisticzs significant (at 00.054 confidence) in both dwarf novae in outburst. (Z Cam and σοςνο) and also in the nova-like LX Vel.," However, the improvement in the fit statistic significant (at $>$ confidence) in both dwarf novae in outburst (Z Cam and SS Cyg) and also in the nova-like IX Vel."
+ Phe only exception to this rule is U Gem. where there is clear evidence for orbital-phase dependent: intrinsic absorption dips in the lhehtceurve (see relscc:orbitalmoctulations ancl reflie:orbital) ).," The only exception to this rule is U Gem, where there is clear evidence for orbital-phase dependent intrinsic absorption dips in the lightcurve (see \\ref{sec:orbitalmodulations} and \\ref{fig:orbital}) )."
+ The GAkkeV line contributed. significantly to our [fits in [our cases: V436 Cen. SS (νο (in outburst). VJ26 Oph and EL UMa.," The keV line contributed significantly to our fits in four cases: V436 Cen, SS Cyg (in outburst), V426 Oph and EI UMa."
+ The best fitting equivalent widths are 17043:50. 95425. 185-zE40. ancl 2003-45eeV. respectively.," The best fitting equivalent widths are $\pm$ 50, $\pm$ 25, $\pm$ 40 and $\pm$ eV respectively."
+ Despite tlose detected Duorescent lines. an X-ray rellection continuum. was not required in order to achieve an statistically acceptab© Lit," Despite these detected fluorescent lines, an X-ray reflection continuum was not required in order to achieve an statistically acceptable fit"
+oomntfs is negative.,points is negative.
+ These sugeest that there are iot a true periodicity in the interval of [200. 250] days.," These suggest that there are not a true periodicity in the interval of [200, 250] days."
+ Tt is difficult to decide αἱ)out the existence of the ooriodicities for &—6 (333 davs) aud &T (286 days) on the basis of above analysis., It is difficult to decide about the existence of the periodicities for $k=6$ (333 days) and $k=7$ (286 days) on the basis of above analysis.
+" The nupication (8) Is vot satisfied for fk=Toand the coucdition (7) is lot satisfied for &=6. although the function ez of he nne seres LS""(t)Sry} is significantly positive for To-—Pag.3:33."," The implication (8) is not satisfied for $k=7$ and the condition (7) is not satisfied for $k=6$ , although the function $c_\tau$ of the time series $\{S^d(t_i)-\overline{S^d(t_i)}\}$ is significantly positive for $\tau=286,333$."
+ The conditious (7) and (8) are satisfiec ork =L(Figure 8)) aud &5., The conditions (7) and (8) are satisfied for $k=4$ (Figure \ref{f8}) ) and $k=5$.
+ Therefore. it is possible o exist the periodicity from the interval oος 500 days.," Therefore, it is possible to exist the periodicity from the interval of $[400,500]$ days."
+ Stuular results were also obtained by Lean(1990) or daily sunspot nunibers and daily sunspot areas., Similar results were also obtained by \citet{lea} for daily sunspot numbers and daily sunspot areas.
+ She considered the meaus of three periodograms of these indexes for data rol. N=31 VCars alc found statistically significant peaks from fιο Interva iH: [100.500]. (sec Lean(1990)... Figure 2).," She considered the means of three periodograms of these indexes for data from $N=31$ years and found statistically significant peaks from the interval of $[400,500]$ (see \citet{lea}, Figure 2)."
+ Isvivova&Solanki(2002) stucied μιlsvot areas from 1576-1999 iid sunspot numbers from 1719-2001 wit1 the help of the wavelet transform., \citet{kri} studied sunspot areas from 1876-1999 and sunspot numbers from 1749-2001 with the help of the wavelet transform.
+ They pointed. οιt that the 151-158-dav period could e the third harimomic of the l-vear (CHrb-dav) period., They pointed out that the 154-158-day period could be the third harmonic of the 1.3-year (475-day) period.
+ Moreover. the both periods fluctuate consideradv with tiue. being «ποσο divine stronger sunspot cycles.," Moreover, the both periods fluctuate considerably with time, being stronger during stronger sunspot cycles."
+ Ticrefore. the wavelet analysis sugecsts a conunon orien of the both periodicities.," Therefore, the wavelet analysis suggests a common origin of the both periodicities."
+ This conclusion coiris the DE method result which indicates that the periodograu peak at 7=151 davs is an alias of the periodicity from the interval of |LOO.500].," This conclusion confirms the DE method result which indicates that the periodogram peak at $\tau =154$ days is an alias of the periodicity from the interval of $[400,500]$ ."
+ Tn order to verify the existence of the periodicity at about 155 davs I cousicer the following time uvEm The value $(f;) is caleuated analogously to Rif;) (sce Sect., In order to verify the existence of the periodicity at about 155 days I consider the following time $\{S(t_i)-\overline{S(t_i)}\}$ – $\{S^d_{n}(t_i)-\overline{S^d_{n}(t_i)}\}$ – $\{S^d_{s}(t_i)-\overline{S^d_{s}(t_i)}\}$ – The value $\overline{S(t_i)}$ is calculated analogously to $\overline{R(t_i)}$ (see Sect.
+ D., 4).
+ The values SU) aud 57(*;) are evaluated from the formula (9)., The values $\overline{S^d_{n}(t_i)}$ and $\overline{S^d_{s}(t_i)}$ are evaluated from the formula (9).
+ In tieupper part of Figure 9 the time series of sunspoareas (5(f;)) of the one rotation time interval from the whole solar disk aid, In theupper part of Figure \ref{f9} the time series of sunspotareas $(S(t_i))$ of the one rotation time interval from the whole solar disk and
+Equations (21)-(24) for the 2-Iaver model contain nine [19ο parameters in addition to the poloidal forcing Ay: equations (29)-(34) lor the 3-Iaver model contain 16 parameters.,Equations (21)-(24) for the 2-layer model contain nine free parameters in addition to the poloidal forcing $A_F$; equations (29)-(34) for the 3-layer model contain 16 parameters.
+ Therefore we must make juclicious choices of parameter values., Therefore we must make judicious choices of parameter values.
+ In making these choices we will be guided by solar conditions as well as the uncertainties in the solar properties (hat deline (hese parameters., In making these choices we will be guided by solar conditions as well as the uncertainties in the solar properties that define these parameters.
+ For example. (he dimensionless frequency w of the top forcing should be approximately the Irequeney of the solar evele. corresponding to a period of 22 vears.," For example, the dimensionless frequency $\omega$ of the top forcing should be approximately the frequency of the solar cycle, corresponding to a period of 22 years."
+ With gc/H7? our frequencyB scale. forB jj=i2X-10cayem?7/see and 4—i2N10Men- this. frequencyB is. 5X10-°/see.(.," With ${\eta}_U/H^2$ our frequency scale, for ${\eta}_U=2X10^{12} 
+cm^2/sec$ and $H=2X10^{10} cm$ this frequency is $5X10^{-9}/sec$ ."
+" TheTl solar larevele frequency frequencyis 9.XIO.ON""/?/seesee, so the dimensionless forcing frequency should FEbe about. 1.8 units."," The solar cycle frequency is $9X10^{-9}/sec$, so the dimensionless forcing frequency should be about 1.8 units."
+ Therefore a lrequency range of 1.5 to 2., Therefore a frequency range of 1.5 to 2.
+ would cover most variability in solar cvcles., would cover most variability in solar cycles.
+ In. all caleulations displaved below. we have chosen a dimensionless frequency w= 1.8.," In all calculations displayed below, we have chosen a dimensionless frequency $\omega$ = 1.8."
+ Specibving the latitudinal wavennuuber of the forcing is more uncertain., Specifying the latitudinal wavemumber of the forcing is more uncertain.
+ The width of the sunspot zone in one hemisphere is about. 30 degrees latitude. or πο.," The width of the sunspot zone in one hemisphere is about 30 degrees latitude, or ${\pi}R/6$."
+ This would be the minimum hall wavelength of the forcing. but that [orcing is seen to be broader in latitude scale than that. due to the dispersal ancl decay of active regions.," This would be the minimum half wavelength of the forcing, but that forcing is seen to be broader in latitude scale than that, due to the dispersal and decay of active regions."
+ Also. we never see surface fields [rom more than 2 sunspot eveles al (he same time. so a more reasonable wavelength might be zA/2. the distance between equator aud pole.," Also, we never see surface fields from more than 2 sunspot cycles at the same time, so a more reasonable wavelength might be ${\pi}R/2$, the distance between equator and pole."
+ Then this wavelength would correspond to a dimensionless wavenumber 7=1.14 units at the surface and κ=1.63 units al the depth of the tachocline., Then this wavelength would correspond to a dimensionless wavenumber $k=1.14$ units at the surface and $k=1.63$ units at the depth of the tachocline.
+ An average value would be about 1.4 units. which is what we use for all caleulations.," An average value would be about 1.4 units, which is what we use for all calculations."
+ As for velocities. the velocity scale jj):/H1 is about 1m/sec. so a (vpical solar meridional [low near the top would be 15 units. and the Iatitudinal differential rotation linear velocity relative to the rotating [rame of about s=70 units.," As for velocities, the velocity scale ${\eta}_U/H$ is about 1m/sec, so a typical solar meridional flow near the top would be 15 units, and the latitudinal differential rotation linear velocity relative to the rotating frame of about $s=70$ units."
+ We will use these dimensionless values to guide our choices of parameter ranges to survey., We will use these dimensionless values to guide our choices of parameter ranges to survey.
+ For some purposes the choice of jj;=10520?/sec may be too high., For some purposes the choice of ${\eta}_U=10^{12}cm^2/sec$ may be too high.
+ Reducing it by a [actor of ten means (hat all dimensionless solar frequencies and velocities are increased by a factor of ten. but dimensionless wavenumbers remain the sanie.," Reducing it by a factor of ten means that all dimensionless solar frequencies and velocities are increased by a factor of ten, but dimensionless wavenumbers remain the same."
+ The l-laver results given above could give us guidance about where to look in parameter space for resonance when there are more lavers than one., The 1-layer results given above could give us guidance about where to look in parameter space for resonance when there are more layers than one.
+ As in that case. we might expect that for resonance to occur in a laver. we must havethe phase speed of the forcing at the," As in that case, we might expect that for resonance to occur in a layer, we must havethe phase speed of the forcing at the"
+in fact requirecl to reconcile some observations with theory (Dahle 2003. Dalal Ixochanek 2002). although this conclusion iw not been universally accepted (Sand 2003: Schechter Wambseanss 2002: Evans Witt. 2003).,"in fact required to reconcile some observations with theory (Dahle 2003, Dalal Kochanek 2002), although this conclusion has not been universally accepted (Sand 2003; Schechter Wambsganss 2002; Evans Witt 2003)."
+ 1. rowever. the lensing detection of halo substructure correc and the overabundant satellite population really cloes exist. it is imperative to understand. the orbital evolution of hese objects and their deviation from the background dark matter cistribution.," If, however, the lensing detection of halo substructure correct and the overabundant satellite population really does exist, it is imperative to understand the orbital evolution of these objects and their deviation from the background dark matter distribution."
+ The work deseribed here. focuses upon a set of numerical simulations of structure formation within the concordance model. analysing in detail the temporal anc spatial properties of satellite galaxies residing within hos dark matter halos.," The work described here focuses upon a set of numerical simulations of structure formation within the concordance model, analysing in detail the temporal and spatial properties of satellite galaxies residing within host dark matter halos."
+ To date. typical satellite properties such as orbital parameters ancl mass loss under the inlluence of the host halo have primarily been investigated using potentials for the dark matter host halo (e.g. Johnston 1996: ανασα 2003)," To date, typical satellite properties such as orbital parameters and mass loss under the influence of the host halo have primarily been investigated using potentials for the dark matter host halo (e.g. Johnston 1996; Hayashi 2003)."
+ We stress that cach of these studies have provided: invaluable insights into the physical processes involved in satellite disruption: our goals was to augment these studies by relaxing the assumption of a static host potential. in deference to the Fact that realistic dark matter halos are not necessarily axis-svmametric.," We stress that each of these studies have provided invaluable insights into the physical processes involved in satellite disruption; our goals was to augment these studies by relaxing the assumption of a static host potential, in deference to the fact that realistic dark matter halos are not necessarily axis-symmetric."
+ Halos constantly erow in mass through slow accretion and violen niergers. possessing rather triaxial shapes (Warren et. al.," Halos constantly grow in mass through slow accretion and violent mergers, possessing rather triaxial shapes (Warren et al."
+ 1992)., 1992).
+ While a κοconsistent cosmological modeling of both rosts and satellites has long been recognised. as optimal. he required. mass ancl force. resolution can be dillicult to accommocate (hence the use of static host potentials in most orevious stucies).," While a self-consistent cosmological modeling of both hosts and satellites has long been recognised as optimal, the required mass and force resolution can be difficult to accommodate (hence the use of static host potentials in most previous studies)."
+ The [first fully self-consistent) simulations targeting he subject were performed. bv. Tormen. (1997). and lormen (1998)., The first fully self-consistent simulations targeting the subject were performed by Tormen (1997) and Tormen (1998).
+ Both studies were landmark etlorts. out. [lacked the temporal. spatial. and mass resolution necessary to explore a wide range of environmental effects.," Both studies were landmark efforts, but lacked the temporal, spatial, and mass resolution necessary to explore a wide range of environmental effects."
+ Unable to follow the satellite distribution within the hosts virial raclius. satellites were instead. tracked only up to and including the point of “accretion”.," Unable to follow the satellite distribution within the host's virial radius, satellites were instead tracked only up to and including the point of “accretion”."
+ “Phis allowed an analysis of the infall pattern. rather than the orbital evolution of the satellites.," This allowed an analysis of the infall pattern, rather than the orbital evolution of the satellites."
+ Chigna (1998) also investigated the cynamics of satellite galaxies in live dark matter host halos., Ghigna (1998) also investigated the dynamics of satellite galaxies in live dark matter host halos.
+ Although ereath increasing the mass and spatial resolution. they still lacked the temporal resolution to explicitly track the satellite orbits.," Although greatly increasing the mass and spatial resolution, they still lacked the temporal resolution to explicitly track the satellite orbits."
+ Instead. the orbits were approximated using a spherical static potential.," Instead, the orbits were approximated using a spherical static potential."
+ More recently. Talloni (2008) usecl ssimulations coupled with semi-analytical tools to explore the evolution of dark matter satellites inside more massive halos.," More recently, Taffoni (2003) used simulations coupled with semi-analytical tools to explore the evolution of dark matter satellites inside more massive halos."
+ However. they focus their elforts on. the interplay between dynamical friction ancl tidal mass loss in determining the final fate of the satellites.," However, they focus their efforts on the interplay between dynamical friction and tidal mass loss in determining the final fate of the satellites."
+ Ixravtsoy (2004) also mainly concentrate on the mass loss history of satellites using fully sellconsistent cosmological ssimulations., Kravtsov (2004) also mainly concentrate on the mass loss history of satellites using fully self-consistent cosmological simulations.
+ In this paper we investigate the evolution of substructure and the orbital parameters of satellites using high spatial. mass.end temporal resolution.," In this paper we investigate the evolution of substructure and the orbital parameters of satellites using high spatial, mass, temporal resolution."
+ As outlined in Paper E (Gill. ποσο Cibson 2004: hereafter GKGI). our suite of simulations has the required. resolution to. follow the satellites even within the very central regions of the host potential (25 of the virial radius) and the time resolution to resolve the satellite dynamics with excellent accuracy (GM2170 Alves).," As outlined in Paper I (Gill, Knebe Gibson 2004; hereafter ), our suite of simulations has the required resolution to follow the satellites even within the very central regions of the host potential $\geq$ of the virial radius) and the time resolution to resolve the satellite dynamics with excellent accuracy $\Delta t \approx$ 170 Myrs)."
+ The outline of the paper is as follows., The outline of the paper is as follows.
+ Section ?? provides a description. of the cosmological simulations emploved.," Section \ref{Computation}
+ provides a description of the cosmological simulations employed."
+ The analysis of the host halo and environment can be found in Section 3. with the satellite orbital parameters presented in Section ??..," The analysis of the host halo and environment can be found in Section \ref{HaloAnalysis}, with the satellite orbital parameters presented in Section \ref{SatAnalysis}."
+ We then investigate the kinematic properties of the dark matter halos and satellites in Section ??.., We then investigate the kinematic properties of the dark matter halos and satellites in Section \ref{velbias}.
+ We finish with our summary and conclusions in Section 6.., We finish with our summary and conclusions in Section \ref{Conclusions}.
+ Our analvsis is based upon a suite of eight high-resolution ssimulations generated using the publicly available adaptive mesh refinement code ((Ixnebe. Green Binney 2001).," Our analysis is based upon a suite of eight high-resolution simulations generated using the publicly available adaptive mesh refinement code (Knebe, Green Binney 2001)."
+ rreaches hieh [force resolution by refining all high-density regions with an automated. refinement algorithm., reaches high force resolution by refining all high-density regions with an automated refinement algorithm.
+ The refinements are recursive: refined regions can also be refined. cach subsequent erid level having cells that are half the size of the cells in the previous level.," The refinements are recursive: refined regions can also be refined, each subsequent grid level having cells that are half the size of the cells in the previous level."
+ This creates a hierarchy. of refinement meshes of cillerent resolutions covering regions of interest., This creates a hierarchy of refinement meshes of different resolutions covering regions of interest.
+ The refinement is done cell-by-cell (individual cells can be refined or de-refined) and meshes are not constrained to have a rectangular (or any other) shape., The refinement is done cell-by-cell (individual cells can be refined or de-refined) and meshes are not constrained to have a rectangular (or any other) shape.
+ The criterion for a cell is simply the number of particles within that cell., The criterion for a cell is simply the number of particles within that cell.
+ A detailed study of the appropriate choice for this number as well as more details about the particulars of the code can be found in Ixnebe (2001)., A detailed study of the appropriate choice for this number as well as more details about the particulars of the code can be found in Knebe (2001).
+ The force resolution is determined by the finest refinement level reached and corresponds to 22h Που 10 simulations presented. here., The force resolution is determined by the finest refinement level reached and corresponds to $\approx$ for the simulations presented here.
+ Phe mass of an individual low-mass particle is ny=1.6.1075. aand the halos are resolved with on the order of millions of these particles.," The mass of an individual low-mass particle is $m_p =
+1.6 \times 10^{8}$ and the halos are resolved with on the order of millions of these particles."
+ In order to investigate the evolution of satellite galaxies and their debris high temporal sampling of the outputs was necessary., In order to investigate the evolution of satellite galaxies and their debris high temporal sampling of the outputs was necessary.
+ From 2=2.5 to 2=0.5 we have 17 equally spaced outputs with Afe O.35Cvrs., From $z=2.5$ to $z=0.5$ we have 17 equally spaced outputs with $\Delta t \approx 0.35$ Gyrs.
+ From 2=0.5 tos—0 we have 30 outputs spaced at Afzz0 TGvrs.," From $z=0.5$ to $z=0$ we have 30 outputs spaced at $\Delta t
+\approx 0.17$ Gyrs."
+ For further details please refer toGKGL., For further details please refer to.
+. For cach of our 376 outputs the satellite galaxies were initially located: using Haalo-Fiinder (MHF)) (GKGL))., For each of our 376 outputs the satellite galaxies were initially located using inder ) ).
+ This provided us with a list o£ all satellites and their internal properties at cach individual redshift; under consideration., This provided us with a list of all satellites and their internal properties at each individual redshift under consideration.
+ Llowever. as we are more interested in orbital information we performed a detailed time analysis of those satellites that were within two times the virial radius of the host halo at its formation time (using MLAPM--LIalo-Fracker: MHT)).," However, as we are more interested in orbital information we performed a detailed time analysis of those satellites that were within two times the virial radius of the host halo at its formation time (using -Halo-Tracker: )."
+ For à, For a
+spectrum. (,spectrum. (
+This is rather obvious. as the two spectra aro by coustruction very similar. and the wo-poiut correlation function. which is really what cuters iuto these calculations. is also correspoudinely simiar.),"This is rather obvious, as the two spectra are by construction very similar, and the two-point correlation function, which is really what enters into these calculations, is also correspondingly similar.)"
+" Second. we note that one can see clirectly from the figures shown in Curzadyan aud Penrose’s paper that their claimed ""ACDM spectrum is still fied: It is cithcult to sav exactly what went wrone du the generation of t18ο updated simulations."," Second, we note that one can see directly from the figures shown in Gurzadyan and Penrose's paper that their claimed $\Lambda$ CDM” spectrum is still flawed: It is difficult to say exactly what went wrong in the generation of these updated simulations."
+ At least they have a nou-flat power svectrmun. Which is a clear nuprovenient over tιο first version.," At least they have a non-flat power spectrum, which is a clear improvement over the first version."
+ Still. it is also clear that the preseut calus are still not correct. and the sale cntieigus that were presented by Welus (2011).. Moss.Scot&Zibin(2011) and Tajian(2010) still apply: When making claims simular to those of Churzadyiui&Peurose(POLL). itis to construct the underlying sinitlations with absolute data fidelity.," Still, it is also clear that the present claims are still not correct, and the same criticisms that were presented by \citet{wehus:2011}, \citet{moss:2011} and \citet{hajian:2010} still apply: When making claims similar to those of \citet{gurzadyan:2011a}, it is to construct the underlying simulations with absolute data fidelity."
+ Wwule he physical importance of Curzadw aud Penrose’s recent claims in our opiuionu are marginal at best. we do believe that there are some interesting poiuts in terius of science sociology aud the currently accepted reterecing process.," While the physical importance of Gurzadyan and Penrose's recent claims in our opinion are marginal at best, we do believe that there are some interesting points in terms of science sociology and the currently accepted refereeing process."
+ While performing the first reanalysis of Cazadvan auc Peurose’s claims. we read through most of the papers cited in their original paper. trviug to undoerstaud the ]ckeround for their clans.," While performing the first reanalysis of Gurzadyan and Penrose's claims, we read through most of the papers cited in their original paper, trying to understand the background for their claims."
+ Iu particular. one apparently ceutral line of reasoning of Curzadveui&Peurose(2010) was based on the notion of the “Noliuogorov statistic’. as introduced by Curzadyviui&Nocheανα(2008):παναιetal.(2009) aud references there3n.," In particular, one apparently central line of reasoning of \citet{gurzadyan:2010} was based on the notion of the “Kolmogorov statistic”, as introduced by \citet{gurzadyan:2008, gurzadyan:2009}
+ and references therein."
+ This statistic nieasures the degree of “randomness” within a set of stochastic variables., This statistic measures the degree of “randomness” within a set of stochastic variables.
+ Tn particular. Comzadyzuetal.(2011) applied this statistic to the inall disks iu the WMATPD sky maps. aud measured the degree of rand:muness within each disk.," In particular, \citet{gurzadyan:2011b} applied this statistic to the small disks in the WMAP sky maps, and measured the degree of randomness within each disk."
+" The main conclusion drawn from this work was tiit only of the signal was ""raidomr. while of the signal was ""non-randonr."," The main conclusion drawn from this work was that only of the signal was “random”, while of the signal was “non-random”."
+ When reading these papers. if seems clear to us that (ανα et οςconfuse randomness with correlation: While the CAIB fehis (most Likely) a raudon field. it is wucorrelated.," When reading these papers, it seems clear to us that Gurzadyan et confuse randomness with correlation: While the CMB field is (most likely) a random field, it is uncorrelated."
+ lustead. the CAIB field is a s1100thi field on scales com]xrable with the iustriuneutal beam. aud it has a well-defined non-fiat power spectrum.," Instead, the CMB field is a smooth field on scales comparable with the instrumental beam, and it has a well-defined non-flat power spectrum."
+ Thus. the veal-space corrclations are strong.," Thus, the real-space correlations are strong."
+ Of course. the iustruineutal is virtually uncorrelated. aud so there are indeed two coim»nents here. oue correlated iud one uncorrelated.," Of course, the instrumental is virtually uncorrelated, and so there are indeed two components here, one correlated and one uncorrelated."
+ But neither is non-riuidoni., But neither is non-random.
+ The interesting wt of this story. though. is the act that at least five papers ou this very topic have (en accepted au published bv the reputable (ancl refereed) journal “Astronomy and Astroplivsics”.," The interesting part of this story, though, is the fact that at least five papers on this very topic have been accepted and published by the reputable (and refereed) journal “Astronomy and Astrophysics”."
+" One of these papers (caHed ""A weakly raudoim universe?"")", One of these papers (called “A weakly random universe?”)
+" was even published asa Letter. with au abstract stating hat ""Deriviug the ορΊσα]. Kohuogorov's function iu he Wilkinson Microwave Anisotropy Probes maps. we obtain the fraction of the random signal to be about 20 ver cent. ie. the «o»nological sky is a weakly random ouc."," was even published as a Letter, with an abstract stating that “Deriving the empirical Kolmogorov's function in the Wilkinson Microwave Anisotropy Probe's maps, we obtain the fraction of the random signal to be about 20 per cent, i.e. the cosmological sky is a weakly random one.”"
+ These are tituv extraordinary claus. and iu our view have no roo in reality.," These are truly extraordinary claims, and in our view have no root in reality."
+ Further. these claims are rot irrelevant. as clearly demoustrated by the most recent developeits concerning the concentre rings: They have. at the ve3v least implicitly. led to an excessive aout of mubhlicity im the eeneral public. potentially damaging the pudie. percexiou of cosimnologists in a wider seuse.," Further, these claims are not irrelevant, as clearly demonstrated by the most recent developments concerning the concentric rings: They have, at the very least implicitly, led to an excessive amount of publicity in the general public, potentially damaging the public perception of cosmologists in a wider sense."
+ In ou view. this is a clear demoustratiou of the poteutial weeusnesses of the established refereciug processes: Àareinal. or even plain wrong. work cau be published die to au unattenive referee.," In our view, this is a clear demonstration of the potential weaknesses of the established refereeing processes: Marginal, or even plain wrong, work can be published due to an unattentive referee."
+ Contrary to this. it is interesting to note the reaction that came after the first Cowzadvan aud Penrose paper was put on tje arXiv in November 2011: In ouly a matter of weeks. three iudepeudenu eroups refuted the original claim.," Contrary to this, it is interesting to note the reaction that came after the first Gurzadyan and Penrose paper was put on the arXiv in November 2011: In only a matter of weeks, three independent groups refuted the original claim."
+ Of course. this reaction was hugelv triggered by the massive licelia atteutiou that the original story cot. and which nujt papers will never experience.," Of course, this reaction was largely triggered by the massive media attention that the original story got, and which most papers will never experience."
+ Nevertheless. we believe tliat this particular case is a eood demonstration of the power of community review outside the established joirhals: The open comunity can be a far more cfficicu reviewer than a somewhat arbitrary referee appointed by a given journal.," Nevertheless, we believe that this particular case is a good demonstration of the power of community review outside the established journals: The open community can be a far more efficient reviewer than a somewhat arbitrary referee appointed by a given journal."
+of the expansion velocity.,of the expansion velocity.
+" The relation between the diffusion time of our objects and SN 1987A is t4(06V)=0.90tg(87A) and t4(06au)=0.85ta(87A), as measured from the time of the bolometric peaks, see top panel in Fig. [I4]."," The relation between the diffusion time of our objects and SN 1987A is $t_d(06V)=0.90~t_d(87A)$ and $t_d(06au)=0.85~t_d(87A)$, as measured from the time of the bolometric peaks, see top panel in Fig. \ref{bolo}."
+" If we now assume the same mean opacity for each SN, and use an average ratio between the expansion velocities as measured from the at 45169, (~ 1.4 for aand ~ 1.7 for see bottom panel in Fig.[I4)),"," If we now assume the same mean opacity for each SN, and use an average ratio between the expansion velocities as measured from the at $\lambda$ 5169, $\sim$ 1.4 for and $\sim$ 1.7 for see bottom panel in Fig. \ref{bolo}) ),"
+ the ejecta mass for both objects is found to be M.;~20 Mc with corresponding (EοςM.jv?) kinetic energies of a few foe., the ejecta mass for both objects is found to be $M_{ej}\sim20$ $\Msun$ with corresponding $E\propto M_{ej}v^2$ ) kinetic energies of a few foe.
+ These simple estimates give a first hint on the nature of the progenitors., These simple estimates give a first hint on the nature of the progenitors.
+ From the bolometric light curve one can also constrain the amount of ssynthesized in the explosion., From the bolometric light curve one can also constrain the amount of synthesized in the explosion.
+" The daughter decay product of is °°Co, and it is the decay of ?9Co to °°Fe that powers the late-time light curve."," The daughter decay product of is $^{56}$ Co, and it is the decay of $^{56}$ Co to $^{56}$ Fe that powers the late-time light curve."
+" Unfortunately our photometric coverage of SN 2006au does not extend beyond ~ 105 days after the explosion, so only an upper limit on the mmass can be estimated."," Unfortunately our photometric coverage of SN 2006au does not extend beyond $\sim$ 105 days after the explosion, so only an upper limit on the mass can be estimated."
+" This is done by assuming the last photometric epoch belongs to the linear decay phase, where LοςMssy;e~*/7."," This is done by assuming the last photometric epoch belongs to the linear decay phase, where $L\propto M_{\Ni}e^{-t/\tau_{}}$."
+ We thus obtain Msex;(06au)<0.073Mo., We thus obtain $M_{\Ni}(06au)\leq 0.073~\Msun$.
+" In contrast, the photometric coverage of SN 2006V (up to 150 days after the explosion) is sufficient to directly measure the mmass from the linear decay phase."," In contrast, the photometric coverage of SN 2006V (up to 150 days after the explosion) is sufficient to directly measure the mass from the linear decay phase."
+" In doing so, we compute Msew;(06V)=0.127+0.010Mo."," In doing so, we compute $M_{\Ni}(06V)=0.127\pm 0.010~\Msun$."
+" Following the plateau luminosity relation presented by FasePopes 1,οςE56y and scaling with the of SN (993),1987A, R((87A)LP?=gara,3x1013 cm progenitor=43 Ro (Woosley)2006V a rough estimate of the radius of SNe {1988),,and 2006au can be inferred."," Following the plateau luminosity relation presented by \citet{popov93}, $L\propto E^{5/6}M_{ej}^{-1/2}R^{2/3}$ , and scaling with the radius of SN 1987A, $R(87A)=3\times 10^{12}$ cm $=43$ $\Rsun$ \citep{woosley88}, a rough estimate of the progenitor radius of SNe 2006V and 2006au can be inferred."
+" We note that this approach has been followed by (2011) for estimating the progenitor radius of SN 2000cb, although (1993) developed the plateau luminosity relation for SNe whose emission was nonradioactive."," We note that this approach has been followed by \citet{kleiser11} for estimating the progenitor radius of SN 2000cb, although \citet{popov93} developed the plateau luminosity relation for SNe whose emission was nonradioactive."
+" Such a simple scaling implies radii of <50R for both objects, which clearly suggests a very compact progenitor for our objects."," Such a simple scaling implies radii of $\lesssim 50~\Rsun$ for both objects, which clearly suggests a very compact progenitor for our objects."
+" To confirm these estimates of the progenitor and explosion parameters, we turn to the semi-analytic model of (1992)."," To confirm these estimates of the progenitor and explosion parameters, we turn to the semi-analytic model of \citet{impop92}."
+". This model includes cooling and recombination, and has been shown to provide a good fit to the bolometric light curve of SN 1987A [[992)."," This model includes cooling and recombination, and has been shown to provide a good fit to the bolometric light curve of SN 1987A \citep{impop92}."
+". As this model is not applicable to the earliest phases, it is fit only to the bolometric light curves of SNe 1987A, 2006V and 2006au at epochs after 40 days past explosion."," As this model is not applicable to the earliest phases, it is fit only to the bolometric light curves of SNe 1987A, 2006V and 2006au at epochs after 40 days past explosion."
+" Adopting a mean opacity &—0.34 cm?g~!, and the same exponential distribution of as used by (1992) to fit the bolometric light curve of SN 1987A, M.;, E, Msew; and R are estimated for all three objects."," Adopting a mean opacity $\kappa=0.34$ $^2$ $^{-1}$, and the same exponential distribution of as used by \citet{impop92} to fit the bolometric light curve of SN 1987A, $M_{ej}$, $E$, $M_{\Ni}$ and $R$ are estimated for all three objects."
+ In the case of wwe assume a Msew; that is equal to the upper limit estimated from the last photometric epoch., In the case of we assume a $M_{\Ni}$ that is equal to the upper limit estimated from the last photometric epoch.
+" On the other hand, for this SN we can provide additional constraints, in particular on the progenitor radius, by using the initial dip in the light curve."," On the other hand, for this SN we can provide additional constraints, in particular on the progenitor radius, by using the initial dip in the light curve."
+" This was done for SN 19874, whose light curves also contain a shock break-out cooling tail."," This was done for SN 1987A, whose light curves also contain a shock break-out cooling tail."
+" provides an analytical expression for the luminosity of the early-time light curve of SN 1987A and similar SNe, in terms of E,Tom and R."," \citet{chevalier92} provides an analytical expression for the luminosity of the early-time light curve of SN 1987A and similar SNe, in terms of $E$, $M_{ej}$ and $R$."
+ The luminosity function given by [Chevalier](1992d is L=3.08x108589!Mgnt?(F1/1.35)91Raorot9 erg s1., The luminosity function given by \citet{chevalier92} is $L=3.08\times10^{43}E_{51}^{0.91}M_{16\Msun}^{-0.40}(F1/1.35)^{-0.17}R_{30\Rsun}t^{-0.34}$ erg $^{-1}$.
+ F1 is the factor by which each gas element increases in velocity from t=0 to very late times., $F1$ is the factor by which each gas element increases in velocity from $t=0$ to very late times.
+ Following [Chevalier] we adopt F1=1.35.," Following \citet{chevalier92}
+ we adopt $F1=1.35$."
+" As this expression overestimates the luminosity of SN 19874 by a factor of 2, we scale it by this factor in order to fit the early epochs for both SNe 1987A and 2006au."," As this expression overestimates the luminosity of SN 1987A by a factor of 2, we scale it by this factor in order to fit the early epochs for both SNe 1987A and 2006au."
+ The best simultaneous fits to the bolometric light curves with the [[mshennikPopov](1992) model and the analytic expression& are shown in Fig., The best simultaneous fits to the bolometric light curves with the \citet{impop92} model and the \citet{chevalier92} analytic expression are shown in Fig.
+[I4] (top panel) as dashed lines., \ref{bolo} (top panel) as dashed lines.
+" The (1907) model allows us to also constrain the ionization temperature, 7;55,, which strongly affects the light curve shape, by fitting the effective temperature before B,,4,; (dashed lines in the middle panel of Fig. [[4))."," The \citet{impop92} model allows us to also constrain the ionization temperature, $T_{ion}$, which strongly affects the light curve shape, by fitting the effective temperature before $B_{max}$ (dashed lines in the middle panel of Fig. \ref{bolo}) )."
+ During the recombination phase the photospheric velocity estimated from A5169 has also been fit in order to better constrain the energy and mass (dashed lines in the bottom panel of Fig. [[4))., During the recombination phase the photospheric velocity estimated from $\lambda$ 5169 has also been fit in order to better constrain the energy and mass (dashed lines in the bottom panel of Fig. \ref{bolo}) ).
+" Our estimates for aare as follows: M.;(06V)=17.0 Mo, E(06V)=2.410?! erg, Mss;(06V)=0.127 Mo and R(06V)=75 Ro."," Our estimates for are as follows: $M_{ej}(06V)=17.0$ $\Msun$, $E(06V)=2.4\times10^{51}$ erg, $M_{\Ni}(06V)= 0.127$ $\Msun$ and $R(06V)=75$ $\Rsun$."
+" The model for ggives: Mej(06au)=19.3 Mo, E(06au)=3.2x10°! erg, Msewi;(06au)=0.073 Mo and R(06au)=90 Ro."," The model for gives: $M_{ej}(06au)=19.3$ $\Msun$, $E(06au)=3.2\times10^{51}$ erg, $M_{\Ni}(06au)= 0.073$ $\Msun$ and $R(06au)=90$ $\Rsun$."
+" The parameters for SN 1987A are: M.;(87A)=11.8 Mo, E(87A)=1.1x10°! erg, Μεονι(δΤΑ)=0.078 Mo and R(87A)=33 Ro."," The parameters for SN 1987A are: $M_{ej}(87A)=11.8$ $\Msun$, $E(87A)=1.1\times10^{51}$ erg, $M_{\Ni}(87A)= 0.078$ $\Msun$ and $R(87A)=33$ $\Rsun$."
+ The latter is in reasonable agreement with the values inferred from hydrodynamical simulations (Blinnikovetal.[2000)., The latter is in reasonable agreement with the values inferred from hydrodynamical simulations \citep{blinnikov00}.
+. Clearly the adopted semi-analytic model relies on significant simplifications and therefore one can not expect an exact fit to the data., Clearly the adopted semi-analytic model relies on significant simplifications and therefore one can not expect an exact fit to the data.
+" However, it does provide a set of reasonable physical parameters."," However, it does provide a set of reasonable physical parameters."
+" The mass and energy estimates of our objects are somewhat lower than those obtained from the simple scaling relations, and at the same time, are larger than those obtained for SN 1987A. The progenitor radii estimates from the semi-analytic models are larger than those estimated from the simple scalings, and somewhat larger than what is computed for the progenitor of SN 1987A. Nevertheless, our estimates onthe radii suggests that the progenitors of SNe 2006V and 2006au were compact stars."," The mass and energy estimates of our objects are somewhat lower than those obtained from the simple scaling relations, and at the same time, are larger than those obtained for SN 1987A. The progenitor radii estimates from the semi-analytic models are larger than those estimated from the simple scalings, and somewhat larger than what is computed for the progenitor of SN 1987A. Nevertheless, our estimates onthe radii suggests that the progenitors of SNe 2006V and 2006au were compact stars."
+" Note that if the same elapsed time between epochs of explosion and maximum for SNe 2006V and 2006au had been assumed as for SN 1987A, larger estimates of mass and energy would have been obtained, especially for ((30 Moa)."," Note that if the same elapsed time between epochs of explosion and maximum for SNe 2006V and 2006au had been assumed as for SN 1987A, larger estimates of mass and energy would have been obtained, especially for $\sim30$ $\Msun$ )."
+" The Msew; estimate would also have been slightly enhanced, while the radius estimate would not change significantly."," The $M_{\Ni}$ estimate would also have been slightly enhanced, while the radius estimate would not change significantly."
+ Even in this case the compact star scenario would be favoured., Even in this case the compact star scenario would be favoured.
+(in contrast (o ~ a lew keV in our case).,(in contrast to $\sim$ a few keV in our case).
+ For comparable £. calculated. from the 1.1000 Iv flux the number of photons above the Fe Ix edge is (hus much smaller (han in our case.," For comparable $\xi$, calculated from the 1—1000 Ry flux the number of photons above the Fe K edge is thus much smaller than in our case."
+" ""Therefore. heavy metals are almost fully ionized in most of our simulations. leading (ο weak or absent line features and a relatively. weak Compton reflection component."," Therefore, heavy metals are almost fully ionized in most of our simulations, leading to weak or absent line features and a relatively weak Compton reflection component."
+ Dallantvne.Turner&Blaes(2004) explored reflections from a hot corona off of an inhomogeneous disk and found the spectrum can differ sienilicantly [rom reflection off a homogeneous disk., \citet{btb04} explored reflections from a hot corona off of an inhomogeneous disk and found the spectrum can differ significantly from reflection off a homogeneous disk.
+ They also had. stronger reflection features than our results. due to lower ionization parameters. and no ~9keV feature.," They also had stronger reflection features than our results, due to lower ionization parameters, and no $\sim 9\ \kev$ feature."
+ Feng&lxaaret(2005) and Stobbart.Roberts&Wilms(2006) fit ULAs’ N-ray spectra with various models: when fit with a MCDDD or MCDDD and low-energy. blackbocly. (heir fits had inner disk temperatures similar to ours.," \citet{fk05} and \citet{srw06} fit ULXs' X-ray spectra with various models; when fit with a MCDBB or MCDBB and low-energy blackbody, their fits had inner disk temperatures similar to ours."
+ However. we find (hat our model is unable to explain a soft-excess.," However, we find that our model is unable to explain a soft-excess."
+ We note that no 9keV. feature has been detected in (he spectra ol ULXs. although features at other energies have been detected (Strohmaver&Mushotzky2003:Agrawal&Misra2006:ontοἱal.2004:Dewangan.GriffithsRao— 2006).. although these features could originate [rom a wind rather than reflection features in the disk.," We note that no $\sim 9 \kev$ feature has been detected in the spectra of ULXs, although features at other energies have been detected \citep{sm03,am06,ketal04,dgr06}, although these features could originate from a wind rather than reflection features in the disk."
+ ULX observations above ~10keV by. or (he next generation of hard X-ray imaging insirumenis might be able to detect the hard. power-laws we predict. although we realize such observations would be difficult.," ULX observations above $\sim 10 \kev$ by, or the next generation of hard X-ray imaging instruments might be able to detect the hard power-laws we predict, although we realize such observations would be difficult."
+ observations of (wo ULXNs in NGC 1313 did nol detect any component above 10 keV above the background (Mizunoetal...2007).. although their MCDDD + power-law [its do seem to agree will our spectra. and the variability of (hese sources is much less than observed in Galactic black hole eandidates.," observations of two ULXs in NGC 1313 did not detect any component above 10 keV above the background \citep{mizuno07}, although their MCDBB + power-law fits do seem to agree with our spectra, and the variability of these sources is much less than observed in Galactic black hole candidates."
+ Unfortunately. our simulations do not reproduce a soft excess as has been observed in many ULXs.," Unfortunately, our simulations do not reproduce a soft excess as has been observed in many ULXs."
+ It is possible that the soft excess originates [rom Compton downscattering of radiation bv wind (e.g..Degelman.|2001).. aud so a lack of its production in our simulations does not disprove the photon bubble model for ULXs.," It is possible that the soft excess originates from Compton downscattering of radiation by wind \citep[e.g.,][]{begelman01}, and so a lack of its production in our simulations does not disprove the photon bubble model for ULXs."
+ If the excess is explained. our model could explain ULXs well fii with a ~1 keV blackbody.," If the excess is explained, our model could explain ULXs well fit with a $\sim 1$ keV blackbody."
+ Higher energy observations (210 keV) could determine this: if they found a stable hard. X-ray power-law D~ 202.4 for a long period of Gime. this woulcl be evidence for the photon bubble model in ULXs. due to the fact Chat a lower mass compact object would have less variabilitv (IxXalogeraetal.," Higher energy observations $\ga10$ keV) could determine this; if they found a stable hard X-ray power-law $\Gamma \sim$ 2—2.4 for a long period of time, this would be evidence for the photon bubble model in ULXs, due to the fact that a lower mass compact object would have less variability \citep{kalogera04}."
+2004)... A ~9keV. [eatiue could also be considered a signature of photon bubbles in ULNs: it has not been found in any other simulation of accreting black holes., A $\sim 9 \kev$ feature could also be considered a signature of photon bubbles in ULXs; it has not been found in any other simulation of accreting black holes.
+ We note that our spectra are similar to the verv high state of N-ray. binaries such as GX 339-4 (Bellonietal.2006) and GRO J1655-40 (Saitoetal...2006).. whieh have similar photon indices as found in our simulations.," We note that our spectra are similar to the very high state of X-ray binaries such as GX 339-4 \citep{belloni06} and GRO J1655-40 \citep{saito06}, which have similar photon indices as found in our simulations."
+ For smaller accretion rates. the photon bubble model may be a viable model to explain the very. high state of X-ray binaries.," For smaller accretion rates, the photon bubble model may be a viable model to explain the very high state of X-ray binaries."
+where ris the radius-vector within the satellite. Ris the distance to the perturber. n=R/R. i—ειMfdluR. and MGR) is the enclosed mass of the NEW ος]: Figure 3 shows that the approximate tidal force calculated iun this ΠΕ is quite accurate. especially near the maxi of the tidal force.,"where $\rr$ is the radius-vector within the satellite, $R$ is the distance to the perturber, ${\bf n} \equiv \RR/R$ , $\acute{\mu}
+\equiv d\ln{M}/d\ln{R}$, and $M(R)$ is the enclosed mass of the NFW model: Figure \ref{fig:tr1} shows that the approximate tidal force calculated in this manner is quite accurate, especially near the maximum of the tidal force."
+ Althoush the tidal force along the satellite trajectory varies rapidly with time. most of the tidal heating of stars and dark matter particles occurs near the stroug peaks of the tidal force.," Although the tidal force along the satellite trajectory varies rapidly with time, most of the tidal heating of stars and dark matter particles occurs near the strong peaks of the tidal force."
+ Each of these tidal peaks can be considered as an independent tidal shock(??)., Each of these tidal peaks can be considered as an independent tidal shock.
+. The amount of tidal heating. such as the increase of the velocity dispersion. is proportional to the integral over the peals of tidal force: where the stun extends over all components of the tidal teusor. a.)=fey.2}.," The amount of tidal heating, such as the increase of the velocity dispersion, is proportional to the integral over the peak of tidal force: where the sum extends over all components of the tidal tensor, $\alpha,\beta=\{x,y,z\}$."
+ The last factor is the correction for the conservation of adiabatic iuvariauts of stellar orbits during the tidal shock?)., The last factor is the correction for the conservation of adiabatic invariants of stellar orbits during the tidal shock.
+". Tere τι is the effective duratiou of peak at time t,,. aud £444 is the dynamical time of the satellite."," Here $\tau_n$ is the effective duration of peak $n$ at time $t_n$, and $t_{\rm dyn}$ is the dynamical time of the satellite."
+" We take tag,πιοyor. Where ry ds the halfanass radius of the stellar disk and 6,454 is the circular velocity of the appropriate NEW uodel at roy."," We take $t_{\rm dyn} = 2\pi
+r_{1/2}/v_{\rm rot}$, where $r_{1/2}$ is the half-mass radius of the stellar disk and $v_{\rm rot}$ is the circular velocity of the appropriate NFW model at $r_{1/2}$."
+ The cumulative tidal heating parameter is the smu. over all tidal peaks: This paraiecter determines the increase of the velocity dispersion of stars (eq. |6]]), The cumulative tidal heating parameter is the sum over all tidal peaks: This parameter determines the increase of the velocity dispersion of stars (eq. \ref{eq:sigma}] ])
+ iu our model of dwarf galaxy formation refsecisal) }., in our model of dwarf galaxy formation \\ref{sec:sam}) ).
+ We estimate the suppression of gas accretion due to the extragalactic UV. backeround using the filtering mass scale. derived by?.," We estimate the suppression of gas accretion due to the extragalactic UV background using the filtering mass scale, derived by."
+. Te defined. Mg as the mass of the halo which would lose half of the barvous. compared to the universal barvou fraction.," He defined $M_{\rm F}$ as the mass of the halo which would lose half of the baryons, compared to the universal baryon fraction."
+ This relates to the Jeans mass of the interealactic eas integrated over the cosinic history (eq. |, This relates to the Jeans mass of the intergalactic gas integrated over the cosmic history (eq. [
+"6] iu. ?)): where Afiy=2.5<loth10,|JT3/72M. µx0.59 is the mean molecular weight of the fully ionized eas. aud the integration exteuds over the expansion factor. o.","6] in ): where $M_{\rm J0} = 2.5\times 10^{11} h^{-1} \Omega_0^{-1/2} \mu^{-3/2} \
+M_{\sun}$, $\mu \approx 0.59$ is the mean molecular weight of the fully ionized gas, and the integration extends over the expansion factor, $a$."
+" The teiiperature of the cosmic gas T, is expressed in units of 10! K for convenience.", The temperature of the cosmic gas $T_4$ is expressed in units of $10^4$ K for convenience.
+" Tere we proposean analytical fit to the results ofον, assuniue a smile depeudence of the temperature on the expansion factor: Tifa)=(αρ). for eziαμ. ια)=1 fora,<axa. and Τι)=(afay)| foy a2a,."," Here we proposean analytical fit to the results of, assuming a simple dependence of the temperature on the expansion factor: $T_4(a) = (a/a_o)^{\alpha}$ for $a \le a_o$, $T_4(a) = 1$ for $a_o \le a \le a_r$, and $T_4(a) = (a/a_r)^{-1}$ for $a \ge a_r$."
+ These three distinct stages cau be clearly seen ou Fie., These three distinct stages can be clearly seen on Fig.
+ 1 of?., 1 of.
+". They correspond. to the 1) epoch before the frst TI regions fori 2>τω 2) the epoch of the overlap of iiultiple ΠΠ resious. 2,<2coL aud 3) the epoch of complete relonization. 2<2,."," They correspond, to the 1) epoch before the first HII regions form, $z>z_o$, 2) the epoch of the overlap of multiple HII regions, $z_r<z<z_o$, and 3) the epoch of complete reionization, $z<z_r$."
+" In the first stage. before redshift +,=Lia, Les. the temperature is rising as the newly-formed. stars jonize their neighboring regions."," In the first stage, before redshift $z_o \equiv 1/a_o-1
+\approx 8$ , the temperature is rising as the newly-formed stars ionize their neighboring regions."
+ The parameter à controls the rate of erowth of the extragalactic UV. flux: we fiud a=6 to be the best fit., The parameter $\alpha$ controls the rate of growth of the extragalactic UV flux; we find $\alpha = 6$ to be the best fit.
+" During the overlap stage. between redshifts L, and :,—αν1lzmT. the temperature is kept constant atroughly Lo! EK as the cosmic WIT reeious overlap."," During the overlap stage, between redshifts $z_o$ and $z_r
+\equiv 1/a_r-1 \approx 7$, the temperature is kept constant atroughly $10^4$ K as the cosmic HII regions overlap."
+ After the universe is fully ionized. at redshifts below τε. the temperature falls adiabatically with the cosmic expansion.," After the universe is fully ionized, at redshifts below $z_r$ , the temperature falls adiabatically with the cosmic expansion."
+ With these analytical expressions for F(a). we integrate," With these analytical expressions for $T_4(a)$ , we integrate"
+Almost all (>85%) MSPs have been confirmed to be in binaries.,Almost all $(>85\%)$ MSPs have been confirmed to be in binaries.
+ The central idea behind the MSP-LMXD connection is that LAINBs can provide the long-livecl phase of moderate mass transfer rates thought necessary (o spin-up the neutron star to millisecond periods., The central idea behind the MSP-LMXB connection is that LMXBs can provide the long-lived phase of moderate mass transfer rates thought necessary to spin-up the neutron star to millisecond periods.
+" The MSP is supposed to be produced when (he accretion shuts olf,", The MSP is supposed to be produced when the accretion shuts off.
+ But does the accretion ever shut-off?, But does the accretion ever shut-off?
+ In order to have a shut-off of the mass accretion rale. M. the accretion evolution time needs to be shorter than (he spin-down lime scale or The accretion rate. AM. decreases as the orbital separation grows.," In order to have a shut-off of the mass accretion rate, $\dot M$, the accretion evolution time needs to be shorter than the spin-down time scale or The accretion rate, $\dot M$, decreases as the orbital separation grows."
+" At late times. for eravitational wave driven orbital evolution in the limit when the low mass companion of the neutron star is verv much less than that of the neutron star. Thus as M decreases. the ratio 7;,,/7,, decreases."," At late times, for gravitational wave driven orbital evolution in the limit when the low mass companion of the neutron star is very much less than that of the neutron star, Thus as $\dot M$ decreases, the ratio $\tau_{spin}/\tau_{\dot M}$ decreases."
+ For a neutron star with a magnetic field B—5x105 G. the two timescales are roughly equal at AZ~10.M./yr.," For a neutron star with a magnetic field $B = 5 \times 10^{8}$ G, the two timescales are roughly equal at $\dot M \sim 10^{-9} M_{\odot}/yr$."
+ The svstem (hus has (ime to come into equilibrium aud accretion never shuts off (Delove2003)., The system thus has time to come into equilibrium and accretion never shuts off \citep{del08}.
+. It was noted by Delove(2008) that in two of the four classes of LMXDBs. mass transler to the neutron star never shuts off.," It was noted by \cite{del08} that in two of the four classes of LMXBs, mass transfer to the neutron star never shuts off."
+ Both of these svstems have gradually declining; mass, Both of these systems have gradually declining mass
+ (45) to derive both «2 and . with yo finally obtained from the SZ protile.,"\ref{eq:tpol}) ) to derive both $\beta$ and $\gamma$, with $y_0$ finally obtained from the SZ profile."
+ In Figure. 7 we show the temperature and Compton—y profiles: for our example cluster. along with the best-fitting predictions of the polytropic ~-model. for the three different definitions of temperature.," In Figure \ref{fi:tmp4} we show the temperature and $y$ profiles for our example cluster, along with the best–fitting predictions of the polytropic $\beta$ –model, for the three different definitions of temperature."
+ The polytropic equation of state provides a reasonable approximation to all temperature profiles and. unlike the isothermal case. allows us to correctly predict also the Compton-j profile.," The polytropic equation of state provides a reasonable approximation to all temperature profiles and, unlike the isothermal case, allows us to correctly predict also the $y$ profile."
+" The corresponding distributions of Pa2 and  are shown in Figure 8 (we do not report the distribution of r,. since it is. by definition. identical to that of the isothermal model)."," The corresponding distributions of $\beta$ and $\gamma$ are shown in Figure \ref{fi:betagamma} (we do not report the distribution of $r_c$, since it is, by definition, identical to that of the isothermal model)."
+ For both quantities. the effect of using different definitions of temperature is rather small.," For both quantities, the effect of using different definitions of temperature is rather small."
+ As expected. using a polytropie temperature profile implies only a modest decrease of the 7 values. because of the weak temperature dependence of the cooling function.," As expected, using a polytropic temperature profile implies only a modest decrease of the $\beta$ values, because of the weak temperature dependence of the cooling function."
+ All the three distributions of  have an average value &1.2. similar to observational estimates (e.g.2).," All the three distributions of $\gamma$ have an average value $\simeq 1.2$, similar to observational estimates \citep[e.g.][]{2002ApJ...567..163D}."
+ Moreover. the seatter in this distribution is so small to make the isothermal ICM an extremely unlikely event.," Moreover, the scatter in this distribution is so small to make the isothermal ICM an extremely unlikely event."
+" The results obtained for 2, are shown in Figure 9.. and also reported in Table l.. using emission-weighted. electron and spectroscopic-like temperatures."," The results obtained for $D_A$ are shown in Figure \ref{fi:da}, and also reported in Table \ref{tab:res2}, using emission–weighted, electron and spectroscopic–like temperatures."
+ Quite interestingly. the improved quality of the fit to the profile of the Compton—y parameter now makes the distribution peak at a value much closer to the correct D4. independently of whether we use the whole sample or the subsample of relaxed clusters.," Quite interestingly, the improved quality of the fit to the profile of the $y$ parameter now makes the distribution peak at a value much closer to the correct $D_A$, independently of whether we use the whole sample or the subsample of relaxed clusters."
+ The angular-diameter distance is correctly recovered when using either the electron or the spectroscopic-like temperature with deviations which are always Z5 per cent. on average.," The angular–diameter distance is correctly recovered when using either the electron or the spectroscopic–like temperature with deviations which are always $\mincir 5$ per cent, on average."
+ This is a rather encouraging result. since it indicates that any bias. induced by using the temperature as measured from X-ray observations.," This is a rather encouraging result, since it indicates that any bias, induced by using the temperature as measured from X–ray observations,"
+is the deacceleration radius.,is the deacceleration radius.
+" 7=AME, i8 the energy in the explosion and 6/j is the initial half-opening angle of the jet.", $E=M_0\Gamma_0$ is the energy in the explosion and $\theta_0$ is the initial half-opening angle of the jet.
+" The above equations show that for the wind and the uniform ISM models Dx£,,/711 BES7. respectively.: as long as D»u6,1 where fun,=[dt(lc) is the observer time. { being the lab frame time and c the jet— velocity in units of e."," The above equations show that for the wind and the uniform ISM models $\Gamma\propto t_{obs}^{-1/4}$ $t_{obs}^{-3/8}$, respectively, as long as $\Gamma\gg\theta_0^{-1}$, where $t_{obs} = \int dt (1-v)$ is the observer time, $t$ being the lab frame time and $v$ the jet velocity in units of $c$."
+ Equations (4)) and (5)) are solved. subject to the boundary conditions yj.=yw1 for. 1. to determine P and 0 as functions of +.," Equations \ref{dynamic1})) and \ref{dynamic2}) ) are solved, subject to the boundary conditions $y_1=y_2=1$ for $x\ll1$ , to determine $\Gamma$ and $\theta$ as functions of $r$."
+ For a relativistic jet with O=@. re. fluid velocity in the radial direction. equation (5)) reduces to The solution of the equations (4)) and (5)) is a two-parameter family of functions. however in the relativistic case the solution depends only on the product (Cy.," For a relativistic jet with $\Theta=\theta$ , i.e. fluid velocity in the radial direction, equation \ref{dynamic2}) ) reduces to The solution of the equations \ref{dynamic1}) ) and \ref{dynamic2}) ) is a two-parameter family of functions, however in the relativistic case the solution depends only on the product $\theta_0\Gamma_0$."
+ One can solve equation (7)) approximately. ignoring the very early time behavior. to determine the time when the sideways expansion alters significantly the jet dynamics.," One can solve equation \ref{dyna}) ) approximately, ignoring the very early time behavior, to determine the time when the sideways expansion alters significantly the jet dynamics."
+ The two relations in equation (7)) can be combined to yield a first order differential equation for yy42=y which is given by with 7=£(3.s)(0)Py). aconstant. and£2.0m7/(3.s).," The two relations in equation \ref{dyna}) ) can be combined to yield a first order differential equation for $y_1 y_2 \equiv y$ which is given by with $\eta=f(3-s)(\theta_0\Gamma_0)$, a constant, and $\xi=x^{3-s}/(3-s)$."
+ An approximate solution to this equation is Thus. κΤό decreases monotonically with radius or time.," An approximate solution to this equation is Thus, $y\propto\Gamma\theta$ decreases monotonically with radius or time."
+" The transition to jet sideways expansion starts when the two terms in the above equation become equal. ie. ©~(j/16)7. and lasts for an interval in ς for which y decreases by a factor of ~3. or .r increases by a factor of ~3°'?ο),"," The transition to jet sideways expansion starts when the two terms in the above equation become equal, i.e. $\xi\sim (\eta/16)^2$, and lasts for an interval in $\xi$ for which $y$ decreases by a factor of $\sim 3$, or $x$ increases by a factor of $\sim 3^{3/(3-s)}$."
+ The Lorentz factor continues to fall during the transition by a factor of a few., The Lorentz factor continues to fall during the transition by a factor of a few.
+" Therefore. the transition time divided by the time at the start of the transition (1n. observer frame). during which a,dla1)/edlutiys increases from (35)/(8.—25) to approximately 1/2. is approximately ο*)."," Therefore, the transition time divided by the time at the start of the transition (in observer frame), during which $\alpha_1 \equiv -d\ln(\Gamma-1)/d\ln 
+t_{obs}$ increases from $(3-s)/(8-2s)$ to approximately $1/2$, is approximately $^{3/(3-s)}$."
+ The solution to 4 and y» can be obtained by inserting the expression for y into equation (7))., The solution to $y_1$ and $y_2$ can be obtained by inserting the expression for $y$ into equation \ref{dyna}) ).
+ However. 4 and y2 determined this way have much larger error than y and should not be used for any serious calculation.," However, $y_1$ and $y_2$ determined this way have much larger error than $y$ and should not be used for any serious calculation."
+" We solve equations (4)) and (5)) numerically and show the results for .c(£,,,) and o4(f,4,) in Figure |.", We solve equations \ref{dynamic1}) ) and \ref{dynamic2}) ) numerically and show the results for $x(t_{obs})$ and $\alpha_1(t_{obs})$ in Figure 1.
+" Note that the change to a, from one asymptotic value. corresponding to spherical shell expansion. to another. when sideways expansion is Well underway. takes a long time: the ratio of the final to the initial time for a change in à4 of 0.1 fora uniform ISM is ~10? whereas for s=2 the ratio is 10°."," Note that the change to $\alpha_1$ from one asymptotic value, corresponding to spherical shell expansion, to another, when sideways expansion is well underway, takes a long time; the ratio of the final to the initial time for a change in $\alpha_1$ of 0.1 for a uniform ISM is $\sim 10^2$ whereas for $s=2$ the ratio is $10^3$."
+ For the parameters chosen here o4=0.5 when E is of order a few., For the parameters chosen here $\alpha_1=0.5$ when $\Gamma$ is of order a few.
+" In the non-relativistic phase of the jet expansion a,1.2. as for a Sedov-Taylor spherical shock wave."," In the non-relativistic phase of the jet expansion $\alpha_1=1.2$, as for a Sedov-Taylor spherical shock wave."
+ The synchrotron spectrum in the co-moving frame is taken to be a sequence of power-laws with breaks at the self-absorption. synchrotron peak. and cooling frequencies. as presented in Sari. Narayan Piran (1998); these frequencies can be found in eg.," The synchrotron spectrum in the co-moving frame is taken to be a sequence of power-laws with breaks at the self-absorption, synchrotron peak, and cooling frequencies, as presented in Sari, Narayan Piran (1998); these frequencies can be found in eg."
+ Panaitescu Kumar (2000)., Panaitescu Kumar (2000).
+ All of our numerical results. unless otherwise stated. are obtained by integrating emission over equal arrival time surface.," All of our numerical results, unless otherwise stated, are obtained by integrating emission over equal arrival time surface."
+" Ignoring the radial structure of the jet. the flux received by an observer located on the jet axis is given by where 0), is the co-moving power per frequency at ν΄=S(lecosce)v.rcoso =(f.cfjj; and r,,;,, and r,,,,,.. are solutions of We ignore the angular integration when discussing the analytical calculation of the observed flux and its power-law decline with time."," Ignoring the radial structure of the jet, the flux received by an observer located on the jet axis is given by where $P'_{\nu'}$ is the co-moving power per frequency at $\nu' = 
+\gamma (1 - v\cos\psi) \nu$ , $r\cos\psi = ct - ct_{obs}$ and $r_{min}$ and $r_{max}$ are solutions of We ignore the angular integration when discussing the analytical calculation of the observed flux and its power-law decline with time."
+" The observed flux at a frequency that is greater than both the cooling frequency. i. and the synchrotron peak. μα. 18 proportional to Atearly times when D0. and PX£|B3oosif{S2s’ the flux decays as 7""! "," The observed flux at a frequency that is greater than both the cooling frequency, $\nu_c$, and the synchrotron peak, $\nu_m$ , is proportional to At early times when $\Gamma\gg\theta^{-1}$ and $\Gamma\propto
+t_{obs}^{-(3-s)/(8-2s)}$, the flux decays as $t_{obs}^{-(3p-2)/4}$ ."
+"Atlate times when Γ01 the power-law index forthe AU flux >=dluf,/dluton.=(0.s)laa(p ns. whereas =2dlny/dlatinsThere are two effects that determine the evolution of ./."," At late times when $\Gamma\theta\lta 1$ the power-law index for the flux $\beta \equiv -d\ln f_\nu/d\ln t_{obs} = (4-s)[\alpha_1(p+2)-1]/2 + sp/4 
++ \alpha_2$ , where $\alpha_2 \equiv -2d\ln y/d\ln t_{obs}$.There are two effects that determine the evolution of $\beta$."
+" One of them. theeffect. is purely geometrical and results from the angular opening ~D! of the relativistic observing cone becoming larger than the jet opening angle 0. 1.8. the observer ""sees"" the edge of the jet."," One of them, the, is purely geometrical and results from the angular opening $\sim\Gamma^{-1}$ of the relativistic observing cone becoming larger than the jet opening angle $\theta$, i.e. the observer “sees"" the edge of the jet."
+ The increase to ./ resulting from it is QoS(3sM(LIosooo decreases with time and therefore the jump in J is smaller for larger Jy., The increase to $\beta$ resulting from it is $\alpha_2\lta (3-s)/(4-s)$; $\alpha_2$ decreases with time and therefore the jump in $\beta$ is smaller for larger $\theta_0$.
+ The dimensionless time for ./ to increase by o» depends on the angular position of the observer w.r.t., The dimensionless time for $\beta$ to increase by $\alpha_2$ depends on the angular position of the observer w.r.t.
+ the Jet axis and is approximately the ratio of the time when the observer sees the far edge of the jet to the time when the near side of the jet becomes visible., the jet axis and is approximately the ratio of the time when the observer sees the far edge of the jet to the time when the near side of the jet becomes visible.
+" This time is given by 1sis therobability probabilitythe that theobserver observerlies lies within àan where P.,, angle oy of the jet axis.", This time is given by where $P_{\phi_0}$ is the probability that the observer lies within an angle $\phi_0$ of the jet axis.
+" For [νι=0.25. Ry, is 18.7 (81) for 5=0 (2). and during this time > increases by approximately 0.7 (0.4)."," For $P_{\phi_0}=0.25$, $R_{t_e}$ is 18.7 (81) for $s=0$ (2), and during this time $\beta$ increases by approximately 0.7 (0.4)."
+" The dependence of Ry, on oy becomes much weaker when the emission is integrated over equal arrival time surface (Figure 2).", The dependence of $R_{t_e}$ on $\phi_0$ becomes much weaker when the emission is integrated over equal arrival time surface (Figure 2).
+" This is because the effect of angular integration is to smear the jet-edge by an angle 1/T— 09/2. which sets the minimum valueof A, to be about 10 (107) for uniform (wind) models."," This is because the effect of angular integration is to smear the jet-edge by an angle $1/\Gamma\sim\theta_0/2$ , which sets the minimum valueof $R_{t_e}$ to be about 10 $^2$ ) for uniform (wind) models."
+ The other effect which leads to a steepening of the afterelow decay is dynamical and ts caused by the lateral spreading of the jet., The other effect which leads to a steepening of the afterglow decay is dynamical and is caused by the lateral spreading of the jet.
+ During the relativistic phasethe increase to ./ fromthe sidewaysexpansion is 6.5=(p|21s)ón4/2 das: nj and day can be read from Figure |., During the relativistic phasethe increase to $\beta$ fromthe sidewaysexpansion is $\delta\beta = (p+2)(4-s)\delta\alpha_1/2+\delta\alpha_2$ ; $\delta\alpha_1$ and $\delta\alpha_2$ can be read from Figure 1.
+" Since a, does| not asymptote to 0.5 so?+y during the relativistic sideways expansionof the jet.", Since $\alpha_1$ does not asymptote to 0.5 so $\beta\not=p$ during the relativistic sideways expansionof the jet.
+ The value of ./ does. however. approach," The value of $\beta$ does, however, approach"
+Rso0.,.
+ Here 7s is the global (spectroscopic) temperature. calculated according to the observrational mass-temperature relation by Arnaud et al. (," Here $T_{\rm x}$ is the global (spectroscopic) temperature, calculated according to the observational mass-temperature relation by Arnaud et al. ("
+2005) (see also Shimizu et al.,2005) (see also Shimizu et al.
+ 2003)., 2003).
+ The adopted initial temperature profile should represen a reasonable initial condition to study the radiative evolution of the cluster cores., The adopted initial temperature profile should represent a reasonable initial condition to study the radiative evolution of the cluster cores.
+ The outer region (x0.15Hua ) is no greatly influenced by radiative cooling and thus we decided to set it to agree with current observations.," The outer region $r\ge 0.15\, R_{500}$ ) is not greatly influenced by radiative cooling and thus we decided to set it to agree with current observations."
+ In section 3.322 we briefly discuss how the results depend on the choice of the initial conditions., In section 3.3 we briefly discuss how the results depend on the choice of the initial conditions.
+ We present results for two custer models., We present results for two cluster models.
+" The “hot”. massive cluster has a virial dark matter mass ον]=1.3107 T.. while the ""cold"" object ws ονομα=dd10 T.."," The “hot”, massive cluster has a virial dark matter mass $M_{\rm DM,hot} = 1.3 \times 10^{15}$ $_\odot$, while the “cold” object has $M_{\rm DM,cold} = 
+4.1 \times 10^{14}$ $_\odot$."
+ The concentration of the cark matter halo is calculated by fitting the cAZ relation given in Bullock et al. , The concentration of the dark matter halo is calculated by fitting the $c-M$ relation given in Bullock et al. (
+2001): ο=ασ(Mp/107AL.) 1H.,"2001): $c=8.35 \, (M_{\rm DM}/10^{14} 
+\;{\rm M}_\odot)^{-0.911}$ ."
+ The central galaxy is the same for he two models and has a stelar density profile following the approximation for a deprojected de Vaucouleurs law. given by Tellier Mathez (1987).," The central galaxy is the same for the two models and has a stellar density profile following the approximation for a deprojected de Vaucouleurs law, given by Mellier Mathez (1987)."
+" The total mass of the galaxy is assumed o be M,=5.8.10! M. and the effective radius is 22.=8.5 kpe.", The total mass of the galaxy is assumed to be $M_* = 5.8 \times 10^{11}$ $_\odot$ and the effective radius is $R_e = 8.5$ kpc.
+ These numbers aretypical or giant elliptical galaxies., These numbers aretypical for giant elliptical galaxies.
+ The initial gas temperature are {0=7.67 keV and 75=4.0 keV ‘or the hot cluster and the cold cluster respectively., The initial gas temperature are $T_0=7.67$ keV and $T_0=4.0$ keV for the hot cluster and the cold cluster respectively.
+ The initial central gas density is set by the requirement that the baryon fraction at the virial radius is similar to the cosmic one. i.e. we require μεν)=0.16 (e.g. Spergel e al.," The initial central gas density is set by the requirement that the baryon fraction at the virial radius is similar to the cosmic one, i.e. we require $f_{\rm b}(R_{\rm vir})=0.16$ (e.g. Spergel et al."
+ 2007)., 2007).
+" This results in p;=2.26«107"" mem  for the cold cluster and po=1.52.10ος26 gom ?for the hot cluster.", This results in $\rho_0=2.26 \times 10^{-26}$ g $^{-3}$ for the cold cluster and $\rho_0=1.82 \times 10^{-26}$ g $^{-3}$ for the hot cluster.
+ The central cooling time. defined as loool=2bpfpmΠΟΠΗACL). is ~3.4 Gyr (cold cluster) and 8.5 Gyr (hot cluster).," The central cooling time, defined as $t_{\rm cool} = 2.5 k\rho T/
+\mu m_p/n_{\rm e}n_{\rm H}\Lambda(T) $, is $\sim 3.4$ Gyr (cold cluster) and $\sim 8.5$ Gyr (hot cluster)."
+" In order to better compare the modelled temperature. with observations we present results for both the emission weighted temperature ρω). and the ""spectroscopic-like"" temperature d.C) defined in Mazzotta et al. (", In order to better compare the modelled temperature with observations we present results for both the emission weighted temperature $T_{ew}(r)$ and the “spectroscopic-like” temperature $T_{sp}(r)$ defined in Mazzotta et al. (
+2004).,2004).
+" Zi... is calculated in the usual way as the line of sight integral Zi.=[1diΓι. where ¢ is the (bolometric) gas emissivity. while 7.,={Loreoy{n dl."," $T_{ew}$ is calculated in the usual way as the line of sight integral $T_{ew}=\int{T \epsilon dl}/\int{\epsilon dl}$, where $\epsilon$ is the (bolometric) gas emissivity, while $T_{sp}=\int{T (n^2 T^{-0.75}) dl}/\int{n^2 T^{-0.75} dl}$ ."
+ According to Mazzotta et al. (, According to Mazzotta et al. (
+2004) ἐν 1s à good approximation to the temperature measured with Chandra and ΑΙ when the thermal structure of the ICM is complex.,2004) $T_{sp}$ is a good approximation to the temperature measured with ${\it Chandra}$ and ${\it XMM}$ when the thermal structure of the ICM is complex.
+ By performing the integral along the line of sight at any radius. we calculate a projected temperature profile which can be directly contrasted with the observed ones.," By performing the integral along the line of sight at any radius, we calculate a projected temperature profile which can be directly contrasted with the observed ones."
+" The global temperature in a given region (Le. O<r« πο). 1s then calculated with a similar weighted average: {0Roo) beep, "," The global temperature in a given region (i.e. $0<r<R_{500}$ ) is then calculated with a similar weighted average: $T_{sp}(0-R_{500}) = 
+\int_0^{R_{500}} T_{sp}(r) (n^2 T^{-0.75}) 4\pi r^2 dr / 
+\int_0^{R_{500}} n^2 T^{-0.75} 4\pi r^2 dr$ ."
+In the following we will normalize the radial distances by [ioo or Foo. the radii at which the overdensity is 200 or 500 respectively.," In the following we will normalize the radial distances by $R_{200}$ or $R_{500}$ , the radii at which the overdensity is 200 or 500 respectively."
+" For our two models we have: Moo,=1.41 Mpc. [ioo=0.93 Mpe (cold cluster) and foo=2.05 Mpe. Mpe thotcluster)."," For our two models we have: $R_{200}=1.41$ Mpc, $R_{500}=0.93$ Mpc (cold cluster) and $R_{200}=2.05$ Mpc, $R_{500}=1.34$ Mpc (hotcluster)."
+presented in Paper 1: 3=0.0.0.5.1.0 and =0.0.6.0(see Eq. 2)).,"presented in Paper 1: $\beta=0.0, -0.5, -1.0$ and $\gamma=0.0, 6.0$(see Eq. \ref{eqn:rin}) ),"
+" optical depth 76,7=1.0.2.0.3.0.6.0. 10.0: torus opening angle O=20"". 40"". 60"": and outer-to-inner radius ratio RawfRin =30 and. 100."," optical depth $\tau_{9.7}=1.0, 2.0, 3.0, 6.0, 10.0$ ; torus opening angle $\Theta$ $^o$, $^o$, $^o$; and outer-to-inner radius ratio $R_{out}/R_{in}$ =30 and 100."
+ This last set of parameters only holds when ;j=0.0. Le. when the dust density is constant with the distance from the centre.," This last set of parameters only holds when $\beta=0.0$, i.e. when the dust density is constant with the distance from the centre."
+" For the moclels with 43«0.0. Roa, is re-caleulatec to be the distance at which the cust density drops at of its value at /2;,."," For the models with $\beta < 0.0$, $R_{out}$ is re-calculated to be the distance at which the dust density drops at of its value at $R_{in}$."
+ For the cold dust component. which is the major contributor to the bolometric emission at wavelengths longer than ~30Πα we choose six well-studied observational SB templates: AIS2 as a representative of a “typical” SD LIU emission. Arp 220 as representative of a very cxtinguished starburst and the templates of NGCL482. ΝαςΕΙ. NGC5253 and NOGCYTIA. as intermediate SB templates.," For the cold dust component, which is the major contributor to the bolometric emission at wavelengths longer than $\sim 30$, we choose six well-studied observational SB templates: M82 as a representative of a “typical” SB IR emission, Arp 220 as representative of a very extinguished starburst and the templates of NGC1482, NGC4102, NGC5253 and NGC7714, as intermediate SB templates."
+ The SED fitting procedure is that of a standard. 47 minimisation., The SED fitting procedure is that of a standard $\chi^2$ minimisation.
+ As a first step. the UV. to ποπ. (rest-frame) data points are fitted. by means of a stellar host component that dominates the emission in this range (since the aceretion emission. of the low luminosity AGN o£ the sample is most probably masked. under the stellar light).," As a first step, the UV to near-IR (rest-frame) data points are fitted by means of a stellar host component that dominates the emission in this range (since the accretion emission of the low luminosity AGN of the sample is most probably masked under the stellar light)."
+" The best fit to the observed data points in this wavelength range is Found by exploring a grid of τος and L(V"") values (see Sec. 3.1)).", The best fit to the observed data points in this wavelength range is found by exploring a grid of $\tau_{sf}$ and $E(B-V)$ values (see Sec. \ref{sec:stars}) ).
+ In the next step. the mic-lh points are then fit by a torus component.," In the next step, the mid-IR points are then fit by a torus component."
+" Since all the objects under study are type 2 AGN and therefore the central source is hidden behind the torus. we will focus on the results of the run realised with torus models of high optical depth. τετ=1.0 with the exception of the Iow-z sample. for which we will allow a full run. including 7o,«1.0 torus models. since for these low luminosity objects obscuration might also be the result. of the dust in the host galaxy."," Since all the objects under study are type 2 AGN and therefore the central source is hidden behind the torus, we will focus on the results of the run realised with torus models of high optical depth, $\tau_{9.7} \ge 1.0$ with the exception of the $z$ sample, for which we will allow a full run, including $\tau_{9.7} < 1.0$ torus models, since for these low luminosity objects obscuration might also be the result of the dust in the host galaxy."
+ In order to properly. constrain the part of the SED where the cold dust is dominant. one would. need at least one EL measurement. in this case a το and/or 160 detection (the mid-Infrared 24 point alone is probably not enough to distinguish between a torus-like ancl a starburst-like emission).," In order to properly constrain the part of the SED where the cold dust is dominant, one would need at least one FIR measurement, in this case a 70 and/or 160 detection (the mid-Infrared 24 point alone is probably not enough to distinguish between a torus-like and a starburst-like emission)."
+ Corresponcingly. an SB component will only be allowed in the presence of a 70 (and/or 160) datapoint(s). along the lines of Paper 1.," Correspondingly, an SB component will only be allowed in the presence of a 70 (and/or 160) datapoint(s), along the lines of Paper 1."
+ A single deviation [rom this will be discussed in Section 4.1.., A single deviation from this will be discussed in Section \ref{sec:samplek}.
+ As detailed in Paper 1. the SED fitting leads to the computation of a series of physical parameters. namely the accretion Luminositv. ος the HX luminosity. Lin. integrated between 1 and 1000puni. the inner ancl outer torus racii. [s (defined in Eq. 2))," As detailed in Paper 1, the SED fitting leads to the computation of a series of physical parameters, namely the accretion luminosity, $L_{acc}$, the IR luminosity, $\rm L_{IR}$, integrated between 1 and 1000, the inner and outer torus radii, $R_{in}$ (defined in Eq. \ref{eqn:rin}) )"
+ and Raw. the optical depth. (or extinction) and hwdrogen column density along the line of sight. the covering factor.CE. ie. the fraction of the racliating source covered by the obscuring material. and the mass of dust. Apu.," and $R_{out}$, the optical depth (or extinction) and hydrogen column density along the line of sight, the covering factor,CF, i.e. the fraction of the radiating source covered by the obscuring material, and the mass of dust, $M_{Dust}$ ."
+ A quantity that was not dealt with, A quantity that was not dealt with
+"are the same for some set of ion species (or if αι/ει is the same), then equation (20)) implies that for these ions, where A=mi/mp, Z=qi/e, e is the proton charge, and /=2a/(1—a).","are the same for some set of ion species (or if $\alpha_{\rm i}/\epsilon_{\rm i}$ is the same), then equation \ref{eq:Tsat}) ) implies that for these ions, where $A = m_{\rm i}/m_{\rm p}$, $Z = q_{\rm i}/e$, $e$ is the proton charge, and $l = 2a/(1-a)$."
+ The value of / can also be expressed as where σ is the spectral index of the power spectrum defined in section 2.., The value of $l$ can also be expressed as where $\sigma $ is the spectral index of the velocity power spectrum defined in section \ref{sec:heating}. .
+" In this paper, it is velocityassumed that a=1/4 and 6=3/2, which gives /=2/3."," In this paper, it is assumed that $a = 1/4$ and $\sigma = 3/2$, which gives $l= 2/3$."
+" If instead a=1/3 and o=5/3, then |—1."," If instead $
+a= 1/3$ and $\sigma = 5/3$, then $l = 1$."
+ The radius at which (19)) and (20)) can be largestapplied is determined by the equationscondition that the parallel thermal speed remain <vA so that ion Landau damping and transit-time vj;damping of KAWs can be neglected., The largest radius at which equations \ref{eq:th2}) ) and \ref{eq:Tsat}) ) can be applied is determined by the condition that the parallel thermal speed $v_{\parallel \rm i}$ remain $\ll v_{\rm A}$ so that ion Landau damping and transit-time damping of KAWs can be neglected.
+" Assuming that mg and Bo are given by equations (8)) and (9)), that (r) for protons is no greater than the r-dependent (isotropic) T;temperature in equation (47) of ?,, and that protons are the ion species with the largest value of vj;, one finds that «vA/3 for all ion species at to at least I5Ro. vj"," Assuming that $n_0$ and $B_0$ are given by equations \ref{eq:n}) ) and \ref{eq:B0}) ), that $T_{\parallel \rm
+ i}(r)$ for protons is no greater than the $r$ -dependent (isotropic) temperature in equation (47) of \cite{cranmer05}, and that protons are the ion species with the largest value of $v_{\parallel \rm i}$, one finds that $v_{\parallel \rm i} < v_{\rm A}/3$ for all ion species at to at least $15 R_{\sun}$."
+;Temperature profiles from equations (19)) and (20)) are plotted in figure 3.., Temperature profiles from equations \ref{eq:th2}) ) and \ref{eq:Tsat}) ) are plotted in figure \ref{fig:Tcomp}.
+" The left panel of this figure shows Τι! for protons and O? ions at 1.8R;«r3.3Ro, and the right panel shows the Τι! profiles for four ion species out to 15Ro."," The left panel of this figure shows $T_{\perp \rm i}$ for protons and ${\rm O}^{+5}$ ions at $1.8
+R_{\sun}< r < 3.3 R_{\sun}$, and the right panel shows the $T_{\perp
+ \rm i}$ profiles for four ion species out to $15 R_{\sun}$."
+ The data points in both panels represent the ion kinetic temperatures obtained by ? fromUVCS observations., The data points in both panels represent the ion kinetic temperatures obtained by \cite{esser99} fromUVCS observations.
+" The model temperature profiles in this figure were constructed using the n, Bo, and U profiles in equations (8)) through (10)) and the values of L, and ὃνρ in (4)) and (5)) with dv,=155 km/s as in figure given1.."," The model temperature profiles in this figure were constructed using the $n$ , $B_0$, and $U$ profiles in equations \ref{eq:n}) ) through \ref{eq:defU}) ) and the values of $L_\perp$ and $\delta v_0$ given in equations \ref{eq:defLperp}) ) and \ref{eq:dv0b}) ) with $\delta v_{\rm a} = 155$ km/s as in figure \ref{fig:swv_corona}."
+ The value equationsc?=0.15 was chosen to match the observed O? temperatures at r 1.8R;., The value $c_2 = 0.15$ was chosen to match the observed ${\rm O}^{+5}$ temperatures at $ r> 1.8 R_{\sun}$.
+ The value αρ=0.71 was then chosen to match the observed at r> protonPossible sources," The value $\alpha_{\rm p} = 0.71$ was then chosen to match the observed proton temperatures at $ r>1.8
+R_{\sun}$."
+" temperaturesof error in 1.8Ro.figure 3 include uncertainties in the values of c5, a, αι, and (η)."," Possible sources of error in figure \ref{fig:Tcomp} include uncertainties in the values of $c_2$ , $a$ , $\alpha_{\rm i}$, and $U_{\rm i}(r)$."
+" For example, if c» is increased from 0.15 to 0.34, but all other model parameters are held fixed, then 7, decreases by a factor of 4.6 at 2Ro."," For example, if $c_2$ is increased from 0.15 to 0.34, but all other model parameters are held fixed, then $T_{\perp
+ \rm p}$ decreases by a factor of 4.6 at $r=2R_{\sun}$ ."
+" If ais increased from 1/4 to 1/3 but all other parameters are fixed at their original values, then Τιρ decreases by a factor of 27 at r=2Ro."," If $a$ is increased from 1/4 to 1/3 but all other parameters are fixed at their original values, then $T_{\perp \rm p}$ decreases by a factor of 27 at $r=2R_{\sun}$."
+" If αι is increased from 0.71 to 1 for protons, then 7y increases by a factor of 2.5 at r=2Ro."," If $\alpha_{\rm i}$ is increased from 0.71 to 1 for protons, then $T_{\perp \rm p}$ increases by a factor of 2.5 at $r=2R_{\sun}$."
+ The value of Uj(r) has less of an impact on the ion temperatures that follow from equations (19)) and (20))., The value of $U_{\rm i}(r)$ has less of an impact on the ion temperatures that follow from equations \ref{eq:th2}) ) and \ref{eq:Tsat}) ).
+" If U; is decreased from 1.75U to 1.5U for O*° ions, then T;; increases by only at r=2Ro."," If $U_{\rm i}$ is decreased from $1.75U$ to $1.5U$ for ${\rm O}^{+5}$ ions, then $T_{\perp \rm i}$ increases by only at $r=2R_{\sun}$ ."
+" There is, however, a potentially larger source of error associated with (η)."," There is, however, a potentially larger source of error associated with $U_{\rm i}(r)$."
+" As the relative velocity AU; between heavy ions and protons approaches va, the electric field in the heavy-ion frame decreases, since most of the Alfvénn-wave fluctuations propagate away from the Sun in the proton frame."," As the relative velocity $\Delta U_{\rm i}$ between heavy ions and protons approaches $v_{\rm A}$, the electric field in the heavy-ion frame decreases, since most of the Alfv\'enn-wave fluctuations propagate away from the Sun in the proton frame."
+" As a result, the stochastic heating rate Qi for heavy ions decreases (?)."," As a result, the stochastic heating rate $Q_{\rm i}$ for heavy ions decreases \citep{chandran10}."
+" The acceleration of heavy ions to relative flow speeds ~va provides a second possible mechanism (in addition to the reduction of ει through the increase of Τι) for saturating stochastic heating, provided the amplitudes of sunward-propagating waves are much less than the amplitudes of anti-Sunward waves at A,~pi."," The acceleration of heavy ions to relative flow speeds $\sim v_{\rm A}$ provides a second possible mechanism (in addition to the reduction of $\epsilon_{\rm i}$ through the increase of $T_{\perp \rm i}$ ) for saturating stochastic heating, provided the amplitudes of sunward-propagating waves are much less than the amplitudes of anti-Sunward waves at $\lambda_\perp \sim
+\rho_{\rm i}$."
+" The effects of AU; on Q| are neglected in figure 3,, which is justified at rS3Ro, where UVCS observations show that AU;<va (?).."," The effects of $\Delta U_{\rm i}$ on $Q_{\perp \rm i}$ are neglected in figure \ref{fig:Tcomp}, which is justified at $r\lesssim 3 R_{\sun}$, where UVCS observations show that $\Delta
+U_{\rm i} \ll v_{\rm A}$ \citep{kohl98}."
+" However, for the velocity profiles assumed in the construction of figure 3,, AU;~να near r=15Ro, that the heavy ion temperatures near r=15Rc maysuggesting be over-estimated by equation (20))."," However, for the velocity profiles assumed in the construction of figure \ref{fig:Tcomp}, $\Delta U_{\rm i} \sim v_{\rm A}$ near $r= 15 R_{\sun}$, suggesting that the heavy ion temperatures near $r= 15 R_{\sun}$ may be over-estimated by equation \ref{eq:Tsat}) )."
+" A self-consistency check on the assumed value of αι for protons can be obtained by comparing the cascade power Τ with the heating rate = required to sustain protons at the temperatureOpikpTip/(Imptexy) T,, given by equations (19)) and (20))."," A self-consistency check on the assumed value of $\alpha_{\rm i}$ for protons can be obtained by comparing the cascade power $\Gamma$ with the heating rate $Q_{\rm p, crit} = k_{\rm B} T_{\perp \rm p} /(m_{\rm
+ p} t_{\rm exp})$ required to sustain protons at the temperature $T_{\perp \rm p}$ given by equations \ref{eq:th2}) ) and \ref{eq:Tsat}) )."
+" If the condition crit«I' were not satisfied, then protons would absorb so muchQp energy from the cascade that αρ would drop below the assumed value of 0.71, thereby reducing Qpcrit below I’ as required by energy conservation."," If the condition $Q_{\rm p, crit} < \Gamma$ were not satisfied, then protons would absorb so much energy from the cascade that $\alpha_{\rm p}$ would drop below the assumed value of 0.71, thereby reducing $Q_{\rm p, crit}$ below $\Gamma$ as required by energy conservation."
+ The value of Qpcrit is shown in figure 4 for 1.8Ro«r15R; for the same /Vset of assumptions used to obtain the proton temperature profile in figure 3..," The value of $Q_{\rm p, crit}/\Gamma$ is shown in figure \ref{fig:HF} for $ 1.8 R_{\sun}< r < 15 R_{\sun}$ for the same set of assumptions used to obtain the proton temperature profile in figure \ref{fig:Tcomp}."
+ It can be seen that Qpcrit«I throughout this range of r.," It can be seen that $Q_{\rm p, crit} < \Gamma$ throughout this range of $r$."
+" However, some caution is warranted here."," However, some caution is warranted here."
+" As noted by ?,, equation (6)) likely overestimates the turbulent dissipation rate for two reasons."," As noted by \cite{chandran09c}, equation \ref{eq:Gamma0}) ) likely overestimates the turbulent dissipation rate for two reasons."
+" First, large-scale Alfvénn waves launched from the Sun must propagate some distance into the corona before their energy cascades all the way to the dissipation scale."," First, large-scale Alfvénn waves launched from the Sun must propagate some distance into the corona before their energy cascades all the way to the dissipation scale."
+" Second, equation (6)) is derived in the limit of small L,."," Second, equation \ref{eq:Gamma0}) ) is derived in the limit of small $L_\perp$."
+" As shown in figure 3 of ? (and in figure 5 of ? for the caseU= 0), finite values of L, reduce I' relative to the small-L, limit."," As shown in figure 3 of \cite{chandran09c} (and in figure 5 of \cite{dmitruk02} for the case$U=0$ ), finite values of $L_\perp$ reduce $\Gamma$ relative to the $L_\perp$ limit."
+ These effects are particularly important close to the Sun., These effects are particularly important close to the Sun.
+ The small value of Op/V atr2Re shown in figure 4 may thus significantly critunderestimate thefractionof the cascade power absorbed by protons at this location.," The small value of $Q_{\rm p, crit} /\Gamma$ at$r\sim 2 R_{\sun}$ shown in figure \ref{fig:HF} may thus significantly underestimate thefractionof the cascade power absorbed by protons at this location."
+" Low-frequency Alfvénn-wave/KAW turbulence offers apromising explanation for the detailed features of ion kinetic temperatures seen in UVCS observations of coronal holes, including the abrupt radial variations in the O*? temperature profile, the widely different temperatures of protons and O? ions, and the large ot temperature anisotropy (T;>> Τι)."," Low-frequency Alfvénn-wave/KAW turbulence offers apromising explanation for the detailed features of ion kinetic temperatures seen in UVCS observations of coronal holes, including the abrupt radial variations in the ${\rm O}^{+5}$ temperature profile, the widely different temperatures of protons and ${\rm O}^{+5}$ ions, and the large ${\rm O}^{+5}$ temperature anisotropy $T_{\perp \rm i} \gg
+T_{\parallel \rm i}$ )."
+" When β< 1, ion heating fromthe linear damping of low- Alfvénn waves and KAWS can be neglected, and stochastic heating is arguably the primary way in which"," When $\beta \ll 1$ , ion heating fromthe linear damping of low-frequency Alfvénn waves and KAWs can be neglected, and stochastic heating is arguably the primary way in which"
+The discovery. of ∙∙∙uiullisecond oscillatious∙∙ in. hermonuclear X-rav bursts provided the first direct measurements of the spin frequencies of jeutrou stars in low-mass N-rav: binaries (LAINBs)° (Strolunaveretal.1996:Strolunaver&\arkwardt 2002).,"The discovery of millisecond oscillations in thermonuclear X-ray bursts provided the first direct measurements of the spin frequencies of neutron stars in low-mass X-ray binaries (LMXBs) \citep{strohmayer96,strohmayer02}."
+". The preseuce of neutrou stars rotating with willisecoudπο periodsnu 03in NS-LAINBsτςPo was τωsuggestedenoonpato nore than a decade previously (Alparetal.1982) MM radio pulsars, xedominately. but", The presence of neutron stars rotating with millisecond periods in NS-LMXBs was suggested more than a decade previously \citep{alpar82} and is key in understanding the relation of LMXBs and millisecond radio pulsars.
+ A-ray burst oscillation sources fori two distinct ‘lasses: “fast oscillators” with frequencies close o twice the frequency difference of the kz quasiperiodic oscillations (QPOs) secu in the yersistent enadssion. and “slow oscillators” with yequencies near the difference of the kITz QPO requeneies (White&ZhaugLo97).," X-ray burst oscillation sources form two distinct classes: “fast oscillators” with frequencies close to twice the frequency difference of the kHz quasiperiodic oscillations (QPOs) seen in the persistent emission, and “slow oscillators” with frequencies near the difference of the kHz QPO frequencies \citep{white97}."
+". Tf the kIIz OPO frequency difference ds close to the spin requency of the neutron star. then the slow oscillators have a burst oscillation frequency close o the spin. frequency,. while. the fast. oscillators: MM⋅ nA] -- ee ↓≼↧↕↴∖↴↨↘↽↸∖⋅↖↽↕∐∏∐≼∐∖↥⋅↴∖↴↑⋜⋯"," If the kHz QPO frequency difference is close to the spin frequency of the neutron star, then the slow oscillators have a burst oscillation frequency close to the spin frequency, while the fast oscillators produce signals at twice the spin frequency."
+≼∐∐∶↴⋁↑↕∐∖↥⋅↸∖↕⋜↧↑↕∪∐∪↕⋟∫⇀⋀∖∟∖⊽↕≧↴∖↴ib nillisecondM pA NN not exclusively. in photospheric radius expausion bursts. while the slow oscillators," The fast oscillators produce burst oscillations predominately, but not exclusively, in photospheric radius expansion bursts, while the slow oscillators"
+As the interferometric gravitational wave detectors LIGO and VIRGO (??) reach their design sensitivities. the first detection of gravitational waves has become more imminent.,"As the interferometric gravitational wave detectors LIGO and VIRGO \citep{2010CQGra..27h4006H,2006CQGra..23S..63A} reach their design sensitivities, the first detection of gravitational waves has become more imminent."
+ Both detectors will undergo serious improvements to increase their sensitivity (??)..," Both detectors will undergo serious improvements to increase their sensitivity \citep{2009CQGra..26k4013S,2005CQGra..22S.461S}."
+ It is therefore important to investigate the properties of the primary candidate sources for detection. namely compact object binaries.," It is therefore important to investigate the properties of the primary candidate sources for detection, namely compact object binaries."
+ There have been a number of papers dealing with several properties of the population of compact binaries (e.g.. 22222222?).. ," There have been a number of papers dealing with several properties of the population of compact binaries \citep[e.g.,][]{2001A&A...376..950N,2003MNRAS.342.1169V,2004NewA....9....1D,2002ApJ...572..962S,2005ApJ...628..343P,2002MNRAS.331.1027D,2005MNRAS.363L..71D,2007ARep...51..308B,2010MNRAS.406..656K}. ."
+In particular. ?. presented their estimated detection rates.," In particular, \citet{2010CQGra..27q3001A} presented their estimated detection rates."
+ In addition. the mass spectrum (?).. and even spin. properties (2?) have been studied.," In addition, the mass spectrum \citep{2007AdSpR..39..285G}, , and even spin properties \citep{2004PhRvD..70l4020S,2010CQGra..27k4007M} have been studied."
+ In this paper. we present yet another aspect of the merging compact object binary population: the distribution of the eccentricity.," In this paper, we present yet another aspect of the merging compact object binary population: the distribution of the eccentricity."
+ From radio observations. we currently know of only six compact object binaries with merger timescales shorter than a Hubble time. all of them NS-NS (neutron star - neutron star) systems. and we know of no BH-NS (black hole - neutron star) nor BH-BH (black hole - black hole) system.," From radio observations, we currently know of only six compact object binaries with merger timescales shorter than a Hubble time, all of them NS-NS (neutron star - neutron star) systems, and we know of no BH-NS (black hole - neutron star) nor BH-BH (black hole - black hole) system."
+ The known NS-NS systems are listed along with their orbital parameters in Table |.., The known NS-NS systems are listed along with their orbital parameters in Table \ref{table:1}.
+ The observed NS-NS binaries have merger times Tierez 100 Myr., The observed NS-NS binaries have merger times $T_{merg}\succsim$ 100 Myr.
+ However. a significant fraction. of the population of the merging NS-NS may originate in. the so-called ultra-compact NS-NS binaries. which have much shorter merger timescales Τουx 100 Myr (2)..," However, a significant fraction of the population of the merging NS-NS may originate in the so-called ultra-compact NS-NS binaries, which have much shorter merger timescales $T_{merg}\precsim$ 100 Myr \citep{2002ApJ...571L.147B}."
+ In addition. the small number of known pulsars are indicative of there being a significant fraction of very eccentric binaries.," In addition, the small number of known pulsars are indicative of there being a significant fraction of very eccentric binaries."
+ Since no BH-NS nor BH-BH binaries are known. we can only rely o evolutionary considerations when estimating their number and properties.," Since no BH-NS nor BH-BH binaries are known, we can only rely on evolutionary considerations when estimating their number and properties."
+ It has been found that their number depends very strongly on the outcome of the common envelope phase when the secondary is on the Hertzsprung gap., It has been found that their number depends very strongly on the outcome of the common envelope phase when the secondary is on the Hertzsprung gap.
+ This phase will very likely end up as a merger and the formation of a Thorne-Zytkow object., This phase will very likely end up as a merger and the formation of a Thorne-Zytkow object.
+ However. it hàs been demonstrated that in a low-metallicity environment the common envelope mergers may be (to some extent) avoided and the BH-BH formation is very effective (?)..," However, it has been demonstrated that in a low-metallicity environment the common envelope mergers may be (to some extent) avoided and the BH-BH formation is very effective \citep{2010ApJ...715L.138B}."
+ The eccentricity of a compact object binary may potentially be derived by analyzing the inspiral signal. provided that the eccentricity is significant.," The eccentricity of a compact object binary may potentially be derived by analyzing the inspiral signal, provided that the eccentricity is significant."
+ In this paper. we investigate the eccentricity distributions in the frequency band of the currently working and future detectors of gravitational waves.," In this paper, we investigate the eccentricity distributions in the frequency band of the currently working and future detectors of gravitational waves."
+ For the currently working detectors (LIGO and VIRGO). we assume that the sensitivity of the detectors will allow us to measure the signal for the frequencies starting at 30 Hz.," For the currently working detectors (LIGO and VIRGO), we assume that the sensitivity of the detectors will allow us to measure the signal for the frequencies starting at 30 Hz."
+ This may not be accurate for the current state of these instruments. but it does accurately represent the predicted sensitivity of the Advanced LIGO/VIRGO detectors.," This may not be accurate for the current state of these instruments, but it does accurately represent the predicted sensitivity of the Advanced LIGO/VIRGO detectors."
+ We consider two future detectors: the Einstein Telescope (?) and DECIGO (??)..," We consider two future detectors: the Einstein Telescope \citep{2010arXiv1003.1386V} and DECIGO \citep{2006AstHe..99..490K,2001PhRvL..87v1103S}."
+ For the Einstein Telescope. we assume that binaries shall be detectable from 3 Hz. and for DECIGO we assume that the lowest frequency detectable is 0.3 Hz.," For the Einstein Telescope, we assume that binaries shall be detectable from 3 Hz, and for DECIGO we assume that the lowest frequency detectable is 0.3 Hz."
+ In all cases. these have to be treated as indicative numbers that roughly describe these instruments.," In all cases, these have to be treated as indicative numbers that roughly describe these instruments."
+ In section 2. we describe the model used to investigate the population of compact object binaries.," In section 2, we describe the model used to investigate the population of compact object binaries."
+" Section 3 presents the results for the current and future gravitational wave detectors,", Section 3 presents the results for the current and future gravitational wave detectors.
+ In section +. we summarize and discuss the results.," In section 4, we summarize and discuss the results."
+ To model the population of compact object binaries. we used theStarTrack population synthesis code (?)..," To model the population of compact object binaries, we used the population synthesis code \citep{2002ApJ...572..407B}."
+ It perform a suite of Monte Carlo simulations of the stellar evolution of stars In. environments of two typical metallicities: Z=Ze=0.02 and Z=10%Z.0.002 (e.g..?)..," It perform a suite of Monte Carlo simulations of the stellar evolution of stars in environments of two typical metallicities: $Z=Z_{\odot}=0.02$ and $Z=10\% \, Z_{\odot}=0.002$ \citep[e.g.,][]{2010ApJ...715L.138B}."
+ In these calculations. we enployed the recent estimates of mass loss rates (2)..," In these calculations, we employed the recent estimates of mass loss rates \citep{2010ApJ...714.1217B}."
+ We calculate a population of2 million massive binary stars. tracking the ensuing formation of relativistic binary compact objects: ¢ouble neutron stars (NS-NS). double black hole binaries (BH-BH). and mixed systems (BH-NS).," We calculate a population of 2 million massive binary stars, tracking the ensuing formation of relativistic binary compact objects: double neutron stars (NS-NS), double black hole binaries (BH-BH), and mixed systems (BH-NS)."
+ Our modeling utilizes vpdated stellar and binary physics. including results from supernova simulations (?)9 and compact object formation (?).. incorporating elaborate mechanisms for treating stellar interactions such as mass transfer episodes (?)9 οὐ tidal synchronization and eircularization (?)..," Our modeling utilizes updated stellar and binary physics, including results from supernova simulations \citep{2001ApJ...554..548F}
+ and compact object formation \citep{1996ApJ...457..834T}, incorporating elaborate mechanisms for treating stellar interactions such as mass transfer episodes \citep{2008ApJS..174..223B} or tidal synchronization and circularization \citep{1981A&A....99..126H}."
+ We place special emphasis on the common envelope evolution phase (?).. which is crucial for close double compact object formation becausethe attendantmass transfer permits an efficient hardening of the binary.," We place special emphasis on the common envelope evolution phase \citep{1984ApJ...277..355W}, which is crucial for close double compact object formation becausethe attendantmass transfer permits an efficient hardening of the binary."
+ This orbital contraction can be sufficiently efficient, This orbital contraction can be sufficiently efficient
+5hort-timescale variability in (he optical afterglow of gammia-ray bursts (GRBs) can bea tool lor understanding the details of GRB origins.,Short-timescale variability in the optical afterglow of gamma-ray bursts (GRBs) can be a tool for understanding the details of GRB origins.
+ The optical afterglow of a GRB stems [rom a relativistic blastwave as it slows down in (he interstellar medium/stellar wind surrounding a source hvpernova (e.g. Stanek οἱ al., The optical afterglow of a GRB stems from a relativistic blastwave as it slows down in the interstellar medium/stellar wind surrounding a source hypernova (e.g. Stanek et al.
+ 2003: Matheson et al., 2003; Matheson et al.
+ 2003)., 2003).
+ Understanding variations in intensiv of the optical alterglow. or lack thereof. can vield insights into the details of the interaction between GRBs and their surroundings (e.g. Wang Loeb 1999).," Understanding variations in intensity of the optical afterglow, or lack thereof, can yield insights into the details of the interaction between GRBs and their surroundings (e.g. Wang Loeb 1999)."
+ Such >104 , Such $> 10$ 
+The formation of massive stars occurs deep inside dense molecular cloud cores.,The formation of massive stars occurs deep inside dense molecular cloud cores.
+ The high luminosity of massive (proto)stars stirs and heats up the surrounding gas. leading to the evaporation of ice mantles around dust grains and to high degrees of excitation of the molecules.," The high luminosity of massive (proto)stars stirs and heats up the surrounding gas, leading to the evaporation of ice mantles around dust grains and to high degrees of excitation of the molecules."
+ This stage. just prior to ionization and destruction of the molecular gas. is characterized by strong line and dust emission from à compact region. and is called a hot molecular core (e.g.?)..," This stage, just prior to ionization and destruction of the molecular gas, is characterized by strong line and dust emission from a compact region, and is called a hot molecular core \citep[e.g.][]{Cesaroni05}."
+ Since these sources are short-lived. they are rare and hence far away. which hampers the determination of their physical structure (density. temperature. velocity. molecular abundances).," Since these sources are short-lived, they are rare and hence far away, which hampers the determination of their physical structure (density, temperature, velocity, molecular abundances)."
+ Nonetheless. such knowledge is a key information for understanding the process of massive star formation (?) and for the interpretation of molecular line surveys (e.g.???).. and has been constrained by radiative transfer modeling of both single-dish (e.g.?2) and interferometry observations (e.g.?)..," Nonetheless, such knowledge is a key information for understanding the process of massive star formation \citep{ZinneckerYorke} and for the interpretation of molecular line surveys \citep[e.g.][]{Schilke06,Belloche07,Bergin10}, and has been constrained by radiative transfer modeling of both single-dish \citep[e.g.][]{Hatchell03,Rolffs11apex} and interferometry observations \citep[e.g.][]{Osorio09}."
+ Especially interesting is the hottest molecular gas. which is closest to the heating sources.," Especially interesting is the hottest molecular gas, which is closest to the heating sources."
+ A good tracer is vibrationally excited HCN (521). whose rotationally excited levels lie. ~1000 K higher in energy than the ground vibrational state. and which has transitions in the cm-wave regime. whose radiation is not obscured by dust.," A good tracer is vibrationally excited HCN $v_2$ =1), whose rotationally excited levels lie $\sim 1000$ K higher in energy than the ground vibrational state, and which has transitions in the cm-wave regime, whose radiation is not obscured by dust."
+ These so-called direct (-type transitions are intrinsically weak. but have been detected in Galactic high-mass star-forming regions (?) as well as toward the proto-planetary nebula CRL618 (?) and the starburst galaxy Arp220 (?)..," These so-called direct $\ell$ -type transitions are intrinsically weak, but have been detected in Galactic high-mass star-forming regions \citep{Thorwirth01} as well as toward the proto-planetary nebula CRL618 \citep{Thorwirth03} and the starburst galaxy Arp220 \citep{Salter08}."
+ Obviously. this hot gas is very compact. and high spatial resolution is needed to investigate it.," Obviously, this hot gas is very compact, and high spatial resolution is needed to investigate it."
+ In this paper. we present VLA observations (with. among the highest resolution ever obtained for thermal molecular line emission) of three massive hot cores which have strong HCN direct. (-type lines (as measured with the Effelsberg 100-m telescope. Sect. ??)).," In this paper, we present VLA observations (with among the highest resolution ever obtained for thermal molecular line emission) of three massive hot cores which have strong HCN direct $\ell$ -type lines (as measured with the Effelsberg 100-m telescope, Sect. \ref{sec:sd}) )."
+ The sources are G10.47+0.03 (e.g. 2)... which lies at a distance of 10.6 kpe (2). and has a luminosity of about 7xΙΟ L. (2).. and SgrB2-N and -M near the Galactic center at 7.8 kpe distance (?).. which have total luminosities of about 8.4»10° and 6.3x10° Ls. respectively (?)..," The sources are G10.47+0.03 \citep[e.g. ][]{Wyrowski99}, , which lies at a distance of 10.6 kpc \citep{Pandian08} and has a luminosity of about $7\times 10^5$ $_\odot$ \citep{Cesaroni10}, and SgrB2-N and -M near the Galactic center at 7.8 kpc distance \citep{Reid09}, which have total luminosities of about $8.4\times 10^5$ and $6.3\times 10^6$ $_\odot$, respectively \citep{Goldsmith92}."
+ All sources are sites of very active high-mass star formation and contain several regions., All sources are sites of very active high-mass star formation and contain several regions.
+ After describing observations and data reduction (Sect. ?2)).," After describing observations and data reduction (Sect. \ref{sec:obs}) ),"
+ we present the results (Sect. ??)), we present the results (Sect. \ref{sec:results}) )
+ and the radiative transfer modeling (Sect. ??)).," and the radiative transfer modeling (Sect. \ref{sec:modeling}) ),"
+ discuss. the implications (Sect. 22)), discuss the implications (Sect. \ref{sec:discussion}) )
+ and draw conclusions (Sect. 7))., and draw conclusions (Sect. \ref{sec:concl}) ).
+ With the National Radio Astronomy Observatory's Very Large Array (VLA) in its BnA configuration we observed three massive star forming regions in the direct (-type line of vibrationally excited HCN at 40.7669 GHz (project AR687)., With the National Radio Astronomy Observatory's Very Large Array (VLA) in its BnA configuration we observed three massive star forming regions in the direct $\ell$ -type line of vibrationally excited HCN at 40.7669 GHz (project AR687).
+ This transition connects the two levels of the w=l. /=13 doublet. whose energies lie 1411 and 1413 K above the ground state. respectively.," This transition connects the two levels of the $v_2$ =1, $J$ =13 doublet, whose energies lie 1411 and 1413 K above the ground state, respectively."
+ Each of the sources was observed during LST 15-21: G10.47+0.03 on 2009 January 30. SgrB2-M on January 3l. and SerB2-N on February |.," Each of the sources was observed during LST 15-21; G10.47+0.03 on 2009 January 30, SgrB2-M on January 31, and SgrB2-N on February 1."
+ Bandpass calibration was done on 3C279 and 3C454.3. and primary flux calibration on 3C286.," Bandpass calibration was done on 3C279 and 3C454.3, and primary flux calibration on 3C286."
+ The observations were performed in fast-switching mode with 150s on the source and 30s on a phase calibrator (1755-225 for G10.47+0.03 and SgrA* for SgrB2). yielding an on-source integration time of about 4.5 hours per source.," The observations were performed in fast-switching mode with 150s on the source and 30s on a phase calibrator (1755-225 for G10.47+0.03 and SgrA* for SgrB2), yielding an on-source integration time of about 4.5 hours per source."
+ For amplitude calibration. 1733-130 (G10.47+0.03) or SgrA* (SgrB2) was observed every 20 minutes. and reference pointing was done every hour on these sources.," For amplitude calibration, 1733-130 (G10.47+0.03) or SgrA* (SgrB2) was observed every 20 minutes, and reference pointing was done every hour on these sources."
+ For G10.47+0.03. we centered one IF with 6.25 MHz bandwidth at 40.762 GHz (shifted from the rest frequency due to the apparent source motion) and one with 25 MHz bandwidth at 40.78 GHz to increase the continuum sensitivity.," For G10.47+0.03, we centered one IF with 6.25 MHz bandwidth at 40.762 GHz (shifted from the rest frequency due to the apparent source motion) and one with 25 MHz bandwidth at 40.78 GHz to increase the continuum sensitivity."
+ For the SgrB2 sources. the line is broader. so we used two overlapping IFs with 6.25 MHz bandwidth centered at 40.761 GHz and 40.765 GHz.," For the SgrB2 sources, the line is broader, so we used two overlapping IFs with 6.25 MHz bandwidth centered at 40.761 GHz and 40.765 GHz."
+ The channel width is0.72 km s7!., The channel width is0.72 km $^{-1}$ .
+uniformly rotating disk withp=10 and v=(2y/Ro.τσ ΠΟ). whore R=Yar| ye.,"uniformly rotating disk with$\rho = 10$ and ${\bf v} = (-2y/R_0, +2x/R_0, 0)$ , where $R =
+\sqrt{x^2 + y^2}$ ."
+ The ring Ry«E«€=0.115 is occupied by a transition region with p—1|OFC?) and v—2(2yf(IO/I.|ef(GII0). where Outside Bo»Ly. the velocity is set to zero and p=1.," The ring $R_0 < R < R_1 = 0.115$ is occupied by a transition region with $\rho = 1 + 9f(R)$ and ${\bf v} = (-2y f(R)/R, +2x f(R)/R, 0)$, where Outside $R > R_1$, the velocity is set to zero and $\rho = 1$."
+ Εις problem uses an ideal gas equation of state with ~=1.4., This problem uses an ideal gas equation of state with $\gamma = 1.4$.
+ The particles are distributed in the same way as in strong blast wave problem refsectiblast))., The particles are distributed in the same way as in strong blast wave problem \\ref{sect:blast}) ).
+ The problem is solved in both rest ancl boosted frames. with (ns=(5/0.15.5/0.15.5/0.15).," The problem is solved in both rest and boosted frames, with $v_{\rm boost} = (5/0.15, 5/0.15, 5/0.15)$."
+ The plots of density. gas and magnetic pressure. and mach number are shown in 14.," The plots of density, gas and magnetic pressure, and mach number are shown in ."
+. The 1Dslicesof magnetic field at =0.5 and jy—0.5 are show in, The 1Dslicesof magnetic field at $x = 0.5$ and $y = 0.5$ are show in
+"star at high densities, capture of degenerate electrons dominates over the neutrino effects, and they turn neutron-rich again.","star at high densities, capture of degenerate electrons dominates over the neutrino effects, and they turn neutron-rich again."
+" The thermodynamic conditions (density, temperature, and entropy), which are responsible for the nucleosynthesis results, are shown in Figure for the ejected mass zones."," The thermodynamic conditions (density, temperature, and entropy), which are responsible for the nucleosynthesis results, are shown in Figure for the ejected mass zones."
+" These figures indicate a temperature, density, and entropy maximum when the quark-hadron phase transition occurs, which causes a second core bounce (about 0.4 s after the first bounce at nuclear densities; see Figure and an outgoing shock front forms."," These figures indicate a temperature, density, and entropy maximum when the quark–hadron phase transition occurs, which causes a second core bounce (about $0.4$ s after the first bounce at nuclear densities; see Figure \ref{fig-properties}) ) and an outgoing shock front forms."
+" The expansion follows2)) a close to constant entropy, i.e. is adiabatic, once matter is ejected (after 0.4 s for the prompt ejection and after 1.5 s for the delayed "," The expansion follows a close to constant entropy, i.e. is adiabatic, once matter is ejected (after $0.4$ s for the prompt ejection and after $1.5$ s for the delayed ejection)."
+"The matter which initially fell back onto the ejection).proto-neutron star and is finally ejected by the NDW, experiences heating and an entropy rise due to this energy deposition by neutrinos."," The matter which initially fell back onto the proto-neutron star and is finally ejected by the NDW, experiences heating and an entropy rise due to this energy deposition by neutrinos."
+" 'The properties of the neutrino and anti-neutrino flux (Iuminosities, average energies and neutrino sphere radii) can be found in Figure"," The properties of the neutrino and anti-neutrino flux (luminosities, average energies and neutrino sphere radii) can be found in Figure"
+2001).,.
+.. The conclusion is that in imost cases the RGD tip detection is quite insensitive to the effect of this contamination., The conclusion is that in most cases the RGB tip detection is quite insensitive to the effect of this contamination.
+ This result is further coufirmed by looking at the results preseuted here., This result is further confirmed by looking at the results presented here.
+ Indeed. the RGB tip positions measured in Fields 2 aud 3 do not significantly differ from thle positions measured im any other field.," Indeed, the RGB tip positions measured in Fields 2 and 3 do not significantly differ from the positions measured in any other field."
+ To estimate the errors connected with the detection ofthe RGB tip. we adopted a bootstrap ο strateev simular to the one prescuted iu Meéndezetal.(2002)..," To estimate the errors connected with the detection of the RGB tip, we adopted a bootstrap resampling strategy similar to the one presented in \citet{2002AJ....124..213M}."
+ The sample of stars chosen to represeut the ROB was resampled 500 times. and the RGB tip measured for cach realization.," The sample of stars chosen to represent the RGB was resampled 500 times, and the RGB tip measured for each realization."
+ The vans., The r.m.s.
+ of the results is then quoted iu columns 3 aud 5 of Table 1.. for ED and ML aethods. respectively.," of the results is then quoted in columns 3 and 5 of Table \ref{tab1}, for ED and ML methods, respectively."
+ The first calibration of the absolute magnitude of the RGD tip dates back to the early 1990s., The first calibration of the absolute magnitude of the RGB tip dates back to the early 1990's.
+ Leeetal.(1993). defined the distance iiodulus based ou the RGD tip as where DC; is the bolometric correction to the I magnitude. and ζωτης is the bolometric magnitude of the TRGB.," \citet{1993ApJ...417..553L} defined the distance modulus based on the RGB tip as where $BC_I$ is the bolometric correction to the I magnitude, and $M_{bol,TRGB}$ is the bolometric magnitude of the TRGB."
+" BC; aud ALiTnOBp are given in DaCosta&Avmandroff(1990) as DC,=(sslοσονTy and Af,=0.19[Fe/TI| 3.81."," $BC_I$ and $M_{bol,TRGB}$ are given in \citet{1990AJ....100..162D} as $\rm{BC_I}=0.881-0.243(V-I)_0$ and $M_{bol}=-0.19\rm{[Fe/H]}-3.81$ ."
+ These calibrations are based on the distance scale of Leeetal.(1990). where the magnitude of RR Lvrae stars is AA(RR)=0.82| O.17|Fe/TI]., These calibrations are based on the distance scale of \cite{1990ApJ...350..155L} where the magnitude of RR Lyrae stars is $M_V(RR)=0.82 + 0.17 \rm{[Fe/H]}$ .
+ All these relations are based on a sual sample of RGB stars observed in a few template elobular clusters. aud they oulv cover the range 2.17<|[Fe/Il]— (0.71.," All these relations are based on a small sample of RGB stars observed in a few template globular clusters, and they only cover the range $-2.17 < \rm{[Fe/H]} < -0.71$ ."
+ An extensive, An extensive
+1999).,.
+". We check that when L.7<u<OL and 1<Q,,50.7. the formmlaecs provided by this work is accurate to 2%."," We check that when $-1.7<w<-0.4$ and $0.1<\Omega_{m0}<0.7$, the formulaes provided by this work is accurate to $2\%$."
+" To relate the mass of a galaxy cluster with its characteristic X-ray temperature. the ratio of cluster/backeround matter densities at the virialization poiut is another very important parameter. It can be shown that Dy(aj,sy»o."," To relate the mass of a galaxy cluster with its characteristic X-ray temperature, the ratio of cluster/background matter densities at the virialization point is another very important parameter, It can be shown that $D_1(a)_{a\rightarrow0}\rightarrow a$."
+" Using eq(1 1). we have a). Substitutiuen eq(21))211) iutoi eq(E7)) ""ELIand ususAue the facfact that. iL""mStittoο=|]?τα we EMcan geton Due) About A's calculation. using definition eq(19)) wo have To calculate the second factor of the above equations’ right-most part. we can use energy conscrving condition aud virial theorcun."," Using \ref{yofxAnsaltz}) ), we have a), Substituting \ref{roaagoesto0}) ) into \ref{deltacDefinition}) ) and using the fact that $[\frac{\rho_{mc}}{\rho_{mb}}]_{a\rightarrow0}=[\frac{a}{a_p}]^3$, we can get ]D_1(a_c) About $\Delta_c^\prime$ s calculation, using definition \ref{DeltacDefinition1}) ) we have To calculate the second factor of the above equations' right-most part, we can use energy conserving condition and virial theorem."
+ If Quintesseuce clusters svuchronously with ordinary iatters. by assuming that at the collapse point. the system virializes fully. we can get the following relations:ιοί just as we do iu (Zeug&Cao2005)..," If Quintessence clusters synchronously with ordinary matters, by assuming that at the collapse point, the system virializes fully, we can get the following relations: just as we do in \citep[]{SphereI}."
+" Now. since Quintessence is asstumed not to cluster ou the scale of. galaxy clusters at all. we can only write. down tlie following relations: in the above four equations. των Cine aud Uto denote the iuatteranatter. matter-Quintesseuce and Quiutesseuce-Quiutessence eravitation potentials respectively,"," Now, since Quintessence is assumed not to cluster on the scale of galaxy clusters at all, we can only write down the following relations: in the above four equations, $U_{mm}$, $U_{mQ}$ and $U_{QQ}$ denote the matter-matter, matter-Quintessence and Quintessence-Quintessence gravitation potentials respectively."
+" The subscripts , aud 42 indicate that quautities carrying them should take values at the collapse aud turi around time respectively."," The subscripts $_{,c}$ and $_{,ta}$ indicate that quantities carrying them should take values at the collapse and turn around time respectively."
+ Explicitly.," Explicitly,"
+were chosen {ο be free of active regions. plage or dominance of a single polarity.,"were chosen to be free of active regions, plage or dominance of a single polarity."
+ No other cata was used to determine if (he region was. for example. part of a coronal hole.," No other data was used to determine if the region was, for example, part of a coronal hole."
+ Given (heir equatorial locations. however. it seems unlikely that much coronal hole data was included in our survev.," Given their equatorial locations, however, it seems unlikely that much coronal hole data was included in our survey."
+ We cannot (therefore sav whether (he null point density. applies to coronal holes as well as to equatorial quiet Sun., We cannot therefore say whether the null point density applies to coronal holes as well as to equatorial quiet Sun.
+ These findings place m a clearer context observations and models with putative relationships to coronal null points., These findings place in a clearer context observations and models with putative relationships to coronal null points.
+ If the quiet Sun corona is being heated primarily τοσα dissipation ab null points. then the majority of this heat is supplied very low in the corona.," If the quiet Sun corona is being heated primarily through dissipation at null points, then the majority of this heat is supplied very low in the corona."
+ The region between 2=1.5 Mm and 2=4.5 Main contains three times as many nulls as the whole space above it: il would therefore probably receive a proportional share of the heating., The region between $z=1.5$ Mm and $z=4.5$ Mm contains three times as many nulls as the whole space above it; it would therefore probably receive a proportional share of the heating.
+ H remains (o be investigated whether (he actual thermal structure of the quiet Sun. corona is consistent with this distribution of heating. and is therefore consistent with heating by enerev dissipation at null points.," It remains to be investigated whether the actual thermal structure of the quiet Sun corona is consistent with this distribution of heating, and is therefore consistent with heating by energy dissipation at null points."
+ Along similar lines. if coronal jets were hypothesized to form anvwhere a coronal null point occurs. (hen our results can be used in an prediction of their frequency.," Along similar lines, if coronal jets were hypothesized to form anywhere a coronal null point occurs, then our results can be used in an prediction of their frequency."
+ Moreover. the distribution of (heir cusp heights should be somehow influenced by the distribution of null heights: short jets should out-number tall ones.," Moreover, the distribution of their cusp heights should be somehow influenced by the distribution of null heights: short jets should out-number tall ones."
+ Actual jets undoubtedly involve many factors. in addition to magnetic geometry. which influence the distribution of their properties.," Actual jets undoubtedly involve many factors, in addition to magnetic geometry, which influence the distribution of their properties."
+ SGLL it would be difficult to invoke null points as a requirement for jets if jets of a given height outnumbered null points at that level.," Still, it would be difficult to invoke null points as a requirement for jets if jets of a given height outnumbered null points at that level."
+ A survey of coronal jets comprehensive enough to determine their Irequency. would thereby cast some light on the importance of null points in their construction., A survey of coronal jets comprehensive enough to determine their frequency would thereby cast some light on the importance of null points in their construction.
+ To estimate the MTE. consider an ideal imagine svstem characterized by Che same point spread function at every point on the image plane aan isoplanatic svstem).," To estimate the MTF, consider an ideal imaging system characterized by the same point spread function at every point on the image plane an isoplanatic system)."
+ In (his case the power spectral density. (psc). SCA). of the image is the product of (the actual ps. Sth). and the square of the modulation transfer ΠΙΠΟΙΟ (MTE). M(/) (BornanclWoll1980).," In this case the power spectral density (psd), $S(k)$, of the image is the product of the actual psd, $\tilde{S}(k)$, and the square of the modulation transfer function (MTF), $M(k)$ \citep{Born1980}."
+. The MTF is the magnitude of the Fourier transform of the normalized point spread function., The MTF is the magnitude of the Fourier transform of the normalized point spread function.
+" Generally, M(7). is a decreasing function of wave-number. causing (he observed. psc. (4). to be sleeper than (the actual one. Sth),"," Generally, $M(k)$, is a decreasing function of wave-number, causing the observed psd, $S(k)$ , to be steeper than the actual one, $\tilde{S}(k)$."
+ If the same field is observed by MDI in both its high and low resolution modesthe, If the same field is observed by MDI in both its high and low resolution modesthe
+sflat spectrum of the⋅ X-ray background.. (XRB) aboveis ‘ 1. keV (Marshall ct al 1980: Cendreau et al 1995: Chen. Fabian Cendreau 1997) is not simply accounted for by the integration. of⋅ known classes of⋅ source. which. generally have much steeper spectra.,"The flat spectrum of the X-ray background (XRB) above 1 keV (Marshall et al 1980; Gendreau et al 1995; Chen, Fabian Gendreau 1997) is not simply accounted for by the integration of known classes of source, which generally have much steeper spectra."
+" A currently. popular model invokes: intrinsic⋠⋠⋠ absorption ""ENin active. galaxies. with. which. to flatten the observed: X-ray spectrum (Setti Woltjer 1989: Macau. CGhisellini Fabian 1994: C'elotti ct al 1995: Comastri et al 1995)."," A currently popular model invokes intrinsic absorption in active galaxies with which to flatten the observed X-ray spectrum (Setti Woltjer 1989; Madau, Ghisellini Fabian 1994; Celotti et al 1995; Comastri et al 1995)."
+" Phe intrinsic absorption must range from column densities""E of⋅ about 10E107cm212. and must cover most (at least ~2/3)' of the sky as seen by the source itself."," The intrinsic absorption must range from column densities of about $10^{22} - 10^{24}\psqcm$, and must cover most (at least $\sim
+2/3$ of the sky as seen by the source itself."
+ The combination of cilferent levels of absorption and. redshift of the objects can then lead. to the observed power-law background spectrum in. the 1LO keV.. band., The combination of different levels of absorption and redshift of the objects can then lead to the observed power-law background spectrum in the 1–10 keV band.
+ qsThe rollover in. the spectrum at about 30. keV e.is due redshift» acting. on an intrinsicMEN LOO keV.: break. such as is. seen in. nearby active. ealaxies (Zdziarski et al 1995).," The rollover in the spectrum at about 30 keV is due redshift acting on an intrinsic 100 keV break, such as is seen in nearby active galaxies (Zdziarski et al 1995)."
+ sPhe geometry of⋅ the obscuring⊀ material⊀ withinsq a typical.. X-ray: backeground-contributing.∢∢ active. galaxy is. unclear.," The geometry of the obscuring material within a typical, X-ray background-contributing, active galaxy is unclear."
+ asThe nucleus must be mostly surrounded. by a typical column. density of sav 107em?., The nucleus must be mostly surrounded by a typical column density of say $10^{23}\psqcm$.
+ M the material freely. orbiting the nucleus. then it should. soon collide Phe with itself. dissipate. and Hatten into a disk. unless either a) the orbits are carefully arranged or b) there is sullicient energv to continually throw matter into a wide range of orbital ⊲⊀⊀⊲inclinations.," If the material is freely orbiting the nucleus, then it should soon collide with itself, dissipate, and flatten into a disk, unless either a) the orbits are carefully arranged or b) there is sufficient energy to continually throw matter into a wide range of orbital inclinations."
+ We. note that what evidence. there isun for a torus of⋅ molecular material. around. active. galaxies. does not indicate that its high column density part covers much of the Sky., We note that what evidence there is for a torus of molecular material around active galaxies does not indicate that its high column density part covers much of the Sky.
+ Indeed. models where the obscuring material is extended on scales LOOpe (Ciranato et al 1997) account for the infrared. properties of Sevfert. galaxies more successfully than ones which postulate a thick compact (<1 pe) torus (Pier Ixrolik. 1992. 1993).," Indeed, models where the obscuring material is extended on scales $\sim 100\, \pc$ (Granato et al 1997) account for the infrared properties of Seyfert galaxies more successfully than ones which postulate a thick compact $<1\, \pc$ ) torus (Pier Krolik, 1992, 1993)."
+ The first case a) may be accounted for by à warped disk., The first case a) may be accounted for by a warped disk.
+ Pringle. (1996). and. Maloney. Begelman Pringle. (1996). have shown that the outer parts of irracliatecl. accretion: disks. are unstable to warping.. the final⋅ result of which. is. that much of. the sky is. covered by the warp.," Pringle (1996) and Maloney, Begelman Pringle (1996) have shown that the outer parts of irradiated accretion disks are unstable to warping, the final result of which is that much of the sky is covered by the warp."
+ uThe details. are currently. uncertain. as to whether most of⋅ the skv can be covered and whether such high. column densitiesu can be attained. in. the warpecl material., The details are currently uncertain as to whether most of the sky can be covered and whether such high column densities can be attained in the warped material.
+. We do not pursue that further here., We do not pursue that further here.
+ Phe second case b) may be accounted for by a nuclear starburst. the supernovae of which can an]provide the energvEZ to 1push clouds ‘around ‘and so obscure the nucleus.," The second case b) may be accounted for by a nuclear starburst, the supernovae of which can provide the energy to push clouds around and so obscure the nucleus."
+ Lt is that possibility we explore further here., It is that possibility we explore further here.
+"clusters form narrow sequences in the color-color and color-age diagrams, in agreement with ?.. (","clusters form narrow sequences in the color-color and color-age diagrams, in agreement with \citet{cervignoluridiana06}. ("
+2.),2.)
+" At intermediate masses, i.e., [Mo|=[3,4], which are typical of the cluster masses observed in extragalactic works, the scatter increases strongly."," At intermediate masses, i.e., $_{\odot}$ ]=[3,4], which are typical of the cluster masses observed in extragalactic works, the scatter increases strongly."
+ The scatter observed in the two-color diagrams for this mass range is similar to what is seen in our observed two-color diagrams. (, The scatter observed in the two-color diagrams for this mass range is similar to what is seen in our observed two-color diagrams. (
+3.),3.)
+" For very low masses (log(Mass)([Mo|=[2]), the scatter again decreases but the model colors now deviate strongly from the SSP colors."," For very low masses $_{\odot}$ ]=[2]), the scatter again decreases but the model colors now deviate strongly from the SSP colors."
+" This is because such clusters have a very low probability of hosting a luminous (but rare and short-lived) post-main sequence star, while the SSP models assume that the colors are an average over all stages of stellar evolution (see ?).. ("," This is because such clusters have a very low probability of hosting a luminous (but rare and short-lived) post-main sequence star, while the SSP models assume that the colors are an average over all stages of stellar evolution \citep[see][]{piskunov09}. . ("
+4.),4.)
+" For ages log(7) S8 and intermediate masses (log(Mass)|Mo]=[3,4]), the color distribution actually becomes bimodal, as observed by ??.."," For ages $\tau$ $\lesssim8$ and intermediate masses $_{\odot}$ ]=[3,4]), the color distribution actually becomes bimodal, as observed by \citet{popescu10a,popescu10}."
+" The blue “peak” in the color distribution is due to clusters without red supergiants, while the presence of even a single red supergiant shifts the colors into the other peak. ("," The blue “peak” in the color distribution is due to clusters without red supergiants, while the presence of even a single red supergiant shifts the colors into the other peak. ("
+5.),5.)
+" The bottom row shows that age estimates are completely dominated by stochastic effects for low-mass clusters (log(Mass)[M]-[2,3])."," The bottom row shows that age estimates are completely dominated by stochastic effects for low-mass clusters $_{\odot}$ ]=[2,3])."
+" For higher masses the ages are more accurately recovered, even if a small scatter is observed compared to the 1:1 relation, especially at ages of a few tens of Myr where the light is strongly dominated by red supergiant stars."," For higher masses the ages are more accurately recovered, even if a small scatter is observed compared to the 1:1 relation, especially at ages of a few tens of Myr where the light is strongly dominated by red supergiant stars."
+" Based on these results it is clear that photometry and the ages estimated from broad band photometry can be heavily affected by the stochastic effects introduced by SIMF sampling, as also shown by ?.."," Based on these results it is clear that photometry and the ages estimated from broad band photometry can be heavily affected by the stochastic effects introduced by SIMF sampling, as also shown by \citet{maizapellaniz09}."
+" We conclude that low-mass clusters (M< 10?Mo) are very strongly affected, reaching age differences up to log(agein)- log(ageout)72.5 dex, for clusters with log(r)[yr]<8.5, while for high-mass clusters (M> 104Mo), this effect gradually diminishes, an effect that is clearly visible as gaps in Fig."," We conclude that low-mass clusters $M \le 10^3 M_{\odot}$ ) are very strongly affected, reaching age differences up to $_{in}$ $_{out})\approx2.5$ dex, for clusters with $\tau$ $\le8.5$, while for high-mass clusters $M \ge 10^4 M_\odot $ ), this effect gradually diminishes, an effect that is clearly visible as gaps in Fig."
+ 8 (last row)., \ref{fig:stoch2c} (last row).
+" Another important effect is observed as a deviation from the 1:1 red line observed for the low-mass clusters, which indicate that the estimated age (Ageour) is again wrongly recovered, even at old ages (1 Gyr)."," Another important effect is observed as a deviation from the 1:1 red line observed for the low-mass clusters, which indicate that the estimated age $Age_{out}$ ) is again wrongly recovered, even at old ages (1 Gyr)."
+" One should be aware of the risk that low-mass, young clusters may erroneously be assigned old ages."," One should be aware of the risk that low-mass, young clusters may erroneously be assigned old ages."
+" When using SSP models to convert their luminosities to masses, based on such wrong age estimates, such objects might be assigned erroneously high masses, thus making it into an observed sample (e.g.?).."," When using SSP models to convert their luminosities to masses, based on such wrong age estimates, such objects might be assigned erroneously high masses, thus making it into an observed sample \citep[e.g.][]{popescu10}."
+ It is important to mention that the SSP models used do not take the binarity or rotation of massive stars into account (seee.g.??)..," It is important to mention that the SSP models used do not take the binarity or rotation of massive stars into account \citep[see e.g.][]{maedermeynet08, eldridgestanway09}."
+" Also, different isochrone assumptions and the techniques used to perform the fit to the ages and masses can be affecting the results (seee.g ??).."," Also, different isochrone assumptions and the techniques used to perform the fit to the ages and masses can be affecting the results \citep[see e.g][]{scheepmaker09,degrijsanders05}."
+" The tests presented in this paper are only intended to address the stochastic sampling effects, and we only rely on Padova isochrones."," The tests presented in this paper are only intended to address the stochastic sampling effects, and we only rely on Padova isochrones."
+" A more detailed study must be made to better account for other effects (e.g. by binaries, rotation, etc)."," A more detailed study must be made to better account for other effects (e.g. by binaries, rotation, etc)."
+ In the next sections we estimate the field star formation histories (SFHs) for each galaxy and determine the ages and masses of the cluster candidates., In the next sections we estimate the field star formation histories (SFHs) for each galaxy and determine the ages and masses of the cluster candidates.
+" We then estimate the cluster formation efficiencies, I."," We then estimate the cluster formation efficiencies, $\Gamma$."
+" To estimate the SFHs, we used the synthetic CMD method."," To estimate the SFHs, we used the synthetic CMD method."
+" We implemented this method using an IDL-based program that was introduced and tested in Paper I. For a description of the program we refer the reader to that paper, but here we summarize the basic functionality."," We implemented this method using an IDL-based program that was introduced and tested in Paper I. For a description of the program we refer the reader to that paper, but here we summarize the basic functionality."
+ The synthetic CMD method (?) takes advantage of the power supplied by the CMDs., The synthetic CMD method \citep{tosi91} takes advantage of the power supplied by the CMDs.
+" The method uses a group of isochrones, together with assumptions about the SIMF, metallicity, distance, extinction, and binarity, to reproduce an observed CMD."," The method uses a group of isochrones, together with assumptions about the SIMF, metallicity, distance, extinction, and binarity, to reproduce an observed CMD."
+ Photometric errors and completeness functions (both being treated as magnitude dependent parameters by our program) are also taken into account., Photometric errors and completeness functions (both being treated as magnitude dependent parameters by our program) are also taken into account.
+" The program searches for the combination of isochrones that best matches the observed CMD, thereby estimating the SFH."," The program searches for the combination of isochrones that best matches the observed CMD, thereby estimating the SFH."
+" The parameters used to estimate the SFH of the galaxies are: a Hess diagram with a resolution of 200x pixels is created, using a Gaussian kernel with a standard deviation of 0.02 magnitudes along the color axis."," The parameters used to estimate the SFH of the galaxies are: a Hess diagram with a resolution of $200\times200$ pixels is created, using a Gaussian kernel with a standard deviation of 0.02 magnitudes along the color axis."
+" The matching is done within the rectangular boxes depicted in Fig. 2,,"," The matching is done within the rectangular boxes depicted in Fig. \ref{fig:hess},"
+" using the V-I vs. V color combination, Padova 2008 isochrones (?) and a? IMF in the mass range 0.1 to 100 Ma."," using the V-I vs. V color combination, Padova 2008 isochrones \citep{marigo08} and a \citet{kroupa02} IMF in the mass range 0.1 to 100 $_{\odot}$."
+" The assumed distance moduli, foreground extinctions, and metallicities are given in Table 1,, while the photometric errors and completeness for each galaxy were determined in Sect."," The assumed distance moduli, foreground extinctions, and metallicities are given in Table \ref{tab:setparam}, while the photometric errors and completeness for each galaxy were determined in Sect."
+ 4., 4.
+" The program also includes a simplified treatment of binaries, in which binary evolution is ignored, but the effect of unresolved binaries on the CMD are modeled."," The program also includes a simplified treatment of binaries, in which binary evolution is ignored, but the effect of unresolved binaries on the CMD are modeled."
+ To account for binarity we used three different assumptions for the binary fraction (f) and mass ratio (q): (1.), To account for binarity we used three different assumptions for the binary fraction $f$ ) and mass ratio $q$ ): (1.)
+" f=0.0 and q= 0.0, (2.)"," $f=0.0$ and $q=0.0$ , (2.)"
+" f=0.5 and q=(0.1,0.9] (assuming a flat distribution), and (3.)"," $f=0.5$ and $q=[0.1,0.9]$ (assuming a flat distribution), and (3.)"
+ f=1 and q=1., $f=1$ and $q=1$.
+ These three assumptions are the same ones as used in Paper I. Figures 9 shows the best-fit Hess diagrams., These three assumptions are the same ones as used in Paper I. Figures \ref{fig:fit_sfhs} shows the best-fit Hess diagrams.
+" Comparing with the observed Hess diagrams (Fig. 2)),"," Comparing with the observed Hess diagrams (Fig. \ref{fig:hess}) ),"
+ we see that the fits are far from perfect., we see that the fits are far from perfect.
+" In particular, all the model Hess diagrams show a clear separation between the blue core He burning (“blue loop”) stars and the main sequence, while this is not obvious in most of the observed diagrams."," In particular, all the model Hess diagrams show a clear separation between the blue core He burning (“blue loop”) stars and the main sequence, while this is not obvious in most of the observed diagrams."
+" This might be partly due to some variation in the internal extinction, which has not been included in our modeling."," This might be partly due to some variation in the internal extinction, which has not been included in our modeling."
+" To infer the SFHs of our sample, we combined all the fields (per galaxy)and passed to our program."," To infer the SFHs of our sample, we combined all the fields (per galaxy)and passed to our program."
+" The star formation rates, normalized to unit area, are shown in Figure 10 and the average values are listed inTable 5 for ages between 10 and 100 Myrs."," The star formation rates, normalized to unit area, are shown in Figure \ref{fig:sfhs} and the average values are listed inTable \ref{tab:sfrcfr} for ages between 10 and 100 Myrs."
+" In this age range, our data"," In this age range, our data"
+"An extensive. literature of analytical ancl numerical investigations of shoek-cloucl interactions now exists in which the clleets of magnetic fields. radiative cooling. hermal conduction and turbulence have all been considered (sce.e.g.Pittarcdetal.2009:παν,Llartquist&Falle2010.anclreferences. therein)...","An extensive literature of analytical and numerical investigations of shock-cloud interactions now exists in which the effects of magnetic fields, radiative cooling, thermal conduction and turbulence have all been considered \citep*[see, e.g.,][and references
+ therein]{Pittard:2009,Pittard:2010}."
+ However. in. many astrophysical objects the cloud size is comparable to or arecr than the depth of the post-shock laver. so idt is surprising that to date almost all numerical investigations nave been performed. in the small-cloud limit where the yost-shock Low has effectively. an infinite depth.," However, in many astrophysical objects the cloud size is comparable to or larger than the depth of the post-shock layer, so it is surprising that to date almost all numerical investigations have been performed in the small-cloud limit where the post-shock flow has effectively an infinite depth."
+" Notable exceptions include investigations where the global Low is simulated. for instance. in the interaction of a supernova remnant (SNR) with a cloud. Clenorio-Lagle&Rozvezka1986:Rozvezka&""Tenorio-lagle1987:Leàoetal. 2009).. though such models suller from low numerical resolution."," Notable exceptions include investigations where the global flow is simulated, for instance, in the interaction of a supernova remnant (SNR) with a cloud \citep{Tenorio-Tagle:1986,Rozyczka:1987,Leao:2009}, though such models suffer from low numerical resolution."
+ Ligher resolution was used in work on the triggered collapse of a cloud by a thin shell (Bossetal.2010).. but the downstream [low was not examined.," Higher resolution was used in work on the triggered collapse of a cloud by a thin shell \citep{Boss:2010}, but the downstream flow was not examined."
+ In this paper we re-examine the interaction of a cold. dense. isothermal shell with a spherical cloud. studying the effect. of the shell thickness ancl Mach number.," In this paper we re-examine the interaction of a cold, dense, isothermal shell with a spherical cloud, studying the effect of the shell thickness and Mach number."
+ The shell is driven by a high pressure. low density. hot bubble.," The shell is driven by a high pressure, low density, hot bubble."
+ Pressure-driven. shells occur in. planetary. nebulae (PNe) ancl around. individual stars or groups of massive stars. ancl have diameters ranging from 0.1 pe to a few kpe.," Pressure-driven shells occur in planetary nebulae (PNe) and around individual stars or groups of massive stars, and have diameters ranging from $\sim 0.1\;$ pc to a few kpc."
+ Since cooling breaks the scale-free nature of acliabatic simulations one is forced to choose a particular lengthscaloe., Since cooling breaks the scale-free nature of adiabatic simulations one is forced to choose a particular lengthscale.
+ In this paper the interaction of a superbubble with a molecular cloud is simulated. though the results are also relevant to the smaller-scale interactions mentioned above.," In this paper the interaction of a superbubble with a molecular cloud is simulated, though the results are also relevant to the smaller-scale interactions mentioned above."
+ The structure. of the superbubble. including the hickness of the shell and the density contrast between the shell and the hot interior. can be obtained [rom a &loba welrodyvnamical simulation.," The structure of the superbubble, including the thickness of the shell and the density contrast between the shell and the hot interior, can be obtained from a global hydrodynamical simulation."
+ However. the shell properties depend on assumptions about the structure of the LSA (c.g.. whether there is a density eracient. the strength anc orientation of magnetic fields. the number. size and density distributions of clouds. etc.)," However, the shell properties depend on assumptions about the structure of the ISM (e.g., whether there is a density gradient, the strength and orientation of magnetic fields, the number, size and density distributions of clouds, etc.)"
+ and the ellicieney of mass ransler from the cool shell and from clouds within the ho xibble to the hot. rarefied gas.," and the efficiency of mass transfer from the cool shell and from clouds within the hot bubble to the hot, rarefied gas."
+ Therefore. we choose insteac o specify instantaneous shell and bubble properties (Mach number. surface density. bubble temperature) in order to ocus on the Κον physics of the interaction.," Therefore, we choose instead to specify instantaneous shell and bubble properties (Mach number, surface density, bubble temperature) in order to focus on the key physics of the interaction."
+ Ehe shell is assumed: to be planar (i.e. its radius is much larger than he cloud)., The shell is assumed to be planar (i.e. its radius is much larger than the cloud).
+ Our investigation is based around a standard: moclel with the following parameters., Our investigation is based around a standard model with the following parameters.
+" The inter-cloud LSAL has an assumed density and temperature of. pus=3.83 and Zia,=SOOOIx."," The inter-cloud ISM has an assumed density and temperature of $\rho_{\rm amb} = 3.33 \times 10^{-25}\;\gpcm3$ and $T_{\rm
+ amb}=8000\;{\rm K}$."
+ The density contrast of the spherical cloud X=107. and its core temperature is d;=Shk.," The density contrast of the spherical cloud $\chi = 10^{3}$, and its core temperature is $T_{\rm c} = 8\;{\rm
+ K}$."
+ The cloud is thus in pressure equilibrium with its surroundings., The cloud is thus in pressure equilibrium with its surroundings.
+ To keep matters as simple as possible in this first investigation all material is assumed to behave as an ideal gas with a ratio of specific heats. 5=5/3. although in reality the cores are molecular.," To keep matters as simple as possible in this first investigation all material is assumed to behave as an ideal gas with a ratio of specific heats, $\gamma=5/3$, although in reality the cores are molecular."
+ This simplification has little effect. on the overall dynamics., This simplification has little effect on the overall dynamics.
+ Ehe. LSAL pressure Duas=2667em I., The ISM pressure $p_{\rm amb}=2667 \;\pcm3{\rm K}$ .
+" The cloud has a mass AM,=300M. ", The cloud has a mass $M_{\rm c}=300 \Msol$ 
+luminosity. function. or may be contaminated by Sevlert 5galaxies.,"luminosity function, or may be contaminated by Seyfert galaxies."
+ This justifies our choice of a low-frequeney selected: sub-m.Jy. radio sample. albeit in retrospect.," This justifies our choice of a low-frequency selected sub-mJy radio sample, albeit in retrospect."
+ We have already. discussed. that. the D93. sample is insullicient. to determine. g(L) at. faint radio uminosities since the volume and thus number of objects at low redshift is small., We have already discussed that the B93 sample is insufficient to determine $\varphi(L)$ at faint radio luminosities since the volume and thus number of objects at low redshift is small.
+ To estimate the corrections to our estimate of he luminosity densitwv. (equation 3)) we are only interested. in the shape of the luminosity function (where shape includes a measure. of the typica Luminosity e.g. £.) and not the overall normalisation which cancels in Equations 4. and 5..., To estimate the corrections to our estimate of the luminosity density (Equation \ref{lumden}) ) we are only interested in the shape of the luminosity function (where shape includes a measure of the typical Luminosity e.g. $L_{*}$ ) and not the overall normalisation which cancels in Equations \ref{eqn:compl} and \ref{ldist}.
+ We thus wish o test the shape of BOS luminosity function agains some reasonable a priori functions., We thus wish to test the shape of B93 luminosity function against some reasonable a priori functions.
+ Vhe LRAS luminosity function is well determined even at the faint end. it is natural to test an IUS luminosity function modified via the ELIt/radio correlation (Llelou. Soifer RowanRobinson 1985).," The IRAS luminosity function is well determined even at the faint end, it is natural to test an IRAS luminosity function modified via the FIR/radio correlation (Helou, Soifer Rowan–Robinson 1985)."
+" lor à galaxy with the model star-burst spectrum of Efstathiou. RowanRobinson Seibenmorgen (1998. in preparation) this relation can be translated to OL We consider two mocoels. the first. (/1) is the ""warm luminosity function of S90(their solution 25). (modified so L.=133Ls04)."," For a galaxy with the model star-burst spectrum of Efstathiou, Rowan–Robinson Seibenmorgen (1998, in preparation) this relation can be translated to or We consider two models, the first $l1$ ) is the “warm” luminosity function of S90(their solution 25), (modified so $L_{*}=133 L_{60*}$ )."
+ Ht is appropriate to use the “warm” sample. since these are the actively star-forming galaxies. similar to this sample.," It is appropriate to use the “warm” sample, since these are the actively star-forming galaxies, similar to this sample."
+" ""This uminositv function is also clearly similar in shape and normalisation to our observed. radio luminosity unction (Figure 3)).", This luminosity function is also clearly similar in shape and normalisation to our observed radio luminosity function (Figure \ref{fig:radlf}) ).
+ Phe use of this model luminosity unction is independent. of the overall normalisation. and assumes the shape is constant with recshilt.," The use of this model luminosity function is independent of the overall normalisation, and assumes the shape is constant with redshift."
+ lt is known that LRAS galaxies strongly evolve (c.g. SOO. Oliver£994. ete).," It is known that IRAS galaxies strongly evolve (e.g. S90, Oliver, etc.),"
+ ancl if. this evolution only changes the normalisation of the uminosity Function then this is immaterial., and if this evolution only changes the normalisation of the luminosity function then this is immaterial.
+ However if shape of the luminosity function. was to change with redshift. this would. alfect our estimate of he luminosity density.," However if shape of the luminosity function was to change with redshift, this would affect our estimate of the luminosity density."
+ Neither LAS samples. nor his sample alone are sullicient to determine. what sort of evolution is appropriate to these Ίνρος of samples: however the faint radio counts do not support pure density evolution as strong as that sugeested in S90. (Rowan-Robinson 1993).," Neither IRAS samples, nor this sample alone are sufficient to determine what sort of evolution is appropriate to these types of samples; however the faint radio counts do not support pure density evolution as strong as that suggested in S90 (Rowan-Robinson 1993)."
+ Thus it is appropriate to consider some mocel with an evolution of the Luminosity with redshift., Thus it is appropriate to consider some model with an evolution of the Luminosity with redshift.
+ Our second model (2 assumes the same zero recishift luminosity function. as /1 but has L=Lo(l:3d JU., Our second model $l2$ assumes the same zero redshift luminosity function as $l1$ but has $L=L_{0}(1+z)^{3.1}$ .
+ Since the S90 2=O luminosity function was not calculated. with this evolution. this moclel is not entirely consistent and. will over-prediet. the numbers of luminous galaxies. and lead to an under-correction. to the Luminosity censity. nevertheless he two models together give a reasonable idea. of he level of uncertainty in the luminosity function. and would bracket a more. natural mocel where Luminosity evolution was included in the estimate of he Luminosity function self consistentIv.," Since the S90 $z=0$ luminosity function was not calculated with this evolution, this model is not entirely consistent and will over-predict the numbers of luminous galaxies, and lead to an under-correction to the Luminosity density, nevertheless the two models together give a reasonable idea of the level of uncertainty in the luminosity function, and would bracket a more natural model where Luminosity evolution was included in the estimate of the Luminosity function self consistently."
+ Our third model. 9. assumes the shape of the ow 2 Conconaf (LOST) radio Luminosity [unction. with no evolution.," Our third model, $l3$, assumes the shape of the low $z$ Condon (1987) radio luminosity function, with no evolution."
+ In addition to the luminosity functions we consider a number of dillerent models for the elfect of the joint optical and radio limits on the observable racio Iuminosities Z5(O|L):, In addition to the luminosity functions we consider a number of different models for the effect of the joint optical and radio limits on the observable radio luminosities $P_i({\cal O}|L)$ :
+We investigate the ionising contribution of quasars ancl the 2]-cm power spectrum as a function of redshift based. on the parameterisecl mocdel in Equation (11)). together with the halo mass function and a minimum halo mass Alywin. ‘Table 1 gives the parameter values used in our simulations. which correspond to the plots in Figure 3. which shows the resulting contributions of stars and quasars to reionisation.,"We investigate the ionising contribution of quasars and the 21-cm power spectrum as a function of redshift based on the parameterised model in Equation \ref{Rq}) ), together with the halo mass function and a minimum halo mass $M_{\rm h,min}$ Table \ref{tab:param} gives the parameter values used in our simulations, which correspond to the plots in Figure \ref{Qevolution} which shows the resulting contributions of stars and quasars to reionisation."
+ ligure 3. shows that the assumption of a lower minimum halo mass allows quasars to contribute more significantly at earlier phases of reionisation., Figure \ref{Qevolution} shows that the assumption of a lower minimum halo mass allows quasars to contribute more significantly at earlier phases of reionisation.
+ Figure 4 shows the relative cilferential ionised fraction due to quasars (ΑΙ)dQu/dAly and the cumulative ionised. fraction due το quasars Qy/Qil<My) as a 'unction of My for simulations S183., Figure \ref{QvsM} shows the relative differential ionised fraction due to quasars $(M_{\rm h}/Q_{\star})dQ_{\rm q}/dM_{\rm h}$ and the cumulative ionised fraction due to quasars $Q_{\rm q}/Q_{\star}(<M_{\rm h})$ as a function of $M_{\rm h}$ for simulations S1–S3.
+ The dillerential— plots demonstrate that the bull of the ionisation by quasars is achieved. by quasars in. lower-mass host halos. despite arecr bubbles being generated by quasars in halos of higher mass.," The differential plots demonstrate that the bulk of the ionisation by quasars is achieved by quasars in lower-mass host halos, despite larger bubbles being generated by quasars in halos of higher mass."
+ This is due to the greater population of Lower-niass ios., This is due to the greater population of lower-mass halos.
+ As mentioned above. the parameter sets used. for hese simulations have been chosen so as to produce a cumulative contribution to reionisation by quasars of 10 »er cent.," As mentioned above, the parameter sets used for these simulations have been chosen so as to produce a cumulative contribution to reionisation by quasars of 10 per cent."
+ We also carried out simulations with a cumulative relative Ionisation contribution by quasars of z 3 per cent., We also carried out simulations with a cumulative relative ionisation contribution by quasars of $\approx$ 3 per cent.
+ The results of these simulations (not presented) demonstrate hat the elfect on the resulting power spectra scale in »oportion to quasar contribution., The results of these simulations (not presented) demonstrate that the effect on the resulting power spectra scale in proportion to quasar contribution.
+ As mentioned. in Section l.. statistical observations of the epoch of reionisation promise to provide a wealth of information about the properties of neutral hydrogen at high. redshift as well as some of the fundamental astrophysics behind. the retonisation process and. the first luminous objects.," As mentioned in Section \ref{Introduction}, statistical observations of the epoch of reionisation promise to provide a wealth of information about the properties of neutral hydrogen at high redshift as well as some of the fundamental astrophysics behind the reionisation process and the first luminous objects."
+" While density perturbations in the matter distribution mediate Ductuations in the 2l-cm signal both prior to and following reionisation. during the reionisation era the relation. between the 21-cm. power spectrum. and the underlying matter power spectrum is complex and. in its late stages. is dominated by the formation of large ionisecl ""bubbles? (??).."," While density perturbations in the matter distribution mediate fluctuations in the 21-cm signal both prior to and following reionisation, during the reionisation era the relation between the 21-cm power spectrum and the underlying matter power spectrum is complex and, in its late stages, is dominated by the formation of large ionised “bubbles"" \citep{furl2004b,mcquinn2006}."
+ Phese bubbles of ionised hydrogen imprint features on the 2l-em power spectrum that reflect the luminosity and clustering of ionising sources responsible [or reionisation., These bubbles of ionised hydrogen imprint features on the 21-cm power spectrum that reflect the luminosity and clustering of ionising sources responsible for reionisation.
+ The 21-cm power spectrum can be used in a statistical test to distinguish. candidate. reionisation models. as well as to constrain the history and morphology of reionisation (?)..," The 21-cm power spectrum can be used in a statistical test to distinguish candidate reionisation models, as well as to constrain the history and morphology of reionisation \citep{barkana2008}."
+ In this paper. we emplov the climensionless power spectrum AGQ(2) as the Κον statistical measure. which is the contribution to the variance of the redshifted 21-cm brightness temperature contrast 971(2) per logarithmic interval in wave-number &=27/A.," In this paper, we employ the dimensionless power spectrum $\Delta_{21}^2(k,z)$ as the key statistical measure, which is the contribution to the variance of the redshifted 21-cm brightness temperature contrast $\delta T_{\rm b}(z)$ per logarithmic interval in wave-number $k = 2\pi/\lambda$."
+ This measure is related to the dimensional form of the power spectrum £2.(4.2) by Po(h.2) is estimated by averaging over all mn niodes of the Fourier transform. (2) of 22) in a thin spherical shell in A-space. Note that we must use caution when emploving this statistical measure since the spherical svmumetry of the signal is broken by redshift evolution over certain ranges of redshift.," This measure is related to the dimensional form of the power spectrum $P_{21}(k,z)$ by $P_{21}(k,z)$ is estimated by averaging over all $m$ modes of the Fourier transform $\hat{T}$ ) of $T(z)$ in a thin spherical shell in $k$ -space, Note that we must use caution when employing this statistical measure since the spherical symmetry of the signal is broken by redshift evolution over certain ranges of redshift."
+ This issue has been ciscussed by a number of authors (see.e.g..ον) and its ellect on the sensitivity of the 21-cm. power spectrum has been investigated. by 2.," This issue has been discussed by a number of authors \citep[see, e.g.,][]{morales2004,bl2005,mcquinn2006} and its effect on the sensitivity of the 21-cm power spectrum has been investigated by \cite{mcquinn2006}."
+ Our numerical power spectrum measurements are subject to sample variance. arising from the finite number of independent modes. counted. in. cach &-shell.. which corresponds to the finite number of independent wavelengths A=οπής that can fit into the simulated. volume.We begin our analvsis by presenting one example of the," Our numerical power spectrum measurements are subject to sample variance, arising from the finite number of independent modes counted in each $k$ -shell, which corresponds to the finite number of independent wavelengths $\lambda = 2\pi/k$ that can fit into the simulated .We begin our analysis by presenting one example of the"
+in the MW.,in the MW.
+" Thus, the rescaling of cross-section of dust by the metal mass using the MW value results in significant overestimate for H2 formation in high-redshift galaxies."," Thus, the rescaling of cross-section of dust by the metal mass using the MW value results in significant overestimate for ${\rm H}_{2}$ formation in high-redshift galaxies."
+ Dust grains injected into the ISM are subjected to destruction by the blast waves (the high-velocity interstellar shocks) driven by the ambient SNe (e.g.Jonesetal.1994)., Dust grains injected into the ISM are subjected to destruction by the blast waves (the high-velocity interstellar shocks) driven by the ambient SNe \citep[e.g.][]{Jon94}.
+. Nozawaetal.(2006) investigated the processing of interstellar dust by sputtering in the hot gas swept up by the SN forward shock., \citet{Noz06} investigated the processing of interstellar dust by sputtering in the hot gas swept up by the SN forward shock.
+" Adopting the dust model by Nozawaetal.(2003) as the size distribution of interstellar dust, Nozawaetal.(2006) have shown that the destruction efficiency of dust depends on the ISM density and the explosion energy of SNe as well as the initial size distribution of dust."," Adopting the dust model by \citet{Noz03} as the size distribution of interstellar dust, \citet{Noz06} have shown that the destruction efficiency of dust depends on the ISM density and the explosion energy of SNe as well as the initial size distribution of dust."
+ It should be noted that the size distribution as well as the destruction efficiency changes as a function of time because interstellar dust are supplied from SNe and processed in ISM successively according to star formation activity., It should be noted that the size distribution as well as the destruction efficiency changes as a function of time because interstellar dust are supplied from SNe and processed in ISM successively according to star formation activity.
+" Thus, we must deal with the destruction process in a way that is applicable to any dust size distribution to explore the global evolution of dust size distribution."," Thus, we must deal with the destruction process in a way that is applicable to any dust size distribution to explore the global evolution of dust size distribution."
+" In order to evaluate the destruction efficiency of interstellar dust for any initial size distribution, here we introduce the conversion efficiency as defined below."," In order to evaluate the destruction efficiency of interstellar dust for any initial size distribution, here we introduce the conversion efficiency as defined below."
+" Consider that the j-th dust species residing in the ISM, whose size distribution is given by the number of dust grains with radii between a and a4-da, f;(a)da, is processed by sputtering in hot plasma produced through the a single passage of SN shock."," Consider that the $j$ -th dust species residing in the ISM, whose size distribution is given by the number of dust grains with radii between $a$ and $a+{\rm d}a$, $f_{j}(a){\rm d}a$, is processed by sputtering in hot plasma produced through the a single passage of SN shock."
+" The conversion efficiency 7;(a,a') is defined as the number fraction of dust grains with radii between a’ and a’+da’ that are converted to grains with radii between a and a+da by sputtering through the passage of a SN shock."," The conversion efficiency $\eta_{j}(a,a^{\prime})$ is defined as the number fraction of dust grains with radii between $a^{\prime}$ and $a^{\prime}+{\rm d}a^{\prime}$ that are converted to grains with radii between $a$ and $a+{\rm d}a$ by sputtering through the passage of a SN shock."
+" The number of dust grains with radii between a and a+da produced by the sputtering is given as n;(a,a’)f;(a’)da’."," The number of dust grains with radii between $a$ and $a+{\rm d}a$ produced by the sputtering is given as $\eta_{j}(a,a^{\prime})f_{j}(a^{\prime}){\rm d}a^{\prime}$."
+" Note that nj(a,a’)=0 for a>a’."," Note that $\eta_{j}(a,a^{\prime})=0$ for $a>a^{\prime}$."
+" Then the change in the number of dust grains with radii between a and à4-da caused by a shock processing is given by and as well the corresponding change of the mass is given by where Ma,;(a)da is the mass of the pre-shocked dust."," Then the change in the number of dust grains with radii between $a$ and $a+{\rm d}a$ caused by a shock processing is given by and as well the corresponding change of the mass is given by where ${\cal M}_{{\rm d},j}(a){\rm d}a$ is the mass of the pre-shocked dust."
+ The size distribution function after the shock processing f/(a) is given by f;(a)=ία)+dN; /da., The size distribution function after the shock processing $f^{\prime}_{j}(a)$ is given by $f^{\prime}_{j}(a)=f_{j}(a)+{\rm d}N_{j}/{\rm d}a$ .
+" The conversion efficiency 7(a,a’) and the mass of ISM gas swept up by shock Mswept depend on the progenitor mass, expanding energy and type of SN as well as the structure, number density"," The conversion efficiency $\eta(a,a^{\prime})$ and the mass of ISM gas swept up by shock $M_{\rm swept}$ depend on the progenitor mass, expanding energy and type of SN as well as the structure, number density"
+velocity. aud two material parameters that define the catastrophic disruption criteria.,"velocity, and two material parameters that define the catastrophic disruption criteria."
+ Catastrophic disruption is defined bv the specific impact cucrey needed to disperse half the total mass., Catastrophic disruption is defined by the specific impact energy needed to disperse half the total mass.
+ For collisions that result In partial accretion or erosion of differeutiated bodies. we estimate the chausge iu the coreanantle mass ratio following the work of 2010).," For collisions that result in partial accretion or erosion of differentiated bodies, we estimate the change in the core-mantle mass ratio following the work of ."
+ The new analytic model now allows plauct formation siuulatious to capture the full diversity of collision outconies., The new analytic model now allows planet formation simulations to capture the full diversity of collision outcomes.
+ To assess the effects of incliding more realistic collision physics in planet formation studies. we analyzed the Hupact parameters from recent N-body simulations of terrestrial planet formation that asstmed perfec moreie aud calculated the amore realistic collision outcomes2009).," To assess the effects of including more realistic collision physics in planet formation studies, we analyzed the impact parameters from recent $N$ -body simulations of terrestrial planet formation that assumed perfect merging and calculated the more realistic collision outcomes."
+. The results provide iusights iuto the most nmuportau collision reguues. quantify the probabilities of differeu outcomes. aud identify future challenges in plane formation simulations.," The results provide insights into the most important collision regimes, quantify the probabilities of different outcomes, and identify future challenges in planet formation simulations."
+ Then. using the distribution of inipact parameters frou the N-body simulations. we modeled individual plane erowtl with a Monte Carlo technique in order to assess the cumulative effects of diverse collision outcomes.," Then, using the distribution of impact parameters from the $N$ -body simulations, we modeled individual planet growth with a Monte Carlo technique in order to assess the cumulative effects of diverse collision outcomes."
+ We compare the distribution of final planet masses from siuulatious with different assuuptious about collision outcomes and tracked changes iu the core mass fraction., We compare the distribution of final planet masses from simulations with different assumptions about collision outcomes and tracked changes in the core mass fraction.
+ The collision plivsics model allows for more detailed investigation of the evolution of the bulk chemistiy and for tighter isotopic constraints on terrestrial planet formation studies., The collision physics model allows for more detailed investigation of the evolution of the bulk chemistry and for tighter isotopic constraints on terrestrial planet formation studies.
+ The outcome of a collision in the eravitv regime was predicted using the new analytic collision pliysics model from(2012). hereafter LS12.," The outcome of a collision in the gravity regime was predicted using the new analytic collision physics model from, hereafter LS12."
+ The model is sunuuanzed here: refer to LS12 for the full details aud limitations., The model is summarized here; refer to LS12 for the full details and limitations.
+ Iu the gravitv reeiue. erosive outcomes are primarily controlled by overcoming eravitational forces rather than overcoming material streneth: bodies larger than about 100 to 1000 12 in size are dominated by eravity1997).," In the gravity regime, erosive outcomes are primarily controlled by overcoming gravitational forces rather than overcoming material strength; bodies larger than about 100 to 1000 m in size are dominated by gravity."
+. The model predicts the collision outcome as a fiction of projectile-to-target mass ratio (7= Ace). impact anele (0). aud impact velocity (1;).," The model predicts the collision outcome as a function of projectile-to-target mass ratio $\gamma=M_{\rm
+ p}/M_{\rm t}$ ), impact angle $\theta$ ), and impact velocity $V_i$ )."
+" By AL,/convention, the projectile is less massive tha thetarget (AL, Af)."," By convention, the projectile is less massive than the target $M_{\rm p} \le M_{\rm t}$ )."
+ The following equations were used to ecucrate theboundaries< and contours on the example collision outcome maps shown in Figure 1authors..., The following equations were used to generate the boundaries and contours on the example collision outcome maps shown in Figure \ref{fig:collmap}.
+ Partial aceretion and crosive outcomes are defined by the catastrophic disruption criteria (Q5). which is the specific impact energv required to disrupt and eravitationally disperse half the total mass.," Partial accretion and erosive outcomes are defined by the catastrophic disruption criteria $Q^*_{RD}$ ), which is the specific impact energy required to disrupt and gravitationally disperse half the total mass."
+" For a particular collision. the total specific impact enerev in the center of mass frame is given by where Mor=My,| aud pp=ALAEMice is the reduced mass."," For a particular collision, the total specific impact energy in the center of mass frame is given by where $M_{\rm tot}=M_{\rm p}+M_{\rm t}$ and $\mu=M_{\rm p}M_{\rm
+ t}/M_{\rm tot}$ is the reduced mass."
+ For bodies in the eravity rege. the head-onu disruption criteria increases with increasing mass aud follows the general form derived by(19903:: where Rey is the spherical radius of the combined projectile and target masses at a density of pj=1000 ke inP and G is the evavitational coustant.," For bodies in the gravity regime, the head-on disruption criteria increases with increasing mass and follows the general form derived by: where $R_{C1}$ is the spherical radius of the combined projectile and target masses at a density of $\rho_1 \equiv 1000$ kg $^{-3}$ and $G$ is the gravitational constant."
+ At the critical iupact velocity (V;=V). the specific impact cherey is equal to the catastrophic disruption threshold (Qiο for the particular iipaet scenario (Rey aud lass ratio).," At the critical impact velocity $(V_i=V^*)$, the specific impact energy is equal to the catastrophic disruption threshold $(Q_R=Q^*_{RD})$ for the particular impact scenario $R_{C1}$ and mass ratio)."
+ Instead of target radius. the variable Rey was introduced by in order to generalize the dependence of 6ης ou size for different density bodies aud for different projectile-to-target mass ratios.," Instead of target radius, the variable $R_{C1}$ was introduced by in order to generalize the dependence of $Q^*_{RD}$ on size for different density bodies and for different projectile-to-target mass ratios."
+ The cocfitcicut yg is of order 1 aud defined below., The coefficient $q_g$ is of order 1 and defined below.
+ fis the velocity exponcut in the point source coupling paralucter from(1990)., $\bar \mu$ is the velocity exponent in the point source coupling parameter from.
+". LS12 fit a value of jj0.3640.01 using the results frou, several munerical studies= ofdisruption iu the gravity roeinie on a wide variety of materials Gveals rock. ice. rubble piles)."," LS12 fit a value of $\bar \mu=0.36\pm0.01$ using the results from several numerical studies of disruption in the gravity regime on a wide variety of materials (weak rock, ice, rubble piles)."
+ LS12 derived the disruption criteria for a reference case of a head-on impact between two equalanass bodies. which represents the mininmni energy scenario to disrupt a particular target.," LS12 derived the disruption criteria for a reference case of a head-on impact between two equal-mass bodies, which represents the minimum energy scenario to disrupt a particular target."
+ The critical impact velocity for on equal mass collisions is given by where e? is a noudimnieusional material parameter that is fitted to simulation data., The critical impact velocity for head-on equal mass collisions is given by where $c^*$ is a nondimensional material parameter that is fitted to simulation data.
+ LS12 fit e=1940.3 for fiuid planets aud e?=5432 for various solid planctesimal analogs using the results frou several uuuerical studies., LS12 fit $c^*=1.9\pm0.3$ for fluid planets and $c^*=5\pm2$ for various solid planetesimal analogs using the results from several numerical studies.
+ eis defined such that it represcuts the value of the equalmass head-on diguption criteria in units of the specific eravitational binding οσον of the combined mass: Equation [ is named the “principal disruption curve.”, $c^*$ is defined such that it represents the value of the equal-mass head-on disruption criteria in units of the specific gravitational binding energy of the combined mass: Equation \ref{eqn:principalq} is named the “principal disruption curve.”
+ Combining the previous equations eives the cocficicut for Equation 2: The derivation of Equation 2. assiuned that all of the projectile’s kinetic energy is deposited in the target., Combining the previous equations gives the coefficient for Equation \ref{eqn:qstarred}: The derivation of Equation \ref{eqn:qstarred} assumed that all of the projectile's kinetic energy is deposited in the target.
+ This assunption is often invalid as a portion of the projectile niüsses the target for many combinations of projectile-to-target nass ratio and impact auele., This assumption is often invalid as a portion of the projectile misses the target for many combinations of projectile-to-target mass ratio and impact angle.
+ For au oblique iupact. LS12 derived the mass fraction of the projectile that gcometrically overlaps with the target. where pis the density of the projectile. b=siu is the impact parzuneter (5=0 for head-on impacts) and (20) is the fraction of the projectile diameter that overlaps the target.," For an oblique impact, LS12 derived the mass fraction of the projectile that geometrically overlaps with the target, where $\rho$ is the density of the projectile, $b=\sin\theta$ is the impact parameter $b=0$ for head-on impacts) and $l/(2R_{\rm p})$ is the fraction of the projectile diameter that overlaps the target."
+ The mipact angle and overlapping ass fraction are defined at the point of first contact assunudue spherical bodies., The impact angle and overlapping mass fraction are defined at the point of first contact assuming spherical bodies.
+ Using just the interacting, Using just the interacting
+"Ih bands. the secondary maximum peak is at 7, = 13.1 and A, = 12.1 ancl (he underlviug (ix magnitudes are J, = 14.3 and A, = I4.1.","IR bands, the secondary maximum peak is at $J_o$ = 13.1 and $K_o$ = 12.1 and the underlying flux magnitudes are $J_o$ = 14.3 and $K_o$ = 14.1."
+ All magnitudes were converted to [Iuxes using zeropoint [Iuxes given by Allen(2000) ancl Campins.Rieke.andLebolsky(1985)., All magnitudes were converted to fluxes using zeropoint fluxes given by \cite{all01} and \cite{cam85}.
+. In each band. the underlying flux was subtracted [rom the peak flux eiving the residual fIux. corresponding to the flux of the secondary maximum only.," In each band, the underlying flux was subtracted from the peak flux giving the residual flux, corresponding to the flux of the secondary maximum only."
+ The SED of the secondary maximum [lis al ils peak is presented in Figure 13.., The SED of the secondary maximum flux at its peak is presented in Figure \ref{fig:sed}.
+" The top panel shows the data in log E, v. log v space and the bottom panel in log pE, v. log » space.", The top panel shows the data in log $_{\nu}$ v. log ${\nu}$ space and the bottom panel in log $\nu$ $_{\nu}$ v. log ${\nu}$ space.
+ In the top panel the optical and II. points lie almost on a straight line., In the top panel the optical and IR points lie almost on a straight line.
+ There is clearly a break somewhere between the A-band point and (he radio but it is impossible to sav where exacilv., There is clearly a break somewhere between the $K$ -band point and the radio but it is impossible to say where exactly.
+ In the bottom panel we can see that most of the power is in the OUR. bands., In the bottom panel we can see that most of the power is in the OIR bands.
+ We attempted to fit the SED with a single blackbodsy. peaking at or around the A-band point since (his has the greatest [Iux.," We attempted to fit the SED with a single blackbody, peaking at or around the $K$ -band point since this has the greatest flux."
+ Two example fits are shown in Figure 1H4.., Two example fits are shown in Figure \ref{fig:sed_onebb}.
+" The solid line for Ty, = 1350IX peaks at Ix. while the dashed line is the best fit to the 1. J and Ix data. wilh a blackbody temperature of Ty, = 16507Ix. Neither of these models fit the optical data."," The solid line for $_{bb}$ = $^o$ K peaks at K, while the dashed line is the best fit to the I, J and K data, with a blackbody temperature of $_{bb}$ = $^o$ K. Neither of these models fit the optical data."
+ For a temperature of 10500. the emitting region would need to have a radius of ~ 1.4 x 10cm given a distance of 9.1 kpe (Oroszetal.1998)..," For a temperature of $^o$ K, the emitting region would need to have a radius of $\sim$ 1.4 x $^{17}$ cm given a distance of 9.1 kpc \citep{oro98}."
+ For an orbital period of 1.123 davs (Oroszοἱal.L998).. the size of the primary Roche lobe is ~ 5 x LOtem.," For an orbital period of 1.123 days \citep{oro98}, the size of the primary Roche lobe is $\sim$ 5 x $^{11}$ cm."
+ Hence this innnedialely rules out anoptically Chick accretion disk as the source of the observed flare radiation., Hence this immediately rules out anoptically thick accretion disk as the source of the observed flare radiation.
+ The SED was fitted with multicolor-blackbody models parameter definitions).., The SED was fitted with multicolor-blackbody models \citep[see][for parameter definitions]{mit84}.
+" We fixed (he inner radius. rj, = g corresponding to the innermost stable orbit of a black hole (where ry, is the gravitational radius)."," We fixed the inner radius, $_{in}$ = $_g$, corresponding to the innermost stable orbit of a black hole (where $_g$ is the gravitational radius)."
+ The mass of the compact object. (black hole) was fixed at 5 M. (Oroszetal.1998)., The mass of the compact object (black hole) was fixed at 5 $_{\odot}$ \citep{oro98}.
+". In Figure 15. we show two fits: the first. (solid line) is for a disk with rj; = 107"" rj, (= 4.9 x 10cm) and M = 10! L. the second (dashed line) is for r4, = 1053 rj, (= L1 x 10cm) and M = 10? |"," In Figure \ref{fig:sed_multi} we show two fits: the first (solid line) is for a disk with $_{out}$ = $^{7.04}$ $_{in}$ (= 4.9 x $^{13}$ cm) and $\dot{M}$ = $^{16}$ $^{-1}$, the second (dashed line) is for $_{out}$ = $^{4.4}$ $_{in}$ (= 1.1 x $^{11}$ cm) and $\dot{M}$ = $^{8.3}$ $^{-1}$."
+ Although the models fit the /-. J- and A-band points. they cannot account lor the B and V points.," Although the models fit the $I$ -, $J$ - and $K$ -band points, they cannot account for the $B$ and $V$ points."
+" In (he first case. r;,;~ LOO times the primary Roche lobe."," In the first case, $_{out} \sim$ 100 times the primary Roche lobe."
+ In the second. case. ro; is within the primary Roche lobe but M is fav less than the average observed accretion rates of X-ray binaries. that is ~ LOP-LO | (Frank.King.andRaine 2002)..," In the second case, $_{out}$ is within the primary Roche lobe but $\dot{M}$ is far less than the average observed accretion rates of X-ray binaries, that is $\sim$ $^{15}$ $^{18}$ $^{-1}$ \citep{fra02}. ."
+ Hence multicolor-blackbody models cannot account lor the secondary maximum (hus., Hence multicolor-blackbody models cannot account for the secondary maximum flux.
+ We fitted multicolor-blackbody models to the underlving outburst [ux to eheck that it could be explained bv an optically (hick accretion disk., We fitted multicolor-blackbody models to the underlying outburst flux to check that it could be explained by an optically thick accretion disk.
+ We chose (o examine data al two, We chose to examine data at two
+"SK acknowledges support from the Royal Commission for the Exhibition of 1851, Imperial College and Worcester College, Oxford.","SK acknowledges support from the Royal Commission for the Exhibition of 1851, Imperial College and Worcester College, Oxford."
+ We thank David Clements and Yvonne Unruh for many constructive comments., We thank David Clements and Yvonne Unruh for many constructive comments.
+" This publication has been made possible by the participation of more than 300,000 volunteers in the Galaxy Zoo project."," This publication has been made possible by the participation of more than 300,000 volunteers in the Galaxy Zoo project."
+ Their contributions are individually acknowledged at http://www.galaxyzoo.org/Volunteers.aspx., Their contributions are individually acknowledged at http://www.galaxyzoo.org/Volunteers.aspx.
+in the z and = sample. respectively.,"in the $-z$ and $-z$ sample, respectively."
+ Note that these values do not include any correction by the baryonic depletion that is expected to raise the gas/baryon fraction by about 8/6 per cent at this overdensity., Note that these values do not include any correction by the baryonic depletion that is expected to raise the gas/baryon fraction by about 8/6 per cent at this overdensity.
+ See also Fig., See also Fig.
+ | and text that follows)., \ref{fig:dat} and text that follows).
+ These estimates are consistent with other. independent. recent determinations (e.g. Allen et al.," These estimates are consistent with other, independent, recent determinations (e.g. Allen et al."
+ 2001. Pratt Arnaud 2002) and consistently lower than the baryon budget required from WMAP results.," 2001, Pratt Arnaud 2002) and consistently lower than the baryon budget required from WMAP results."
+ It is worth noticing that only the highest estimates of the gas (baryon) mass fraction (e.g. A426. A2142. RXT1350) are perfectly consistent with WMAP results. whereas the other clusters are systematically below them.," It is worth noticing that only the highest estimates of the gas (baryon) mass fraction (e.g. A426, A2142, RXJ1350) are perfectly consistent with WMAP results, whereas the other clusters are systematically below them."
+ To investigate the systematics that could affect this estimate. I change the values of the factors 2 and C' and study their influence on the baryon fraction.," To investigate the systematics that could affect this estimate, I change the values of the factors $B$ and $C$ and study their influence on the baryon fraction."
+ The factor. D. which parametrizes the uncertainties on Adj.) IS expected to be between | and 1.15 from the cluster mass profiles recovered from both X-ray and lensing data (e.g. Allen et al.," The factor $B$ , which parametrizes the uncertainties on $M_{\rm tot}$ is expected to be between 1 and 1.15 from the cluster mass profiles recovered from both X-ray and lensing data (e.g. Allen et al."
+ 2001)., 2001).
+ The factor C' represents the level of clumpiness that affects the estimate of Ad... in X-ray analysis and that simulations show to be lower than 1.2 (Mathiesenet al., The factor $C$ represents the level of clumpiness that affects the estimate of $M_{\rm gas}$ in X-ray analysis and that simulations show to be lower than 1.2 (Mathiesenet al.
+ 1999)., 1999).
+Having established the basic stellar population properties of the short GRB hosts in our sample we now address the fundamental questions listed in refsec:intro through a comparison to the general galaxy mass function and the hosts of long GRBs.,Having established the basic stellar population properties of the short GRB hosts in our sample we now address the fundamental questions listed in \\ref{sec:intro} through a comparison to the general galaxy mass function and the hosts of long GRBs.
+ The question of whether the short GRB rate tracks stellar mass alone can be addressed through two statistical tests: (1) the fraction of short GRBs in early- and late-type galaxies: and (11) the relation of the host stellar masses to the general galaxy mass function., The question of whether the short GRB rate tracks stellar mass alone can be addressed through two statistical tests: (i) the fraction of short GRBs in early- and late-type galaxies; and (ii) the relation of the host stellar masses to the general galaxy mass function.
+ In the redshift range under consideration here (z~0.2— 1) about an equal fraction of the cosmic stellar mass budget resides in early- and late-type galaxies (e.g.. [bertetal. 2010).," In the redshift range under consideration here $z\sim 0.2-1$ ) about an equal fraction of the cosmic stellar mass budget resides in early- and late-type galaxies (e.g., \citealt{isl+10}) )."
+ Therefore. if short GRBs track stellar mass alone we expect a roughly one-to-one ratio of galaxy types.," Therefore, if short GRBs track stellar mass alone we expect a roughly one-to-one ratio of galaxy types."
+ This does not appear to be the case., This does not appear to be the case.
+ For example. within the sample of short GRBs with optical afterglows (19 events). only 2 are unambiguously hosted by early-type galaxies (GRBs 050724 and 100117: Bergeretal.2005.. Fong et al.," For example, within the sample of short GRBs with optical afterglows (19 events), only 2 are unambiguously hosted by early-type galaxies (GRBs 050724 and 100117; \citealt{bpc+05}, Fong et al."
+ in prep.).," in prep.),"
+ while 8 are unambiguously hosted by late-type galaxies: the probability of obtaining this ratio from an intrinsic one-to-one distribution is only 0.04., while 8 are unambiguously hosted by late-type galaxies; the probability of obtaining this ratio from an intrinsic one-to-one distribution is only 0.04.
+ The identity of the remaining 9 hosts ts unclear at the present due to their faintness or the lack of underlying galaxies at the burst positions., The identity of the remaining 9 hosts is unclear at the present due to their faintness or the lack of underlying galaxies at the burst positions.
+ Still. unless nearly all of these bursts were hosted by early-type galaxies. the resulting ratio appears to be skewed in favor of late-type host galaxies with on-going star formation activity.," Still, unless nearly all of these bursts were hosted by early-type galaxies, the resulting ratio appears to be skewed in favor of late-type host galaxies with on-going star formation activity."
+ We note that the same result holds true if we consider the bursts with only X-ray afterglow positions and identified hosts., We note that the same result holds true if we consider the bursts with only X-ray afterglow positions and identified hosts.
+ Thus. the host galaxy demographics suggest that short GRBs do not track stellar mass alone. or phrased alternatively. they do not have a delay time distribution that is skewed to old ages of ~few Gyr.," Thus, the host galaxy demographics suggest that short GRBs do not track stellar mass alone, or phrased alternatively, they do not have a delay time distribution that is skewed to old ages of $\sim {\rm few}$ Gyr."
+ It is possible. however. that this result is influenced by secondary factors such as the typical circumburst density or intrinsic differences in the energy scale and afterglow brightness as a function of galaxy type (possibly reminiscent of the differences in peak luminosity for Type Ia supernovae in early- and late-type galaxies; Hamuyetal.2000:Mannucei 2006)).," It is possible, however, that this result is influenced by secondary factors such as the typical circumburst density or intrinsic differences in the energy scale and afterglow brightness as a function of galaxy type (possibly reminiscent of the differences in peak luminosity for Type Ia supernovae in early- and late-type galaxies; \citealt{htp+00,mdp06}) )."
+ If such differences lead to fainter afterglows (or prompt emission) for short GRBs in early-type galaxies. this would suppress the early-type fraction.," If such differences lead to fainter afterglows (or prompt emission) for short GRBs in early-type galaxies, this would suppress the early-type fraction."
+ Although the modest size of the host sample. and the substantial fraction of short GRBs with only 7-ray positions (~1/3 of all events). prevent definitive conclusions. it does not appear that the optical afterglows of short GRBs in early- and late-type galaxies are distinct (Berger2010a).," Although the modest size of the host sample, and the substantial fraction of short GRBs with only $\gamma$ -ray positions $\sim 1/3$ of all events), prevent definitive conclusions, it does not appear that the optical afterglows of short GRBs in early- and late-type galaxies are distinct \citep{ber10}."
+. Thus. our preliminary conclusion from. the host demographics is that short GRBs do not appear to track the cosmic stellar mass fraction in. early- and late-type galaxies.," Thus, our preliminary conclusion from the host demographics is that short GRBs do not appear to track the cosmic stellar mass fraction in early- and late-type galaxies."
+ We next turn to a comparison of the inferred stellar masses with the galaxy mass function., We next turn to a comparison of the inferred stellar masses with the galaxy mass function.
+ The cumulative distribution of stellar masses for the short GRB hosts 1s shown in Figure 5. with the range of possible masses bounded by the single age SSP and maximal values., The cumulative distribution of stellar masses for the short GRB hosts is shown in Figure \ref{fig:masses_2} with the range of possible masses bounded by the single age SSP and maximal values.
+ We also present a breakdown of the sample into early- and late-type galaxies. each spanning the same range.," We also present a breakdown of the sample into early- and late-type galaxies, each spanning the same range."
+ For the late-type hosts. the intermediate young+told values are also shown.," For the late-type hosts, the intermediate young+old values are also shown."
+" To compare these distributions to the distribution of galaxy masses we also plot the cumulative stellar mass functionmass. Le.. the fraction of stellar mass in galaxies above some mass. FO»M). given by the equation: where (M) is the Schechter mass function: We use several determinations of (M) in this comparison.including the Coleetal.(2001) mass function from the 2MASS/2dF catalogs for all galaxy types at z0 10/75 Mo. 621.18: the nearly identical Panteretal.(2004) mass function from SDSS for all galaxy types at z0 (M*=10! M... à2-1.16): the Belletal.(2003). mass function for late-type galaxies from 2MASS/SDSS converted to a Salpeter IMF for comparison with our inferred. values (M*=10!) M... à21.27: and the Ibertetal.(2010) mass functions from the COSMOS survey for quiescent galaxies at z—0.3 (M=10! M.. a=-091) and intermediate-activity galaxies at z—0.5 (M=10! ,, o-—].02). matched to the redshifts of the early- and late- short GRB host galaxies in our sample."," To compare these distributions to the distribution of galaxy masses we also plot the cumulative stellar mass function, i.e., the fraction of stellar mass in galaxies above some mass, $f(>M)$ , given by the equation: where $\Phi(M)$ is the Schechter mass function: We use several determinations of $\Phi(M)$ in this comparison,including the \citet{cnb+01} mass function from the 2MASS/2dF catalogs for all galaxy types at $z\sim 0$ $M^*=10^{11.16}$ $_\odot$, $\alpha=-1.18$ ); the nearly identical \citet{phj04} mass function from SDSS for all galaxy types at $z\sim 0$ $M^*=10^{11.19}$ $_\odot$, $\alpha=-1.16$ ); the \citet{bmk+03} mass function for late-type galaxies from 2MASS/SDSS converted to a Salpeter IMF for comparison with our inferred values $M^*=10^{10.97}$ $_\odot$, $\alpha=-1.27$ ); and the \citet{isl+10} mass functions from the COSMOS survey for quiescent galaxies at $z\sim 0.3$ $M^*=10^{11.13}$ $_\odot$, $\alpha=-0.91$ ) and intermediate-activity galaxies at $z\sim 0.5$ $M^*=10^{10.93}$ $_\odot$, $\alpha=-1.02$ ), matched to the redshifts of the early- and late-type short GRB host galaxies in our sample."
+ The resulting distributions of fC»M) for the various mass functions are shown in Figure 5.., The resulting distributions of $f(>M)$ for the various mass functions are shown in Figure \ref{fig:masses_2}.
+ The median of each distribution is close to the value of M. but also depends on the value of & (Coleetal.2001: (M;=1095 M.; Panteretal.2004: (M)=10? .: Belletal.2003: (M5=10/79! .; Ilbertetal.2010:: (M)=10!79 ΜΜ... for intermediate-activity and (M=10! .. for quiescent).," The median of each distribution is close to the value of $M^*$, but also depends on the value of $\alpha$ \citealt{cnb+01}: $\langle M^*\rangle=10^{10.88}$ $_\odot$; \citealt{phj04}: $\langle M^*\rangle=10^{10.92}$ $_\odot$; \citealt{bmk+03}: $\langle M^*\rangle=10^{10.61}$ $_\odot$; \citealt{isl+10}: $\langle M^*\rangle=10^{10.76}$ $_\odot$ for intermediate-activity and $\langle M^*\rangle=10^{11.02}$ $_\odot$ for quiescent)."
+ Ins refsec:res we demonstrated that the median single age SSP mass for the short GRB host sample is 1.3«I0 M. z0.2(M5. well below the expected value if the short GRB rate tracked stellar mass.," In \\ref{sec:res} we demonstrated that the median single age SSP mass for the short GRB host sample is $1.3\times 10^{10}$ $_\odot$, $\approx 0.2\langle M^*\rangle$, well below the expected value if the short GRB rate tracked stellar mass."
+ The result is even more discrepant if we restrict the sample to only the late-type hosts. with a median of about 0.06:M5.," The result is even more discrepant if we restrict the sample to only the late-type hosts, with a median of about $0.06\langle M^*\rangle$."
+ On the other hand. for the early-type hosts we find a median of about 1.60M suggesting that they are drawn from the early-type galaxy nass function.," On the other hand, for the early-type hosts we find a median of about $1.6\langle M^*\rangle$, suggesting that they are drawn from the early-type galaxy mass function."
+ Using the maximal masses for the full sample we find a nedian stellar mass of about 1.20M5. while for the subset of late-type hosts the median is about 0.5(M5: the maximal nasses of the early-type hosts lead to a large median mass of about 5(M). suggesting that they are likely over-estimates of the true stellar mass.," Using the maximal masses for the full sample we find a median stellar mass of about $1.2\langle M^*\rangle$, while for the subset of late-type hosts the median is about $0.5\langle M^*\rangle$; the maximal masses of the early-type hosts lead to a large median mass of about $5\langle M^*\rangle$, suggesting that they are likely over-estimates of the true stellar mass."
+ Finally. using the young+old stellar nasses for the late-type hosts the median mass is about 0.3(M5.," Finally, using the young+old stellar masses for the late-type hosts the median mass is about $0.3\langle M^*\rangle$."
+ Thus. we find that the distribution of stellar masses for the late-type hosts 1s shifted to à somewhat lower value than the expected median of the galaxy mass function. whilethe early-type hosts have à similar median to the galaxy mass function.," Thus, we find that the distribution of stellar masses for the late-type hosts is shifted to a somewhat lower value than the expected median of the galaxy mass function, whilethe early-type hosts have a similar median to the galaxy mass function."
+ Beyond the comparison of median values. we further assess the agreement (or lack thereof) between the short GRB hosts and galaxy mass functions using the Kolmogorov-Smirnov (K-S) test.," Beyond the comparison of median values, we further assess the agreement (or lack thereof) between the short GRB hosts and galaxy mass functions using the Kolmogorov-Smirnov (K-S) test."
+ For the full sample there is negligible probability that the distribution of single age SSP masses is drawn from the galaxy mass function. with Pz8« 107.," For the full sample there is negligible probability that the distribution of single age SSP masses is drawn from the galaxy mass function, with $P\approx
+8\times 10^{-5}$ ."
+ On the other, On the other
+bulk fluid speed with a gradient of velocity in the precursor region is Uy>U0.,bulk fluid speed with a gradient of velocity in the precursor region is $U_{0}>U>0$.
+" Because of the preseribed scattering law instead of the particle’s movement in the fields, the injected particles from the thermalized downstream into the precursor for diffusive processes are controlled by the free clastic scatter mechanism."," Because of the prescribed scattering law instead of the particle's movement in the fields, the injected particles from the thermalized downstream into the precursor for diffusive processes are controlled by the free elastic scatter mechanism."
+" To obtain the information ofthe total particles in different regions at any time, we build a database for recording the velocities, positions, and time of the all particles, as well as the index and the bulk speeds of the grids."," To obtain the information of the total particles in different regions at any time, we build a database for recording the velocities, positions, and time of the all particles, as well as the index and the bulk speeds of the grids."
+" The scattering angle distributions are presented by Gaussian distribution function with a standard deviation 6, and an average value gd involving four cases: (1) Case A: 6= q/4, u= 0. ("," The scattering angle distributions are presented by Gaussian distribution function with a standard deviation $\sigma$, and an average value $\mu$ involving four cases: (1) Case A: $\sigma=\pi$ /4, $\mu=0$ . ("
+2) Case B: 6=x/2. u=0. (,"2) Case B: $\sigma=\pi$ /2, $\mu=0$. ("
+"3) Case C: o=x, u=0. (","3) Case C: $\sigma=\pi$, $\mu=0$. ("
+"4) Case D: 6=ee, ui=0.","4) Case D: $\sigma=\infty$, $\mu=0$."
+ These presented simulations are all based on a one-dimensional simulation box and the specific parameters are based on Wang&Yan(2011)., These presented simulations are all based on a one-dimensional simulation box and the specific parameters are based on \cite{wang11}.
+ We present the entire shock evolution with the velocity profiles of the time sequences in each case as shown in Figure |.., We present the entire shock evolution with the velocity profiles of the time sequences in each case as shown in Figure \ref{fig:shock}.
+ The total velocity profiles are divided into three parts with respect to the shock front and the FEB locations., The total velocity profiles are divided into three parts with respect to the shock front and the FEB locations.
+" The upstream bulk speed Uy dynamically slows down by passing though the precursor region, and its value decreases to zero in the downstream region (1.c. U,= 0)."," The upstream bulk speed $U_{0}$ dynamically slows down by passing though the precursor region, and its value decreases to zero in the downstream region (i.e. $U_{d}=0$ )."
+ The precursor explicitly shows a different slope of the bulk velocity and different final FEB locations in different cases., The precursor explicitly shows a different slope of the bulk velocity and different final FEB locations in different cases.
+ The present velocity profiles are similar to the density profiles in the previous simulations by Wang&Yan(201I).. and the different prescribed scattering law leads to the different shock structures.," The present velocity profiles are similar to the density profiles in the previous simulations by \citet{wang11}, and the different prescribed scattering law leads to the different shock structures."
+ The various losses of the particles and the calculated results of the shocks at the end of the simulation for the four cases are listed in Table 1.., The various losses of the particles and the calculated results of the shocks at the end of the simulation for the four cases are listed in Table \ref{tab:res}.
+ The initial box energy is Ey., The initial box energy is $E_{0}$.
+" The subshock's compression ratio 7,5 and the total compression ratio r,,, are calculated from the fine velocity structures in the shock frame in each case, respectively."," The subshock's compression ratio $r_{sub}$ and the total compression ratio $r_{tot}$ are calculated from the fine velocity structures in the shock frame in each case, respectively."
+" The total energy spectral index I;,, and the subshock's energy spectral index Ty, are deduced from the corresponding total compression ratio 7;,, and the subshock's compression ratio ry, in cach case.", The total energy spectral index $\Gamma_{tot} $ and the subshock's energy spectral index $\Gamma_{sub} $ are deduced from the corresponding total compression ratio $r_{tot}$ and the subshock's compression ratio $r_{sub}$ in each case.
+" The Mipsy. Proxy and Ej,,, are the mass, momentum, and energy losses which are produced by the escaped particles via the FEB, respectively."," The $M_{loss}$, $P_{loss}$ and $E_{loss}$ are the mass, momentum, and energy losses which are produced by the escaped particles via the FEB, respectively."
+" We have monitored the mass, momentum and energy of the total particles in cach time step in each case."," We have monitored the mass, momentum and energy of the total particles in each time step in each case."
+ Figure 2. shows all the types of energy functions with respect to time., Figure \ref{fig:eng} shows all the types of energy functions with respect to time.
+ The total energy Ey 1s the energy summation over the time in the entire simulation box at any instant in time., The total energy $E_{tot}$ is the energy summation over the time in the entire simulation box at any instant in time.
+" The box energy Ej, presents the actual energy in the simulationbox at any instant in time.", The box energy $E_{box}$ presents the actual energy in the simulationbox at any instant in time.
+ The supplement energy Epig is the summation of the amount of energy from the pre-inflow box (at the, The supplement energy $E_{PIB}$ is the summation of the amount of energy from the pre-inflow box (at the
+third light 7 flux's uncertainty.,third light $I$ flux's uncertainty.
+ The formal error of modified potentials ts at the level of 0.04 and the rest is the contribution from the third light., The formal error of modified potentials is at the level of 0.04 and the rest is the contribution from the third light.
+ As for ASAS-04. changing temperatures by 200 K did not affect the potentials significantly.," As for ASAS-04, changing temperatures by 200 K did not affect the potentials significantly."
+ Uncertainty in the absolute photometric calibration was however added to the colour and the absolute magnitudes error., Uncertainty in the absolute photometric calibration was however added to the colour and the absolute magnitudes error.
+ Absolute bolometric and V magnitudes were again calculated with JKTABSDIM for the two cases with and without spots., Absolute bolometric and $V$ magnitudes were again calculated with JKTABSDIM for the two cases with and without spots.
+ In the case with spots. we found the distance to ASAS-08 to be 44 + 9 pc. which agrees very well with the one from (45.8+4.7 pe).," In the case with spots, we found the distance to ASAS-08 to be 44 $\pm$ 9 pc, which agrees very well with the one from $45.8 \pm 4.7$ pc)."
+ On this basis. we also calculated. M for the third star.," On this basis, we also calculated $M_V$ for the third star."
+ Its uncertainty given in Table 5 includes the measurement error. whichis however not incorporated into the V—/ error because the colour is calculated strarghtforwardly from the model.," Its uncertainty given in Table \ref{tab_3rd} includes the measurement error, whichis however not incorporated into the $V-I$ error because the colour is calculated straightforwardly from the model."
+ The error of the absolute photometric calibration its added to both mentioned parameters., The error of the absolute photometric calibration is added to both mentioned parameters.
+ We also note that values of Mj; and My for the binary’s components are again consistent within error bars but. as for ASAS-04. we decided to use the “rectified” ones in further analysis.," We also note that values of $M_{bol}$ and $M_V$ for the binary's components are again consistent within error bars but, as for ASAS-04, we decided to use the ”rectified” ones in further analysis."
+ In order to show the influence of spots in both investigated systems quantitatively. in we plot Figure 4+ the absolute difference in magnitudes between our final models and models of binaries with the same absolute physical parameters but without spots.," In order to show the influence of spots in both investigated systems quantitatively, in we plot Figure \ref{fig_spot} the absolute difference in magnitudes between our final models and models of binaries with the same absolute physical parameters but without spots."
+ Additionally we plot differential RV curves constructed with PHOEBE as a difference between the spotted and unspotted model with the eclipse proximity effects turned off in both cases so that the Rossiter-McLaughlin effect is not shown., Additionally we plot differential RV curves constructed with PHOEBE as a difference between the spotted and unspotted model with the eclipse proximity effects turned off in both cases so that the Rossiter-McLaughlin effect is not shown.
+ One can deduce that for ASAS-08 a spotted surface is seen all the time because the spot on the primary covers the polar region., One can deduce that for ASAS-08 a spotted surface is seen all the time because the spot on the primary covers the polar region.
+ The magnitude of the light variation — 0.03—0.05 — is not very high., The magnitude of the light variation – $0.03 - 0.05$ – is not very high.
+ It is in deed lower than the peak-to-peak scatter of the ASAS measurements. so from the ASAS data only it would be difficult to deduce the presence of spots in ou systems.," It is in deed lower than the peak-to-peak scatter of the ASAS measurements, so from the ASAS data only it would be difficult to deduce the presence of spots in our systems."
+ Especially if the spots were moving during the nin0 years of observations. which is presumably possible.," Especially if the spots were moving during the nine years of observations, which is presumably possible."
+ For thisnr reason our RV measurements could have been obtained wheο the spot configuration was different., For this reason our RV measurements could have been obtained when the spot configuration was different.
+ We cannot say much abot= the spot's pattern at the moment of spectroscopic observation., We cannot say much about the spot's pattern at the moment of spectroscopic observation.
+ From the Fig., From the Fig.
+ 4. we can however estimate that the RV curve's modulation is smaller than about 0.5 km s! for the observed configuration of spots in our photometric data set., \ref{fig_spot} we can however estimate that the RV curve's modulation is smaller than about 0.5 km $^{-1}$ for the observed configuration of spots in our photometric data set.
+ Obviously. this is within the total RV errorbars.," Obviously, this is within the total RV errorbars."
+ We also found that turning on the proximity effects while fitting to RVs in PHOEBE does not significantly change the results and in the case of 08 makes the rns of the RV fit a bit larger., We also found that turning on the proximity effects while fitting to RVs in PHOEBE does not significantly change the results and in the case of ASAS-08 makes the $rms$ of the RV fit a bit larger.
+ We believe the total error of the RV measurements does account for the overall activity of the targets and therefore the resulting uncertainties of the best-fitting parameters are realistic., We believe the total error of the RV measurements does account for the overall activity of the targets and therefore the resulting uncertainties of the best-fitting parameters are realistic.
+ Figure 4. also shows three-dimensional reconstructions of ASAS-04 and ASAS-08 in two phases — @=0.02 (primary eclipse) and @=0.48 (secondary eclipse)., Figure \ref{fig_spot} also shows three-dimensional reconstructions of ASAS-04 and ASAS-08 in two phases – $\phi=0.02$ (primary eclipse) and $\phi=0.48$ (secondary eclipse).
+ Values on axis are solar radii and stars are shown on the same seale., Values on axis are solar radii and stars are shown on the same scale.
+ The properties of the spots can be found to be slightly puzzling., The properties of the spots can be found to be slightly puzzling.
+ Among several late-type detached eclipsing binaries known to have spots. for example20093).. (Kron1952:Torres&Ribas2002).. (Ribas2003).. (López-Morales 2005).. (Vaccaroetal.2007) or (Shkolniketal.2008).. none is reported to be covered to an extent such large as ASAS-08.," Among several late-type detached eclipsing binaries known to have spots, for example, \citep{kro52,tor02}, \citep{rib03}, \citep{lop05}, , \citep{vac07} or \citep{shk08}, none is reported to be covered to an extent such large as ASAS-08."
+ Some of seven eclipsing binaries listec by Coughlin&Shaw(2007) have more than half of their surface covered by spots. but their real masses are only estimated from the evolutionary models.," Some of seven eclipsing binaries listed by \citet{cou07} have more than half of their surface covered by spots, but their real masses are only estimated from the evolutionary models."
+" In some other cases. like GU Boo (López-Morales&Ribas2005).. the temperature factor is higher than 1 (a ""bright"" spot)."," In some other cases, like GU Boo \citep{lop05}, the temperature factor is higher than 1 (a ”bright” spot)."
+ It is not clear if the existence of such a feature is physically valid. but an explanation given by López-Morales&Ribas(2005) suggests that almost whole star's surface is covered by many small spots of high contrast and only a small fraction of “clear” photosphere is seen.," It is not clear if the existence of such a feature is physically valid, but an explanation given by \citet{lop05} suggests that almost whole star's surface is covered by many small spots of high contrast and only a small fraction of ”clear” photosphere is seen."
+ This seems to be the case of ASAS-08 as well especially having the temperature factors close to 1., This seems to be the case of ASAS-08 as well especially having the temperature factors close to 1.
+ As mentioned above. radii of both ASAS-04 components and their higher temperature indicate late main-sequence or even à post-main-sequence evolutionary stage.," As mentioned above, radii of both ASAS-04 components and their higher temperature indicate late main-sequence or even a post-main-sequence evolutionary stage."
+ Inflated radii are commonly observed in low-mass eclipsing binaries. but together with lower temperatures (Torres&Ribas2002:Chabrieretal. 2007)... which for ASAS-04 seem to. be higher than predicted.," Inflated radii are commonly observed in low-mass eclipsing binaries, but together with lower temperatures \citep{tor02,cha07}, which for ASAS-04 seem to be higher than predicted."
+ This may still be the effect of the absolute magnitude calibration. but it would mean that the real temperatures differ by almost 3o.," This may still be the effect of the absolute magnitude calibration, but it would mean that the real temperatures differ by almost $\sigma$."
+ Keeping in mind the uncertainty in the temperatures one cannot be convinced that ASAS-04 is indeed an old system., Keeping in mind the uncertainty in the temperatures one cannot be convinced that ASAS-04 is indeed an old system.
+ But other observational facts seem to support this hypothesis., But other observational facts seem to support this hypothesis.
+ Young objects of this mass should exhibit quite a high level of activity., Young objects of this mass should exhibit quite a high level of activity.
+ A substantial emission in X-rays and UV should be observed and a strong emission in spectral lines should also be present., A substantial emission in X-rays and UV should be observed and a strong emission in spectral lines should also be present.
+ No significant UV emission is seen around ASAS-O4's position. and at a distance of 137 a faint X-ray source can be found -- J045304.," No significant UV emission is seen around ASAS-04's position, and at a distance of 13” a faint X-ray source can be found -- ."
+1-070011.. López-Morales(2007) found a correlation between the activity level and radii of low-mass stars in close binaries., \citet{lop07} found a correlation between the activity level and radii of low-mass stars in close binaries.
+ The fractional deviation of a star’s radius as compared to | Gyr. Z= 0.02. Ξ1.0 BCAH98 isochrone (Baratffeetal.1998) for main-sequence stars turned out to be a function of the X-ray-to-bolometrie luminosity ratio — Ly/Lj.," The fractional deviation of a star's radius as compared to 1 Gyr, $Z = 0.02$ , $\alpha = 1.0$ BCAH98 isochrone \citep{bar98} for main-sequence stars turned out to be a function of the X-ray-to-bolometric luminosity ratio – $L_X/L_{bol}$."
+" For a linear correlation of Ly and rotational velocity (asfoundbyFlemingetal. 1989).. and stars M.«0.7 M... the relation was where Rup. Is the observed radius and Aj, Is the radius predicted by the BCAH98 isochrone for a given mass."," For a linear correlation of $L_X$ and rotational velocity \citep[as found by][]{fle89}, and stars $M<0.7$ $_{\odot}$, the relation was where $R_{obs}$ is the observed radius and $R_{mod}$ is the radius predicted by the BCAH98 isochrone for a given mass."
+ When stars with ή>0.7 M. were included. the slope changed to 62.511 (López-Morales 2007)..," When stars with $M>0.7$ $_{\odot}$ were included, the slope changed to 62.511 \citep{lop07}. ."
+ The uncertainties of the slope’s values were not given. nor was the zero-point of therelation. but from Fig.," The uncertainties of the slope's values were not given, nor was the zero-point of therelation, but from Fig."
+ 3., 3.
+ in Lopez-Morales(2007) one can deduce that itwas about 0.005 and 0.015 for the cases without and with the more Massive stars respectively., in \citet{lop07} one can deduce that itwas about 0.005 and 0.015 for the cases without and with the more massive stars respectively.
+The polar field is defined as the average over a polar cap of 15? latitude extension.,The polar field is defined as the average over a polar cap of $15^{\circ}$ latitude extension.
+ The total surface flux is obtained by integrating the unsigned magnetic field over the whole surface., The total surface flux is obtained by integrating the unsigned magnetic field over the whole surface.
+" Here we describe how we model the magnetic flux source term, S(A,¢,t), in Equation (1))."," Here we describe how we model the magnetic flux source term, $S(\lambda,\phi,t)$, in Equation \ref{eqn:SFT}) )."
+" We use the semi-synthetic sunspot group records determined in Paper I on the basis of either the group sunspot number, Rg (?) or the Wolf sunspot number, Rz, the 12-month running averages of which are shown in Figure 1. from 1700 onwards."," We use the semi-synthetic sunspot group records determined in Paper I on the basis of either the group sunspot number, $R_G$ \citep{Hoyt98} or the Wolf sunspot number, $R_Z$, the 12-month running averages of which are shown in Figure \ref{fig:WolfVSgroup} from 1700 onwards."
+ Each sunspot group in the semi-synthetic record is taken to represent a bipolar magnetic region (BMR)., Each sunspot group in the semi-synthetic record is taken to represent a bipolar magnetic region (BMR).
+" The area of a BMR, being the sum of the umbral, penumbral and facular areas, is calculated from the area of the corresponding sunspot group using the empirical relation found by ?.."," The area of a BMR, being the sum of the umbral, penumbral and facular areas, is calculated from the area of the corresponding sunspot group using the empirical relation found by \citet{Chapman97}."
+" The tilt angle of the BMR, a, is assumed to be of that of the sunspot group."," The tilt angle of the BMR, $\alpha$, is assumed to be of that of the sunspot group."
+ This factor (70%)) was empirically determined in CJSS10 and reflects the effect of the localized inflows into active regions., This factor ) was empirically determined in CJSS10 and reflects the effect of the localized inflows into active regions.
+" In Paper I, the tendency of sunspots to appear in activity nests was measured by considering which combination of uniformly distributed random emergence longitudes and completely ordered longitudes match the degree of non-randomness in the sunspot record of the Royal Greenwich Observatory since 1874."," In Paper I, the tendency of sunspots to appear in activity nests was measured by considering which combination of uniformly distributed random emergence longitudes and completely ordered longitudes match the degree of non-randomness in the sunspot record of the Royal Greenwich Observatory since 1874."
+ Here we follow the same approach and use both the semi-synthetic datasets with ordered and random emergence longitudes., Here we follow the same approach and use both the semi-synthetic datasets with ordered and random emergence longitudes.
+ How the resulting fields are combined is discussed in Section 3., How the resulting fields are combined is discussed in Section 3.
+" The magnetic field distribution of each new BMR is given by Β(1.4)=B*(A,d)—B""(A,Φ) with where 6.(A,@) are the heliocentric angles between the location of the sunspot group,(A,9), and the centers of each polarity, (44,94)."," The magnetic field distribution of each new BMR is given by $B(\lambda,\phi)=B^+(\lambda,\phi)-B^-(\lambda,\phi)$ with where $\beta_{\pm}(\lambda,\phi)$ are the heliocentric angles between the location of the sunspot $(\lambda,\phi)$, and the centers of each polarity, $(\lambda_{\pm},\phi_{\pm})$."
+ A8 is the separation between the two polarities.," $\Delta
+\beta$ is the separation between the two polarities."
+" We take AB=0.45A)/2 where Ag is the total area of the sunspot groups, and 6=4? (??).."," We take $\Delta\beta=0.45A_R^{1/2}$, where $A_R$ is the total area of the sunspot groups, and $\delta=4^\circ$ \citep{van_Ballegooijen98,
+Baumann04}."
+ Bmax is the peak field strength of a BMR., $B_\mathrm{max}$ is the peak field strength of a BMR.
+" Following CJSS10, we take Bmax=374 G. AB and the tilt angle, a, determine the latitudinal separation, Afsina, of the two polarities of a BMR."," Following CJSS10, we take $B_\mathrm{max}=374$ G. $\Delta\beta$ and the tilt angle, $\alpha$, determine the latitudinal separation, $\Delta\beta\sin\alpha$, of the two polarities of a BMR."
+ This quantity is important for the axial dipole moment of the global magnetic field and affects the polar field and the open flux during activity minima., This quantity is important for the axial dipole moment of the global magnetic field and affects the polar field and the open flux during activity minima.
+" The longitudinal separation, given by ABcosa(cos1):1, dominates the equatorial dipole moment and the open flux during activity maximum phases."," The longitudinal separation, given by $\Delta\beta\cos\alpha(\cos\lambda)^{-1}$, dominates the equatorial dipole moment and the open flux during activity maximum phases."
+" Associated with the n BMR in the semi-synthetic sunspot record is the time £, when the BMR appears.", Associated with the $n^{\mathrm{th}}$ BMR in the semi-synthetic sunspot record is the time $t_n$ when the BMR appears.
+" The contribution to the source term 5(4,9,1) from the n BMR is taken to be ó(t—t,)(B*(4,9)-B9)), where 6 is the Dirac delta function."," The contribution to the source term $S(\lambda,\phi,t)$ from the $n^{\mathrm{th}}$ BMR is taken to be $\delta(t - t_n) \left(B^{+}(\lambda,\phi)-B^{-}(\lambda,\phi)\right)$, where $\delta$ is the Dirac delta function."
+ This corresponds to modeling the sunspot group emergence process as being instantaneous., This corresponds to modeling the sunspot group emergence process as being instantaneous.
+ The SFTM describes the evolution of the magnetic field on the Sun's surface., The SFTM describes the evolution of the magnetic field on the Sun's surface.
+ To obtain the heliospheric open flux we have to extrapolate the surface field outward., To obtain the heliospheric open flux we have to extrapolate the surface field outward.
+" Since the often used potential field source surface model does not well represent the Ulysses spacecraft data (?),, we use the better suited current sheet source surface (CSSS) extrapolation (??).."," Since the often used potential field source surface model does not well represent the Ulysses spacecraft data \citep{Schuessler06}, we use the better suited current sheet source surface (CSSS) extrapolation \citep{Zhao95a,Zhao95b}."
+" There are three parameters in the CSSS extrapolation, namely, one associated with the thickness of the current sheet, a, the radius of the cusp surface, Κος, and the radius of the source surface, R4."," There are three parameters in the CSSS extrapolation, namely, one associated with the thickness of the current sheet, $a$, the radius of the cusp surface, $R_{\rm{cs}}$, and the radius of the source surface, $R_{\rm{ss}}$."
+" The values for the three parameters are taken from the reference case in CJSS10, i.e, a=0.2Ro, Res=1.55Ro, Rss=10.ΌΚο."," The values for the three parameters are taken from the reference case in CJSS10, i.e., $a=0.2\rsun$, $R_{\rm{cs}}=1.55\rsun$, $R_{\rm{ss}}=10.0\rsun$."
+" The location of the cusp surface determines the amount of the open flux, which is calculated by the integration of unsigned magnetic field over the whole cusp surface."," The location of the cusp surface determines the amount of the open flux, which is calculated by the integration of unsigned magnetic field over the whole cusp surface."
+" In CJSS10, we used the SFTM with a source term based on the actually observed sunspot longitudes, latitudes, areas and cycle-averaged tilt angles to reconstruct the surface field and open flux for the period 1913-1986."," In CJSS10, we used the SFTM with a source term based on the actually observed sunspot longitudes, latitudes, areas and cycle-averaged tilt angles to reconstruct the surface field and open flux for the period 1913–1986."
+" The time evolution of the open flux derived from geomagnetic indices (??),, and the reversal times of polar fields (?) were well reproduced."," The time evolution of the open flux derived from geomagnetic indices \citep{Lockwood03, Lockwood09b}, and the reversal times of polar fields \citep{Makarov03} were well reproduced."
+ We therefore begin by comparing the results obtained using the current model based on sunspot numbers with those in CJSS10., We therefore begin by comparing the results obtained using the current model based on sunspot numbers with those in CJSS10.
+" To do so we have to consider the tendency for sunspots to occur in activity nests at specific longitudes, because this affects equatorial dipole moment and the open flux around activity maxima."," To do so we have to consider the tendency for sunspots to occur in activity nests at specific longitudes, because this affects equatorial dipole moment and the open flux around activity maxima."
+ We performed simulations with BMRs emerging at randomly distributed longitudes and simulations where all BMRs appear at longitudes 90° in the northern hemisphere and 270° in the southern hemisphere., We performed simulations with BMRs emerging at randomly distributed longitudes and simulations where all BMRs appear at longitudes $^{\circ}$ in the northern hemisphere and $^{\circ}$ in the southern hemisphere.
+" Since both the SFTM and the CSSS extrapolation are linear, we can combine the fields from the two kinds of simulations using where c is a measure of the strength of the activity nesting."," Since both the SFTM and the CSSS extrapolation are linear, we can combine the fields from the two kinds of simulations using where $\mathrm{c}$ is a measure of the strength of the activity nesting."
+ In Paper I we found that the value c=0.15 is appropriate for a determination of the open flux., In Paper I we found that the value $\mathrm{c}=0.15$ is appropriate for a determination of the open flux.
+ This is confirmed here., This is confirmed here.
+ Figure 2 shows the evolution of the open flux from CJSS10 compared with cases with c=0 (purely random) and c= 0.15., Figure \ref{fig:calibrate_c} shows the evolution of the open flux from CJSS10 compared with cases with $\mathrm{c}=0$ (purely random) and $\mathrm{c}=0.15$ .
+two spectra dis the scale at which structures initially form.,two spectra is the scale at which structures initially form.
+ The A.7 spectra has more energy αἱ smaller scales. leading to smaller characteristie structure size.," The $k^{-2}$ spectra has more energy at smaller scales, leading to smaller characteristic structure size."
+ After a few tens of Alfvénn limes these smaller-scale strictures merge and the svstem resembles the initial A.? spectra., After a few tens of Alfvénn times these smaller-scale structures merge and the system resembles the initial $k^{-3}$ spectra.
+" The n, and vy phases can be either eross-correlated or uncorrelated.", The $n_k$ and $\psi_k$ phases can be either cross-correlated or uncorrelated.
+ By cross-correlated we mean (hat (he phase anele for each Fourier component of the η ancl 6. fields are equal., By cross-correlated we mean that the phase angle for each Fourier component of the $n_k$ and $\psi_k$ fields are equal.
+ In general. Craddocketal.(1991). focused on the formation and longevity of current filaments in a turbulent KAW system.," In general, \citet{crad91} focused on the formation and longevity of current filaments in a turbulent KAW system."
+ To preserve small-scale structure in the current filaments. these numerical solutions set jj—0 and had j(~10 . with a resolution of 128x128. corresponding (o a μον Of 44.," To preserve small-scale structure in the current filaments, these numerical solutions set $\eta = 0$ and had $\mu \sim 10^{-3}$ , with a resolution of $128 \times 128$, corresponding to a $k_{max}$ of 44."
+ Laree-amplitude density structures that would have arisen were dampecl {ο preserve munerical stability up to an advecüve instability Gime of a few hundred. Alfvéenn times. for the parameter values therein.," Large-amplitude density structures that would have arisen were damped to preserve numerical stability up to an advective instability time of a few hundred Alfvénn times, for the parameter values therein."
+ The numerical solutions presented here explore a range of parameter values for 7j and ye., The numerical solutions presented here explore a range of parameter values for $\eta$ and $\mu$.
+ They make use of hvper-diffusivity and hyper-resistivily of appropriate strengths to preserve structures in n. D and J.," They make use of hyper-diffusivity and hyper-resistivity of appropriate strengths to preserve structures in $n$, $B$ and $J$."
+ An advective instability is excited alter ~107 Allvénn times if resistive damping is negligible., An advective instability is excited after $\sim 10^2$ Alfvénn times if resistive damping is negligible.
+ The 1=0 solutionsnot presented here due to their poor resolution of small-scale structuressee current filaments arise. but thev can be poorly resolved at this grid spacing.," The $\eta = 0$ solutions–not presented here due to their poor resolution of small-scale structures–see large-amplitude current filaments arise, but they can be poorly resolved at this grid spacing."
+ With no resistivity. the finite number of Fourier modes cannot resolve arbitrarily small structures without Gibbs phenomena resulting aud distorting (he current [ield.," With no resistivity, the finite number of Fourier modes cannot resolve arbitrarily small structures without Gibbs phenomena resulting and distorting the current field."
+ We have found through experience that small hyper-resistivity and small hyper-diffusivity preserve large-amplitude density. structures and (their spatial correlation with the magnetic and current structures. while preventing the distortion resulüing rom poorly-resolved current sheets and filaments.," We have found through experience that small hyper-resistivity and small hyper-diffusivity preserve large-amplitude density structures and their spatial correlation with the magnetic and current structures, while preventing the distortion resulting from poorly-resolved current sheets and filaments."
+ They allow the numerical solutions to run for arbitrarily long times. €xl the effects of structure mergers become apparent.," They allow the numerical solutions to run for arbitrarily long times, and the effects of structure mergers become apparent."
+ These occur on a longer timescale than the slowest eddy turnover times., These occur on a longer timescale than the slowest eddy turnover times.
+ The results presented here will consider (wo regimes of parameter values. the azzji and ηδνqi regimes.," The results presented here will consider two regimes of parameter values, the $\eta \approx \mu$ and $\eta \gg \mu$ regimes."
+The effect of cross-correlated and uncorrelated initial conditions will be addressed presently.,The effect of cross-correlated and uncorrelated initial conditions will be addressed presently.
+"paranueters,",parameters.
+ While there were mauy parameter sets in the | success rauese in the supersonic case. we choose r-—1l.dl aud 51.08 and box size=61 pixels. as it was the best match we found with the subsonic cases (67%))," While there were many parameter sets in the + success range in the supersonic case, we choose $\gamma_T$ =1.1 and $\beta_T$ =1.58 and box size=64 pixels, as it was the best match we found with the subsonic cases )"
+The study of protoplauetary disks is currently uudergonmg an exciting∙∙ stage partlv propelled |by the discovery. of| extrasolar plauctary svsteiis following|. the detection: of m51 Peg D by Mayor Queloz| (1995).,The study of protoplanetary disks is currently undergoing an exciting stage partly propelled by the discovery of extrasolar planetary systems following the detection of 51 Peg B by Mayor Queloz (1995).
+", The possible: discovery| of telluric planets iu the near future. in addition to the jovian-like planets already. detected. will pose interesting questionCSTIOUS∙ ο nthee formation.Qrluie() Qof. CXTTOasSOIarextrasolar planctaryaletary svstenis."," The possible discovery of telluric planets in the near future, in addition to the jovian-like planets already detected, will pose interesting questions on the formation of extrasolar planetary systems."
+ IKuowledge of the properties of protoplauctary ∙ ≼∐↴∖↴↘↽↴∖↴⋜⋯≼↧∐∪↖↖⇁↑∐↸∖⋅↖⇁↸∖↖↽∪↕↖↽↸∖↑∪≺∐∖↴⋝⊔↴∖↴≺∐↴∖↴↘↽↴∖↴⋜∐⋅⋯∐∐↧⊔⋜⊔∐≓∙∙ ∙sequence. (MS). stars is one of the fools required. for modeling such a process., Knowledge of the properties of protoplanetary disks and how they evolve to debris disks around main-sequence (MS) stars is one of the tools required for modelling such a process.
+ Observed: spectral energy distributions (SEDs) of pre-main sequence (PALS) stars are widely used to study the properties of protoplauctary disks aud to classify Ἑ Tauri and Herbig AcBe stars (TAcBe hereafter) iuto απ evolutionary scheme (Adams et al., Observed spectral energy distributions (SEDs) of pre-main sequence (PMS) stars are widely used to study the properties of protoplanetary disks and to classify T Tauri and Herbig AeBe stars (HAeBe hereafter) into an evolutionary scheme (Adams et al.
+ 1987: Hilleubraud et al., 1987; Hillenbrand et al.
+ 1992)., 1992).
+ Although the interpretation of a given SED is. to sole extent. model dependent. the theoretical modeling of SEDs coustitnes ani invaluable tool for understaudiug the structure iand properties: of protoplauetaryiarvo disks. οςe.g. Chiang Coldreich. (L997. 1999). D'Alessio et al. (," Although the interpretation of a given SED is, to some extent, model dependent, the theoretical modelling of SEDs constitutes an invaluable tool for understanding the structure and properties of protoplanetary disks, e.g. Chiang Goldreich (1997, 1999), D'Alessio et al. ("
+1998 1999. 2001). Dullemond et al. (,"1998, 1999, 2001), Dullemond et al. ("
+2001).,2001).
+ The EXPORT consortium (Eiroa et al., The EXPORT consortium (Eiroa et al.
+ 2000) observed a large sample of PAISd and Veea-type stars duringH. the 11998d Iuternational: ⊺↕⋯↸∖↕⋟↥⋅∪∶↴∙↥∡∐⊔⊔↸∖∪↕↑↕∐∖≼⊲⋜⊔⋜∐⋅⋅↖⇁↕↴∖↴↕⋜⋯≼↧↴∖↴∎≼⊽∏, 2000) observed a large sample of PMS and Vega-type stars during the 1998 International Time Programme of the Canary Islands' Observatories.
+⋝↴∖↴↸∖↥⋅↖↽⋜↧↑∪↥⋅↕↸∖↴∖↴⋅ One of t1ο drvius goals of this effort was the study o£ the evolution and xoperties of protoplanetary disks by analysing the SEDs of the wouug stars. taking advantage. of the fact that the EXPORT optical and pearIR photometry were obtained simultaucoushv.," One of the driving goals of this effort was the study of the evolution and properties of protoplanetary disks by analysing the SEDs of the young stars, taking advantage of the fact that the EXPORT optical and near-IR photometry were obtained simultaneously."
+ This ohsemecitioua iyproach is appropriate since T Tami aud [TAeBe stars vary markedly in these spectral regimes aud a sienificaut. part of the total huninosity of the object is radiated by the PMS stellar photosphere at these wavelengths., This observational approach is appropriate since T Tauri and HAeBe stars vary markedly in these spectral regimes and a significant part of the total luminosity of the object is radiated by the PMS stellar photosphere at these wavelengths.
+ Among the stars in the EXPORT sample withE measuredj parallaxes.30 the ILAcBestare stars IID 34982 and TID L11569 are the only ones whose positions fall below the zero-age main sequence ina logLTig UR diagram or equivalent (c.g. van deu Aucker et al.," Among the stars in the EXPORT sample with measured parallaxes, the HAeBe stars HD 34282 and HD 141569 are the only ones whose positions fall below the zero-age main sequence in a $\log L - \log T_{\rm eff}$ HR diagram or equivalent (e.g. van den Ancker et al."
+"1998,"," 1998,"
+Globular clusters (GC's) are known to be a rich source of ow mass X-rav binaries (LAINDs).,Globular clusters (GCs) are known to be a rich source of low mass X-ray binaries (LMXBs).
+ In carly studies of the Alilky Way's X-ray population. it was realised. that many more LMXDs were located in GCs than would be expected owed. on their masses alone.," In early studies of the Milky Way's X-ray population, it was realised that many more LMXBs were located in GCs than would be expected based on their masses alone."
+ Of the ~150 bright LAINBs currently known in the Alilky Way. 14 reside in 12 of its GCs (227).," Of the $\sim$ 150 bright LMXBs currently known in the Milky Way, 14 reside in 12 of its GCs \citep{Liu01,Heinke09}."
+ Since the Galactic GC's are thought to containess than 0.1% of the stars in the Galaxy. this suggests hat the formation of LAINBs is more than two orders. of magnitude more cllicient in GC's than in the field of the Galaxy.," Since the Galactic GCs are thought to containless than $\%$ of the stars in the Galaxy, this suggests that the formation of LMXBs is more than two orders of magnitude more efficient in GCs than in the field of the Galaxy."
+ Ht has long been proposed that this extra cllicicney is due to dvnamical formation of these binaries in the dense cluster cores (?2)..," It has long been proposed that this extra efficiency is due to dynamical formation of these binaries in the dense cluster cores \citep{Katz75,Clark75}."
+ Phe proposed mechanisms for dvnamical formation of LMND svstems include: a donor star capturing the neutron star through tidal capture (e.g.2): exchange interactions between a neutron star and a primordial binary system (22): or direct. collisions of a neutron star with the envelope of a giant star. (2)...," The proposed mechanisms for dynamical formation of LMXB systems include: a donor star capturing the neutron star through tidal capture \citep[e.g.][]{Fabian75}; exchange interactions between a neutron star and a primordial binary system \citep{Hills76,Hut83}; or direct collisions of a neutron star with the envelope of a giant star \citep{Verbunt87}."
+ These interactions are more likely to occur in the cores of GC's. than in field star populations. due to their high stellar densities.," These interactions are more likely to occur in the cores of GCs, than in field star populations, due to their high stellar densities."
+ The properties of LAINB hosting clusters can be used to investigate the formation and evolution of these svstenmis., The properties of LMXB hosting clusters can be used to investigate the formation and evolution of these systems.
+ For example. if cvnamical formation is the primary method of forming LAINBs in GC's. then it is expected that there should be a direct relationship between the stellar collision rate in a cluster and the presence of an LAINB.," For example, if dynamical formation is the primary method of forming LMXBs in GCs, then it is expected that there should be a direct relationship between the stellar collision rate in a cluster and the presence of an LMXB."
+ Previous work has shown that the LAINBs in the Alilky Way's GCs are consistent with clvnamical formation scenarios (?).., Previous work has shown that the LMXBs in the Milky Way's GCs are consistent with dynamical formation scenarios \citep{Verbunt87}.
+ Llowever. studsing such relationships in the Galactic CC svstem is limited by the low number of LAINBs in its GCs.," However, studying such relationships in the Galactic GC system is limited by the low number of LMXBs in its GCs."
+ In order to study more significant sample sizes in the Milky Mes GCs. it is necessary to combine dillerent. types of compact object binary systems (c.g. 2)..," In order to study more significant sample sizes in the Milky Way's GCs, it is necessary to combine different types of compact object binary systems \citep[e.g.][]{Pooley03}. ."
+ In the era of and it has, In the era of and it has
+The periodogram for sshows (wo clear solutions (hat are nearly equivalently οσους fits: 9.1 hours ancl. secondarily. 7.3 hours (Figure 1)).,"The periodogram for shows two clear solutions that are nearly equivalently good fits: 9.1 hours and, secondarily, 7.3 hours (Figure \ref{period}) )."
+ Both solutions have amplitudes of 1.09 magnitudes., Both solutions have amplitudes of 1.09 magnitudes.
+ The secondary peak is non-resonant with the best fit. ie. not an obvious harmonic of the best-fit period. ad is also equite siegnilicant compared to the randomized data.," The secondary peak is non-resonant with the best fit, i.e., not an obvious harmonic of the best-fit period, and is also quite significant compared to the randomized data."
+ We therefore searched further for a best-lit solution that consisted of (wo independent sinusoids with independent phases. amplitudes. ancl periods. though the periods were restricted to a small range around each of best fits derived above.," We therefore searched further for a best-fit solution that consisted of two independent sinusoids with independent phases, amplitudes, and periods, though the periods were restricted to a small range around each of best fits derived above."
+ Formally. the 47 improves significantly through allowing a (wo-sine fit. but the data may not warrant attaching too much importance to a multiple rotation pole interpretation.," Formally, the $\chi^2$ improves significantly through allowing a two-sine fit, but the data may not warrant attaching too much importance to a multiple rotation pole interpretation."
+ In the following analvsis. we use the single-sinusoid better-fitting 9.1 hour period.," In the following analysis, we use the single-sinusoid better-fitting 9.1 hour period."
+ None of our conclusions depend upon the choice between the (wo best-fit: periods. nor particularly on (he choice of single- over couble-sinusoid fit.," None of our conclusions depend upon the choice between the two best-fit single-sinusoid periods, nor particularly on the choice of single- over double-sinusoid fit."
+ The best-fit sinusoid solution to the photometry of hhas a period of 2.8 hours and amplitude of 0.42 magnitudes., The best-fit sinusoid solution to the photometry of has a period of 2.8 hours and amplitude of 0.42 magnitudes.
+ However. the significance of (his solution is only when compared to LOO randomized trials of the ddata.," However, the significance of this solution is only when compared to 100 randomized trials of the data."
+ We therefore conclude (hat we failed to detect. significant. periodic variability [orΕμμ., We therefore conclude that we failed to detect significant periodic variability for.
+.. To place an upper limit on the amplitude of an undetected periodic variation. we augmented the ddata with svnthetc lighteurves wilh various amplitudes ancl periods of 4 hours (a typical KBO photometric variation. period) and carried oul Che best-fit solution search. described above.," To place an upper limit on the amplitude of an undetected periodic variation, we augmented the data with synthetic lightcurves with various amplitudes and periods of 4 hours (a typical KBO photometric variation period) and carried out the best-fit solution search described above."
+ We successfully recovered all synthetic lighteurves with amplitudes larger than around154., We successfully recovered all synthetic lightcurves with amplitudes larger than around.
+" We can therefore place an upper limit on possible lighteurve amplitudes forBIlo,. requiring that anv such variation must have an amplitude less than 0.15 magnitudes to be undetected by us."," We can therefore place an upper limit on possible lightcurve amplitudes for, requiring that any such variation must have an amplitude less than 0.15 magnitudes to be undetected by us."
+ The albedo variations and/or asphericity of this body must be less (han15'4., The albedo variations and/or asphericity of this body must be less than.
+. The best-fit solution for exeeives a period of 7.5 hours., The best-fit solution for gives a period of 7.5 hours.
+ There is a peak of nearly equal signifieanceex al 3.79 hours. almost exaclly half the best solution.," There is a peak of nearly equal significance at 3.79 hours, almost exactly half the best solution."
+ The amplitude of the lighteurve is identical (0.07 mag) for, The amplitude of the lightcurve is identical (0.07 mag) for
+the X-ray emission up to 10.5 keV. We do this by presenting deconvolved images m narrow bands obtained by the MECS instrument on board BeppoSAX.,the X-ray emission up to 10.5 keV. We do this by presenting deconvolved images in narrow bands obtained by the MECS instrument on board BeppoSAX.
+ Before discussing our analysts. a short characterization of Cas A. The remnant is at a distance of 3.1!07 kpe (Reed et al. 1995))," Before discussing our analysis, a short characterization of Cas A. The remnant is at a distance of $3.4^{+0.3}_{-0.1}$ kpc (Reed et al. \cite{Reed}) )"
+ and is between 340 and 318 years old., and is between 340 and 318 years old.
+ The younger age Is based on the possible explosion date of 1680. when Flamsteed observed an unidentified Gth magnitude star near the current position of Cas A (Ashworth 1980)).," The younger age is based on the possible explosion date of 1680, when Flamsteed observed an unidentified 6th magnitude star near the current position of Cas A (Ashworth \cite{Ashworth}) )."
+ The older age is based on the expansion analysis of optical filaments (van den Bergh Kamper 1983))., The older age is based on the expansion analysis of optical filaments (van den Bergh Kamper \cite{vdBergh83}) ).
+ The current wisdom ts that the progenitor was probably an early type Wolf-Rayet star (Fesen et al. 1987)), The current wisdom is that the progenitor was probably an early type Wolf-Rayet star (Fesen et al. \cite{Fesen87}) )
+ with an initial mass of about ((Gareiá-Segura et al. 1996)), with an initial mass of about (Garciá–Segura et al. \cite{GLM96}) )
+ and a final mass less than (Vink et al. 1996)., and a final mass less than (Vink et al. \cite{VKB96}) ).
+ Recent expansion. studies of high resolution X-ray images indicate. that the current overall expansion velocity is ~ 3200 km/s for the bright shell and ~ 5200 km/s for the fainter emission associated with the blast wave (Vink et al. 1998..," Recent expansion studies of high resolution X-ray images indicate that the current overall expansion velocity is $\sim$ 3200 km/s for the bright shell and $\sim$ 5200 km/s for the fainter emission associated with the blast wave (Vink et al. \cite{Vink98},"
+ Koralesky et al. 1998))., Koralesky et al. \cite{Koralesky}) ).
+ This is considerably greater than the expansion of the bright radio ring which ts about 2000 km/s (Anderson Rudnick 1995))., This is considerably greater than the expansion of the bright radio ring which is about 2000 km/s (Anderson Rudnick \cite{AR95}) ).
+ The BeppoSAX satellite (Boella et al. 1997a)), The BeppoSAX satellite (Boella et al. \cite{Boella97a}) )
+ contains four X-ray detectors: the LECS (0.1-10 keV) and MECS (1.5-12 keV). composed of three units. the HPGSPC (4-120 keV) and the PDS (13-300 keV).," contains four X-ray detectors: the LECS (0.1–10 keV) and MECS (1.5--12 keV), composed of three units, the HPGSPC (4–120 keV) and the PDS (13–300 keV)."
+ The overall spectrum of Cas A as observed by BeppoSAX is shown in the paper by Favata et al. (1997))., The overall spectrum of Cas A as observed by BeppoSAX is shown in the paper by Favata et al. \cite{Favata}) ).
+ Only the LECS and MECS have imaging capabilities., Only the LECS and MECS have imaging capabilities.
+ Here we have opted for only using the MECS (Boella et al. 1997b)), Here we have opted for only using the MECS (Boella et al. \cite{Boella97b}) )
+ instrument. since the supporting structure of the LECS blocks part of the field of view. whereas the MECS has an unobstructed view of Cas A. which simplifies the image analysis considerably.," instrument, since the supporting structure of the LECS blocks part of the field of view, whereas the MECS has an unobstructed view of Cas A, which simplifies the image analysis considerably."
+ Furthermore. the MECS has a higher sensitivity above 5 keV than the LECS.," Furthermore, the MECS has a higher sensitivity above 5 keV than the LECS."
+ BeppoSAX observed Cas A five times: four observations during the performance verification phase in August and September 1996 and one observation on November 26. 1997.," BeppoSAX observed Cas A five times; four observations during the performance verification phase in August and September 1996 and one observation on November 26, 1997."
+ The total observation time is different for each set of instruments. but data presented here comprise 161 ks of MECS observation time.," The total observation time is different for each set of instruments, but data presented here comprise 161 ks of MECS observation time."
+" The data analyzed here are based on the MECS2 and MECS3 units (labeled together as MECS23) for uniformity: the MECSI was no longer functioning at the time of the last observation,", The data analyzed here are based on the MECS2 and MECS3 units (labeled together as MECS23) for uniformity: the MECS1 was no longer functioning at the time of the last observation.
+ Comparing the various data sets we found that there were rather large inaccuracies in the aspect solutions., Comparing the various data sets we found that there were rather large inaccuracies in the aspect solutions.
+ This was most distinct for the last observation (~1.1/)) made after a gyroscope failure. which may have caused the relatively large error in the aspect solution.," This was most distinct for the last observation $\sim 1.4$ ) made after a gyroscope failure, which may have caused the relatively large error in the aspect solution."
+ Using a correlation analysis similar to that employed by Vink et al. (1998)).," Using a correlation analysis similar to that employed by Vink et al. \cite{Vink98}) ),"
+ we shifted the observations in order to match in position., we shifted the observations in order to match in position.
+ The MECS point spread function (PSF) consists of two components. representing the gaussian shaped intrinsic," The MECS point spread function (PSF) consists of two components, representing the gaussian shaped intrinsic"
+ , 
+Lab. Los Alamos National Lab. and the University of Michigan. (http://www.rotse.net).,"Lab, Los Alamos National Lab, and the University of Michigan (http://www.rotse.net)."
+ We thank the Turkish National Observatory of TUBBITTAK for running the optical facilities., We thank the Turkish National Observatory of TÜBBİTTAK for running the optical facilities.
+ We acknowledge support from TUBBITTAK. The Scientific and Technological Research Council of Turkey. through project 106TO40.," We acknowledge support from TÜBBİTTAK, The Scientific and Technological Research Council of Turkey, through project 106T040."
+ We also acknowledge the RXTE/ASM team for the X-ray monitoring data., We also acknowledge the /ASM team for the X-ray monitoring data.
+ SWIFT/BAT transient monitor results were provided by the SWIFT/BAT team., /BAT transient monitor results were provided by the /BAT team.
+available in its full content as Supplementary Material in the electronic edition of this article.,available in its full content as Supplementary Material in the electronic edition of this article.
+ The Las Campanas plates had 103aO0 emulsion and were exposed through a GG 385 filter. so the magnitudes are 1 magnitudes.," The Las Campanas plates had 103aO emulsion and were exposed through a GG 385 filter, so the magnitudes are $B$ magnitudes."
+ A total of 508 plates was obtained. ancl they were all measured on a Culley iris photometer (Culley1956. 1961).," A total of 508 plates was obtained, and they were all measured on a Cuffey iris photometer \citep{c56,c61}."
+. The plates from. 1972 to 1980 were measured. by sunmimer students. and the data were used. but not published. in the study ofthe Fourier parameter ys. by Clementet (1992).," The plates from 1972 to 1980 were measured by summer students, and the data were used, but not published, in the study of the Fourier parameter $\varphi_{31}$ by \cite{cj92}."
+. The 1981 to 1991 plates were measured later by €. Clement., The 1981 to 1991 plates were measured later by C. Clement.
+ Since the plates were measured at dillerent times bv dillerent. people. there might be svstematic magnitude shifts from vear to vear. particularly for stars that were subject to crowding.," Since the plates were measured at different times by different people, there might be systematic magnitude shifts from year to year, particularly for stars that were subject to crowding."
+ The 1981.1991 magnitudes should be on a consistent system., The 1981–1991 magnitudes should be on a consistent system.
+ However. even then. for. erowcded stars. Lb was not always possible to be consistent because of cilferences in the seeing conditions.," However, even then, for crowded stars, it was not always possible to be consistent because of differences in the seeing conditions."
+ Altogether 73 Rt Όντας stars were measured., Altogether 73 RR Lyrae stars were measured.
+ The photographic observations at. Budapest with the GO-cm telescope were exposed on Guilleminot Superfulgur or Agla Astro Spezial plates., The photographic observations at Budapest with the 60-cm telescope were exposed on Guilleminot Superfulgur or Agfa Astro Spezial plates.
+ Since the sensitivity of the emulsions is similar. the measurements or estimates should be on à consistent svstem.," Since the sensitivity of the emulsions is similar, the measurements or estimates should be on a consistent system."
+ On 292 plates the brightness of 39 lut Lyrae stars was measured., On 292 plates the brightness of 39 RR Lyrae stars was measured.
+ The magnitudes of a further 35 variables that were close to the erowcded central region or had close companion(s) were estimated by M. Lovas., The magnitudes of a further 35 variables that were close to the crowded central region or had close companion(s) were estimated by M. Lovas.
+ The observations with the l-m telescope have been made on 103aO0 plates., The observations with the 1-m telescope have been made on 103aO plates.
+ In the vears 1976... 1988. and 190293 altogether 94. 103aO plates were exposed. without filters.," In the years 1976, 1988 and 1992–93 altogether 94 103aO plates were exposed without filters."
+ In the other vears 104. photographs were obtained on 103aO plates through a 2mm GG1S8 filter., In the other years 104 photographs were obtained on 103aO plates through a 2mm GG13 filter.
+ In. all. 93 variables were measured on these plates.," In all, 93 variables were measured on these plates."
+ We found no serious systematic cüllerence between the filtered and unfiltered blue magnitudes that exceeded the measurement errors., We found no serious systematic difference between the filtered and unfiltered blue magnitudes that exceeded the measurement errors.
+ The measuring procedure of the photographic plates obtained with the GO-cm ancl 1-m telescopes was the Following., The measuring procedure of the photographic plates obtained with the 60-cm and 1-m telescopes was the following.
+" The plates were digitized on a Umax PowerLook 3000 [Latbed-transpareney scanner with 3.13 and 1.29 pixel-per-aresce resolution on the 1-m and 60-em. photographs. respectively,"," The plates were digitized on a Umax PowerLook 3000 flatbed-transparency scanner with 3.13 and 1.29 pixel-per-arcsec resolution on the 1-m and 60-cm photographs, respectively."
+ Digital. aperture photometry was applied. to the images using standard packages., Digital aperture photometry was applied to the images using standard packages.
+ In this wav. the photographic densities of the variables and comparison stars in M5 have been determined.," In this way, the photographic densities of the variables and comparison stars in M5 have been determined."
+ The brightness of variables on the Asiago plates that were obtained in bad secing conditions have been estimated bv M. Lovas., The brightness of variables on the Asiago plates that were obtained in bad seeing conditions have been estimated by M. Lovas.
+ These. plates had. HaO emulsion. and were exposed without filter., These plates had IIaO emulsion and were exposed without filter.
+ In. all 67 variables were estimated on 35 plates., In all 67 variables were estimated on 35 plates.
+ In the course of measuring anc estimating the brightness of the variables. the comparison stars were chosen from Arp's photometry (Arp1955.1962).," In the course of measuring and estimating the brightness of the variables, the comparison stars were chosen from Arp's photometry \citep{ar55,ar62}."
+. Llis photographic (myg4) comparison sequence was applied to the Budapest and Asiago plates., His photographic $m_{pg}$ ) comparison sequence was applied to the Budapest and Asiago plates.
+ The Las Campanas ancl Piszkésstetó observations are on the D magnitude scale., The Las Campanas and Piszkésstető observations are on the $B$ magnitude scale.
+ No systematic cilference exceeding the 0.02-mag observational error. has been found between the filtered and. untilterecl Piszkésstetó blue observations., No systematic difference exceeding the 0.02-mag observational error has been found between the filtered and unfiltered Piszkésstető blue observations.
+ These observations supplemented: with the published ones provide an almost. continuous coverage for the second half of the last. century., These observations supplemented with the published ones provide an almost continuous coverage for the second half of the last century.
+ We consider only those II Lyrae stars of M5 that have a long enough historical record to follow their period changes Le. variables with numbering below 100., We consider only those RR Lyrae stars of M5 that have a long enough historical record to follow their period changes i.e. variables with numbering below 100.
+ V53. νου V93 and VOL are omitted because their photometry is seriously allectecd by the contamination of close companions.," V53, V86, V93 and V94 are omitted because their photometry is seriously affected by the contamination of close companions."
+ In all. we investigate the period changes of 65 Rab and 21 RRe stars of M5.," In all, we investigate the period changes of 65 RRab and 21 RRc stars of M5."
+ We used. the original photometric data. (reviewed. in Sections 2.1 and 2.2) to derive the O Cs. ‘Table 2 summarizes the available photometry for each RR Lyrae star of M5. the sources are denoted by 7|," We used the original photometric data (reviewed in Sections 2.1 and 2.2) to derive the $O-C$ s. Table \ref{data} summarizes the available photometry for each RR Lyrae star of M5, the sources are denoted by `+'."
+ The procedure applied to the determination of the OC diagrams and to explore the period changes is basically the same as used. by Juresiketal.(2001). in the study. of g Centauri's variables., The procedure applied to the determination of the $O-C$ diagrams and to explore the period changes is basically the same as used by \cite{ocen} in the study of $\omega$ Centauri's variables.
+ As the first. step. the zero-points of the magnitudes of the different observations of cach RR Lyrac star were homogenized.," As the first step, the zero-points of the magnitudes of the different observations of each RR Lyrae star were homogenized."
+ The mean light curves used for the magnitude transformations were the Las Campanas data from. 1991 and for the 14 stars (V5. νο V54. V65. V6T. VSs5. VST. V90. VOL. V95. V96. WOT. V9s. V99) that were not or only partly measured on the Las Campanas plates. they were constructed from the 1976.1993 Piszkésstetó data.," The mean light curves used for the magnitude transformations were the Las Campanas data from 1981--1991 and for the 14 stars (V5, V36, V54, V65, V67, V85, V87, V90, V91, V95, V96, V97, V98, V99) that were not or only partly measured on the Las Campanas plates, they were constructed from the 1976–1993 Piszkésstető data."
+ The light curves of the RRab and RRe stars have been decomposed by fifth and third order Fourier series. respectively. to fix the mean light curves.," The light curves of the RRab and RRc stars have been decomposed by fifth and third order Fourier series, respectively, to fix the mean light curves."
+ For the sake of uniformity. we used these relatively low-order fits for the Rab stars. because higher-order solutions gave wavy curves for some of the stars.," For the sake of uniformity, we used these relatively low-order fits for the RRab stars, because higher-order solutions gave wavy curves for some of the stars."
+ We checked. however. to ensure that the zero-points and. phase shifts would remain unchanged if higher order fits had been used for the variables with good-quality light curves.," We checked, however, to ensure that the zero-points and phase shifts would remain unchanged if higher order fits had been used for the variables with good-quality light curves."
+ Both vertical and horizontal shifts were determined for all the dillerent. photometric data sets of the variables in order to match their normal curves best., Both vertical and horizontal shifts were determined for all the different photometric data sets of the variables in order to match their normal curves best.
+ “Phe first. part of the Las Campanas observations. from. 1972 to 1980. were," The first part of the Las Campanas observations, from 1972 to 1980 were"
+Faraday rotated by an equal amount (defined by Eq. 155).,Faraday rotated by an equal amount (defined by Eq. \ref{eq:FR}) ).
+ Since there is no differential Faraday rotation. depolarization does not occur and the polarized intensity of the synchrotron emission is unchanged (see Fig. 4.," Since there is no differential Faraday rotation, depolarization does not occur and the polarized intensity of the synchrotron emission is unchanged (see Fig. \ref{fig:PImodels}."
+.a and Table 1)., .a and Table \ref{tab:maps}) ).
+ Morphology of the polarization angles over the map is determined only by the spatial ARAL variations of the plasma., Morphology of the polarization angles over the map is determined only by the spatial $RM$ variations of the plasma.
+ In order to test the above results. we correlate a polarized emission image with a total intensity. and polarization angles with the image of free-free emission.," In order to test the above results, we correlate a polarized emission image with a total intensity, and polarization angles with the image of free-free emission."
+ The cross correlation is C.=| in both cases., The cross correlation is $C=1$ in both cases.
+ Since the polarized emission s completely correlated with the total intensity. and the free-free emission. which traces thermal plasma. with polarization angles. we conclude that our algorithm is preforming as expected.," Since the polarized emission s completely correlated with the total intensity, and the free-free emission, which traces thermal plasma, with polarization angles, we conclude that our algorithm is preforming as expected."
+ The spatial structures of polarized emission are unchanged. while the structures of polarization angle follow thermal plasma.," The spatial structures of polarized emission are unchanged, while the structures of polarization angle follow thermal plasma."
+ In B and C models there are regions where both CR and thermal electrons are mixed., In B and C models there are regions where both CR and thermal electrons are mixed.
+ In those regions the polarization angle of the synchrotron radiation is then differentially Faraday rotated along the LOS., In those regions the polarization angle of the synchrotron radiation is then differentially Faraday rotated along the LOS.
+ As a result. the polarized synchrotron radiation is quenched and the level of the polarized emission is weaker than intrinsic polarization emission (see Table |).," As a result, the polarized synchrotron radiation is quenched and the level of the polarized emission is weaker than intrinsic polarization emission (see Table \ref{tab:maps}) )."
+ In addition to depolarization produced by the mixed regions. in model C depolarization is also produced by variations of the Galactic magnetic field along the LOS.," In addition to depolarization produced by the mixed regions, in model C depolarization is also produced by variations of the Galactic magnetic field along the LOS."
+ As for model A. we calculate correlations between the total and polarized images.," As for model A, we calculate correlations between the total and polarized images."
+ In both models «B. C) we find much weaker correlation (6ρε= 3.0.1) than in model A. The weaker correlation results from differential Faraday rotation along the LOS. which changes the morphology of polarized emission both in intensity and angle.," In both models (B C) we find much weaker correlation $C_{B,C}=0.3,0.1$ ) than in model A. The weaker correlation results from differential Faraday rotation along the LOS, which changes the morphology of polarized emission both in intensity and angle."
+ To illustrate this we correlate the polarized synchrotron emission with the free-free emission., To illustrate this we correlate the polarized synchrotron emission with the free-free emission.
+ We find anti-correlation in both models (Cpe:=0.7. 0.5).," We find anti-correlation in both models $C_{B,C}=-0.7,-0.5$ )."
+ The intensity of the polarized emission is more quenched in the regions that show stronger free-free emission (regions of denser plasma)., The intensity of the polarized emission is more quenched in the regions that show stronger free-free emission (regions of denser plasma).
+ Moreover. model C shows weaker (anti)correlations than model B. due to Faraday structures produced by magnetic field.," Moreover, model C shows weaker (anti)correlations than model B, due to Faraday structures produced by magnetic field."
+ The last aspect of the algorithm. which needs to be verified. is the frequency behavior of the simulated Galactic emission.," The last aspect of the algorithm, which needs to be verified, is the frequency behavior of the simulated Galactic emission."
+ It is expected that the synchrotron emission shows roughly a power law along the frequency (see Eq. 3)), It is expected that the synchrotron emission shows roughly a power law along the frequency (see Eq. \ref{eq:syn}) ).
+ By examining Fig. 6..," By examining Fig. \ref{fig:los},"
+ one can see that simulated synchrotron emission indeed shows a power law behavior in total intensity., one can see that simulated synchrotron emission indeed shows a power law behavior in total intensity.
+ However. in polarized intensity only model A shows a power law behavior.," However, in polarized intensity only model A shows a power law behavior."
+ All the other models show characteristic fluctuations caused by differential Faraday rotation (?).., All the other models show characteristic fluctuations caused by differential Faraday rotation \citep{burn66}.
+ Note that fluctuations show richer structures. as in model C. if the differential Faraday rotation is more prominent.," Note that fluctuations show richer structures, as in model C, if the differential Faraday rotation is more prominent."
+ The calculated brightness temperature spectral index. taiji. variations of the simulated Galactic synchrotron emission are given in the Fig. 3.," The calculated brightness temperature spectral index, $\beta_{syn}$, variations of the simulated Galactic synchrotron emission are given in the Fig. \ref{fig:synff}."
+c., .c.
+" The variations of ;,,, are estimated by fitting the data with a linear functions along the frequency in logarithmic scale.", The variations of $\beta_{syn}$ are estimated by fitting the data with a linear functions along the frequency in logarithmic scale.
+" The mean value of 2,4; is -2.50. which is in a good agreement with the expected theoretical value ο—(p|3)/2=2.55."," The $mean$ value of $\beta_{syn}$ is -2.50, which is in a good agreement with the expected theoretical value $\beta_{p=2}=-(p+3)/2=-2.55$ ."
+ A slight difference between the two is caused by the 3D spatial variations of the spectral index p., A slight difference between the two is caused by the 3D spatial variations of the spectral index $p$.
+ Based on discussion and tests of the algorithm presented in this subsection. we can conclude that our algorithm is simulating Galactic emission as expected.," Based on discussion and tests of the algorithm presented in this subsection, we can conclude that our algorithm is simulating Galactic emission as expected."
+ In the following subsection. we will compare the simulated data cubes with the observations.," In the following subsection, we will compare the simulated data cubes with the observations."
+ Input parameters of the simulation (properties of the CR electrons. hermal electrons. and Galactic magnetic field. see subsections ??.. ?? 22)) are chosen in such way that the global properties (boundary conditions) of the simulated maps. e.g. morphology and requency behavior of the Galactic emission. are in agreement with he observations overviewed in Sec.," Input parameters of the simulation (properties of the CR electrons, thermal electrons, and Galactic magnetic field, see subsections \ref{sec:CRe}, \ref{sec:MF} \ref{sec:Te}) ) are chosen in such way that the global properties (boundary conditions) of the simulated maps, e.g. morphology and frequency behavior of the Galactic emission, are in agreement with the observations overviewed in Sec."
+ 3: We also quantitatively compare the simulated emission in the Fan region tour model D) with the observations of this region at 150MlIz (2x:, 3: We also quantitatively compare the simulated emission in the Fan region (our model D) with the observations of this region at $150~{\rm MHz}$ \citep{bernardi09a}: :
+The recent detection of Cyl] in the carbon-rich AGB star and the cold. dense interstellar cloud by MeCarthyetal.(2006) with abundance ratios relative to Call of 0.01.0.1 indicates (hat anions may play. a more significant role in interstellar plivsies and chemistry than heretofore believed.,"The recent detection of $_6$ $^-$ in the carbon-rich AGB star and the cold, dense interstellar cloud by \citet{mcc06} with abundance ratios relative to $_6$ H of 0.01–0.1 indicates that anions may play a more significant role in interstellar physics and chemistry than heretofore believed."
+ The possibility that a relatively large Iraction of molecular material in interstellar clouds might be in the form of anions was first suggested bv Herbst(1981). who pointed out that carbon. chain molecules and other radicals have large electron. affinities. leading to high radiative attachment rates such as those measured by Woodinetal.(1980). with the attendant possibility of anion/neutral fractions on the order of a few percent.," The possibility that a relatively large fraction of molecular material in interstellar clouds might be in the form of anions was first suggested by \citet{her81} who pointed out that carbon chain molecules and other radicals have large electron affinities, leading to high radiative attachment rates such as those measured by \cite{wood80} with the attendant possibility of anion/neutral fractions on the order of a few percent."
+ More recently. it has been recognized that PAIL anions could soak up a significant fraction of the [ree electrons in interstellar clouds and alter (he charge balance to a significant degree (Lepp as well as providing significant. heating through. photocdetachment of electrons in diffuse clouds (Lepp&Dalgarno1985b).," More recently, it has been recognized that PAH anions could soak up a significant fraction of the free electrons in interstellar clouds and alter the charge balance to a significant degree \citep{lep88a} as well as providing significant heating through photodetachment of electrons in diffuse clouds \citep{lep88b}."
+. Petrie(1996) has investigated the svnthesis of by the dissociative attachment of MeCN and MgNC and other mechanisms. while Petrie&Herbst(1997). showed that C4N σοι] be detectable in interstellar clouds.," \citet{pet96} has investigated the synthesis of $^-$ by the dissociative attachment of MgCN and MgNC and other mechanisms, while \citet{pet97} showed that $_3$ $^-$ could be detectable in interstellar clouds."
+ The study of large hydrocarbons anions received some attention following the observation bv Tulejetal.(1998). that absorption bands in several carbon-chain anions coincided with several of the diffuse interstellar bands., The study of large hydrocarbons anions received some attention following the observation by \citet{tul98} that absorption bands in several carbon-chain anions coincided with several of the diffuse interstellar bands.
+ Subsequently. (he formation of such species in diffuse clouds was studied by Ruffleetal.(1999) who showed. in particular. that Cl was unlikely to be a source for any DIDs.," Subsequently, the formation of such species in diffuse clouds was studied by \citet{ruf99} who showed, in particular, that $_7^-$ was unlikely to be a source for any DIBs."
+ Millaretal.(2000) considered the formation of hvdrocarbon anions containing more than six carbon atoms in the carbon-rich. circumstellar envelope of and showed that appreciable column densities could arise in the outer envelope., \citet{mil00} considered the formation of hydrocarbon anions containing more than six carbon atoms in the carbon-rich circumstellar envelope of and showed that appreciable column densities could arise in the outer envelope.
+ The basic route to the formation of anions is electron radiative attachment: which has been discussed in the context of bare carbon chains by Terzieva&Herbst(2000)., The basic route to the formation of anions is electron radiative attachment: which has been discussed in the context of bare carbon chains by \citet{ter00}.
+. The rate coellicients of relevant processes have not been experimentally determined at low temperatures and in (this work we have used plase-state theory assuming s-wave attachment with radiative stabilisation occurring by vibrational and electronic transitions llerbst 2000)., The rate coefficients of relevant processes have not been experimentally determined at low temperatures and in this work we have used phase-state theory assuming $s$ -wave attachment with radiative stabilisation occurring by vibrational and electronic transitions \citep{ter00}.
+. Table 6 contains some relevant attachment rate coefficients., Table \ref{tab:rates} contains some relevant attachment rate coefficients.
+ The ealeulations of rates [ον these and other processes will be discussed in a separate paper (Ilerbst. in preparation).," The calculations of rates for these and other processes will be discussed in a separate paper (Herbst, in preparation)."
+" For C,II radicals. we find that radiative attachment occurs al the collisional"," For $_{n}$ H radicals, we find that radiative attachment occurs at the collisional"
+The EVLA software is divided into two major areas: the control and monitoring of the hardware. or Monitor and Control (M&CC). and the software generally. accessed by the user. or User Software.,"The EVLA software is divided into two major areas: the control and monitoring of the hardware, or Monitor and Control C), and the software generally accessed by the user, or User Software."
+ Figure 1. shows a very high level view of the EVLA software system. along with how it fits in with the hardware. and with the information flow highlighted.," Figure \ref{fig:blockdiagram} shows a very high level view of the EVLA software system, along with how it fits in with the hardware, and with the information flow highlighted."
+ Two overriding principles guide the implementation of the software., Two overriding principles guide the implementation of the software.
+ First. the software must be loosely coupled. so that if a particular subsystem is not funcüoning. the rest ean continue {ο run.," First, the software must be loosely coupled, so that if a particular subsystem is not functioning, the rest can continue to run."
+ Second. the software must make it easier to use and operate the EVLA. including lack of human intervention where possible.," Second, the software must make it easier to use and operate the EVLA, including lack of human intervention where possible."
+ The EVLA monitor and control software design concept is a highly distributed. loosely coupled svstem.," The EVLA monitor and control software design concept is a highly distributed, loosely coupled system."
+ Most. hardware modules contain a lairly capable. single board computer svstem. the Monitor/Control Interface Board. or MID (for the correlator. called CMIDBs. for Correlator MIB).," Most hardware modules contain a fairly capable, single board computer system, the Monitor/Control Interface Board, or MIB (for the correlator, called CMIBs, for Correlator MIB)."
+ The MIDs receive commands via datagram (under the UDP protocol) from the central svstem. and send monitor data by datagram to a central location. which archives il in a general purpose database.," The MIBs receive commands via datagram (under the UDP protocol) from the central system, and send monitor data by datagram to a central location, which archives it in a general purpose database."
+" The MIDs also support a telnet protocol. which provides a uselul debugeine interlace,"," The MIBs also support a telnet protocol, which provides a useful debugging interface."
+ Commands sent to devices/MIDs are labeled with a time at which thev are to be executed., Commands sent to devices/MIBs are labeled with a time at which they are to be executed.
+ Thev are typically sent to the devices several seconds in advance. so the central entity itself does not. normally. have serious real-time deadlines to meet.," They are typically sent to the devices several seconds in advance, so the central entity itself does not, normally, have serious real-time deadlines to meet."
+ Because sell-senerated REI is a major problem for a sensitive. broad. band radio receiving svstenm. a central consideration for the MID design was the minimization of generated RFI.," Because self-generated RFI is a major problem for a sensitive, broad band radio receiving system, a central consideration for the MIB design was the minimization of generated RFI."
+ For that reason. a svslenr-on-a-chip was chosen. rather (han the more common design with memory on an adjacent chip.," For that reason, a system-on-a-chip was chosen, rather than the more common design with memory on an adjacent chip."
+ MIBs also perform limit checking on the monitor points available to them., MIBs also perform limit checking on the monitor points available to them.
+ If a monitor point [alls outside the range of permitted values. an alert is sent bv datagram.," If a monitor point falls outside the range of permitted values, an alert is sent by datagram."
+ These alerts are classified as to whether they are informational. require telescope operator action. or cause (he data to be flageed as ba.," These alerts are classified as to whether they are informational, require telescope operator action, or cause the data to be flagged as bad."
+ There are six central elements of the overall control svstem:, There are six central elements of the overall control system:
+lligh resolution spectral anc astrometric measurements of such stars would also aid in the modeling of the mass distribution in the Galactic center.,High resolution spectral and astrometric measurements of such stars would also aid in the modeling of the mass distribution in the Galactic center.
+ Other consequences of general relativitv. (he possible form of the metric itself (aud. the corresponding theories to which the metric is à solution) could be studied.," Other consequences of general relativity, the possible form of the metric itself (and the corresponding theories to which the metric is a solution) could be studied."
+ The prospect lor measuring general relativistic pericenter precession. an O(:37) effect (where 9 is the pericenter velocity in light units) for the Galactc-center stars is widely appreciated (see.[orexample.??) ancl is anticipated to be observable by interferometric instruments currently under clevelopment (?)..," The prospect for measuring general relativistic pericenter precession, an $\O(\beta^2)$ effect (where $\beta$ is the pericenter velocity in light units) for the Galactic-center stars is widely appreciated \citep[see, for example,][]{jaroszynski, fragile} and is anticipated to be observable by interferometric instruments currently under development \citep{eisenhauer}."
+ Should stars further in be detected in the future. precession effects αἱ Q(3!) would become measurable. enabling tests of no-hair theorems (?)..," Should stars further in be detected in the future, precession effects at $\O(\beta^4)$ would become measurable, enabling tests of no-hair theorems \citep{will08}."
+ Relativity also perturbs the kinematics of a star., Relativity also perturbs the kinematics of a star.
+ ? drew attention to the rapil velocity changes that stus undergo around pericenter passage. and argued that the O(9)3. elect of time dilation in the stars [rame would be measurable as a perturbation of the redshift.," \cite{gillessenredshift} drew attention to the rapid velocity changes that stars undergo around pericenter passage, and argued that the $\O(\beta^2)$ effect of time dilation in the star's frame would be measurable as a perturbation of the redshift."
+" ?. calculated the O03"") kinematic contribution of space curvature g;; and Q(3!) effect of black hole spin gj. suggesting that even the latter may become measurable with [uture interlerometric instruments."," \cite{saha} calculated the $\O(\beta^3)$ kinematic contribution of space curvature $g_{ij}$ and $\O(\beta^4)$ effect of black hole spin $g_{0j}$ , suggesting that even the latter may become measurable with future interferometric instruments."
+ ? presented a new orbit integration method in the presence of space curvature. spin. aud. Newtonian Galactic perturbations.," \cite{preto} presented a new orbit integration method in the presence of space curvature, spin, and Newtonian Galactic perturbations."
+ In this paper. we extend previous work to include the redshift contributions that come from the effect of space curvature and spin on the light path between the star and the observer.," In this paper, we extend previous work to include the redshift contributions that come from the effect of space curvature and spin on the light path between the star and the observer."
+ Specifically. we will compute the redshift of a moving point source in a weak-fiekl approximation of a Ixerr metric. identilv the various contributions. and investigate the scaling of these signals with orbital size.," Specifically, we will compute the redshift of a moving point source in a weak-field approximation of a Kerr metric, identify the various contributions, and investigate the scaling of these signals with orbital size."
+ We formulate (he problem as (wo photons emitted from a stellar orbit. an inlinitesimal proper time interval Ar apart.," We formulate the problem as two photons emitted from a stellar orbit, an infinitesimal proper time interval $\Delta \tau$ apart."
+ Both photons hit the observer. but with a difference A/ in arrival time.," Both photons hit the observer, but with a difference $\Delta t$ in arrival time."
+ The redshift z is then The physical process is the same as in the calculation of the spectrum of an acecretion disk (e.@..?).. only the computational strategv. needed is different.," The redshift $z$ is then The physical process is the same as in the calculation of the spectrum of an accretion disk \citep[e.g.,][]{mueller}, only the computational strategy needed is different."
+ In the accretion-disk case. photons are shot backwards [rom the observer in a range of directions. but in the general direction of the extended. source (the disk).By," In the accretion-disk case, photons are shot backwards from the observer in a range of directions, but in the general direction of the extended source (the disk).By"
+those of theC!8O cores in the OMC-1 region of the Orion A cloud 2009)..,those of the cores in the OMC-1 region of the Orion A cloud \citep{ike09b}.
+" To fairly compare the 9140 cores with the OMC-1 ones, we smoothed the OMC-1 data by a Gaussian function with a FWHM size of 37"".3 so that the OMC-1 cloud is put at the distance of 910 pc to 9140."," To fairly compare the S140 cores with the OMC-1 ones, we smoothed the OMC-1 data by a Gaussian function with a FWHM size of $''$ .3 so that the OMC-1 cloud is put at the distance of 910 pc to S140."
+" Furthermore, we randomly picked up 25 of the original spectra to make the cube data for OMC-1 so as to have the same signal to noise ratio as that of the S140 data."," Furthermore, we randomly picked up 25 of the original spectra to make the cube data for OMC-1 so as to have the same signal to noise ratio as that of the S140 data."
+ Finally we re-identified 44 cores from the smoothed cube data in OMC-1., Finally we re-identified 44 cores from the smoothed cube data in OMC-1.
+" The left panel of Figure 12 shows that the distribution of the cores in $140 has a single peak at 0.37 pc, which is much larger than the peak radius of 0.26 pc in OMC-1: the mean value of 0.36+0.07 pc for in 8140 is larger than that in OMC-1 of 0.27+0.05 pc."," The left panel of Figure \ref{histograms} shows that the distribution of the cores in S140 has a single peak at 0.37 pc, which is much larger than the peak radius of 0.26 pc in OMC-1: the mean value of $\pm$ 0.07 pc for in S140 is larger than that in OMC-1 of $\pm$ 0.05 pc."
+ The Kolmogorov-Smirnov (K-S) test applied to the histograms demonstrates that the distributions are considerably different from each other with a significance level of 1%., The Kolmogorov-Smirnov (K-S) test applied to the histograms demonstrates that the distributions are considerably different from each other with a significance level of 1.
+". The mass of the cores in the $140 region tends to be also larger than that in OMC-1, as shown in the middle panel of Figure 12;; the mean value in $140 of 41429 is twice larger than that in OMC-1 of 18+9Mo."," The mass of the cores in the S140 region tends to be also larger than that in OMC-1, as shown in the middle panel of Figure \ref{histograms}; ; the mean value in S140 of $\pm$ 29 is twice larger than that in OMC-1 of $\pm$ 9."
+" In addition, the K-S test shows that the two distributions are considerably different from each other."," In addition, the K-S test shows that the two distributions are considerably different from each other."
+" Although only one broad peak at c 20 seems to exist for OMC-1, we can see two broad peaks at ~ 20 and ~ 50 for $140."," Although only one broad peak at $\sim$ 20 seems to exist for OMC-1, we can see two broad peaks at $\sim$ 20 and $\sim$ 50 for S140."
+ The above comparisons suggest that the large > 0.4 pc) and massive > 50 Mo) cores exist only in the $140 region., The above comparisons suggest that the large $>$ 0.4 pc) and massive $>$ 50 ) cores exist only in the S140 region.
+" Furthermore, the mass fraction of the cores in 9140 is very large: the total mass of the cores is 5000 Mo, of the total mass traced by the emission (see §3.2))."," Furthermore, the mass fraction of the cores in S140 is very large: the total mass of the cores is 5000 $M_{\odot}$, of the total mass traced by the emission (see \ref{clumpProperties}) )."
+" This fraction is twice larger than that of the cores in OMC-1 (Ikeda&Kitamura2009),, but is roughly consistent with the fraction of 60 for the cores (Ikedaetal.2007,2009)."," This fraction is twice larger than that of the cores in OMC-1 \citep{ike09b}, but is roughly consistent with the fraction of 60 for the cores \citep{ike07,ike09}."
+. The large massive cores and the high mass fraction ofthe cores in the $140 region are likely to be caused by high column, The large massive cores and the high mass fraction ofthe cores in the S140 region are likely to be caused by high column
+The full predictive poteutial of our framework can be realized if it is combines witha detailed analysis of morphological. polarization and multifrequeucy observations.,"The full predictive potential of our framework can be realized if it is combined with a detailed analysis of morphological, polarization and multifrequency observations."
+ We offer one concrete sugeestion in this comuection., We offer one concrete suggestion in this connection.
+ The total intensity profile of a pulsar could be separated iuto individua Gaussian components., The total intensity profile of a pulsar could be separated into individual Gaussian components.
+ The circular polarization signature could be used to identity the core component., The circular polarization signature could be used to identify the core component.
+ A precise determination of the above offset wil then inunediately lead to an altitude difference., A precise determination of the above offset will then immediately lead to an altitude difference.
+ A polarization position auele curve for cach commponcut can be constructed using the rotating vector mode aud |-j0wii à and . values., A polarization position angle curve for each component can be constructed using the rotating vector model and known $\alpha$ and $\beta$ values.
+ A composite of these which reproduces the observe position angle curve can be coustructed., A composite of these which reproduces the observed position angle curve can be constructed.
+ Then the BCW offsets between intensity and the position angle curve centroids for cach component will allow determination of μονxdividual altitudes., Then the BCW offsets between intensity and the position angle curve centroids for each component will allow determination of individual altitudes.
+ It may be noted that the BOW offsets can be unanibieuouslv determined oulv if position augle sweeps are clean. which typically occurs for conal components. (," It may be noted that the BCW offsets can be unambiguously determined only if position angle sweeps are clean, which typically occurs for conal components. ("
+I£ BOW imnothod is applied. aud very huge altitucles are considered. then its modification to include will be needed).,"If BCW method is applied and very large altitudes are considered, then its modification to include will be needed)."
+ With the thus obtained conal altitude combined with A. the core altitude can be fixed., With the thus obtained conal altitude combined with $\Delta_{cc}$ the core altitude can be fixed.
+ The filline factors can then be deteziiued by the 1/3 rule proposed here., The filling factors can then be determined by the 1/3 rule proposed here.
+ Consistency of this information with multfrequeuey. observatious will permit a further look iuto issues like radius to frequency mapping. tine delays etc.," Consistency of this information with multfrequency observations will permit a further look into issues like radius to frequency mapping, time delays etc."
+ We hope to have demonstrated in previous sectious that the kinematical effects considered here have the potential to explain the core/cone longitude offsets aud also shed light ou the nature of the eumissiou niechiauisu. aud therefore a detailed observational study of these offsets holds much promise in clicidating the intricacies of pulsar emission.," We hope to have demonstrated in previous sections that the kinematical effects considered here have the potential to explain the core/cone longitude offsets and also shed light on the nature of the emission mechanism, and therefore a detailed observational study of these offsets holds much promise in elucidating the intricacies of pulsar emission."
+For the proof. see Simon|?.Corollary4.Pageση].,"For the proof, see Simon\cite[Corollary 4, Page 85]{SI87}."
+" In order to show the existence of a solution to svstem (1.1)) in Subsection ??.. we will apply this lemma {ο each scalar component of uw in the particular case where NeW yp. Bab and Y2WMQsareny. We denote by Ανν} the class of all measurable ΠΟΙΟ uv. defined on J. for which. The space PNP(I)={puspe>0andwel,sp)H ihe linear hull of ᾖν,p)."," In order to show the existence of a solution to system \ref{EM:burger}) ) in Subsection \ref{EM:preuvf}, we will apply this lemma to each scalar component of $u^\e$ in the particular case where $X=W^{1,1}(I)$ , $B= L^{1}(I)$ and $Y=W^{-1,1}(I):=(W^{1,\infty}_0(I))'$ We denote by $K_{exp}(I)$ the class of all measurable function $u$, defined on $I$, for which, The space $EXP(I)= \left\{\mu u:\quad \mu\ge 0 \quad \mbox{and} \quad u \in K_{exp}(I) \right\}$ the linear hull of $K_{exp}(I)$ ."
+ This space is supplemented with the following Luxemburg norm (see dams ): Let us recall some useful propertiesof this space.," This space is supplemented with the following Luxemburg norm (see Adams \cite[(13), Page
+234]{Adams} ): Let us recall some useful propertiesof this space."
+ For all h>0 aud (10.7)€(0.7)x R. wehave:," For all $h>0$ and $(t,x)\in (0,T)\times \R$ , wehave:"
+ , .
+where ).," where ), )."
+ The angle cosines in the colinear limit reduce to VOD) and the angular terms are presented in equations (AT5)-(AT8)).," The angle cosines in the colinear limit reduce to _q, _q, _q and the angular terms are presented in equations \ref{TsTsTs-Colinear}) \ref{TtTtTt-Colinear}) )."
+ In particular. note that the tensor auto-correlation explicitly vanishes identically. as was argued before in BCOS.," In particular, note that the tensor auto-correlation explicitly vanishes identically, as was argued before in BC05."
+ In the colinear limit where r=1. this simplifies further.," In the colinear limit where $r=1$, this simplifies further."
+"Under the transformation y;—v, the angular terms TueCuaeasdnd7.;.., we invariant even though the underlying angles change: therefore when Ξ| and ὁ=0 the bispectrum becomesΑπ.","Under the transformation $y_k\rightarrow y_p$ the angular terms $\FTrTrTr$, $\FTsTsTs$ and $\FTtTtTt$ are invariant even though the underlying angles change; therefore when $r=1$ and $\phi=0$ the bispectrum becomes."
+ There is a strong. intuitive difference between a squeezed bispectrum with r=|. which corresponds to the vanishing of q. and a squeezed bispectrum with r>1.," There is a strong, intuitive difference between a squeezed bispectrum with $r=1$, which corresponds to the vanishing of $\mathbf{q}$, and a squeezed bispectrum with $r>1$."
+ At 7= 1. the configuration is approaching a degenerate state where g—0 and at @=z reaches k= -—p.g=0.," At $r=1$ , the configuration is approaching a degenerate state where $q\rightarrow 0$ and at $\phi=\pi$ reaches $k=-p$, $q=0$."
+ At r>|. however. the bispectrum is in a squeezed state with g=p— κ. The symmetries of the bispectrum equate these squeezed configurations with an equivalent colinearbispectrum.," At $r>1$, however, the bispectrum is in a squeezed state with $q=p-k$ The symmetries of the bispectrum equate these squeezed configurations with an equivalent colinearbispectrum."
+ Specifically. settingm," Specifically, setting,"
+ Specifically. settingma," Specifically, setting,"
+ Specifically. settingmat," Specifically, setting,"
+ Specifically. settingmath," Specifically, setting,"
+ Specifically. settingmathb," Specifically, setting,"
+ Specifically. settingmathbf," Specifically, setting,"
+ Specifically. settingmathbfp," Specifically, setting,"
+ Specifically. settingmathbfp.," Specifically, setting,"
+ Specifically. settingmathbfp..," Specifically, setting,"
+10 leading edge of the profile.,the leading edge of the profile.
+ As shown in Fig. d.," As shown in Fig. \ref{fg:1326},"
+ it displays an intriguing combination of nulling. and. mode changing., it displays an intriguing combination of nulling and mode changing.
+ Since it is relatively strong. averaging of data into subintegrations was not required. and Fig.," Since it is relatively strong, averaging of data into subintegrations was not required and Fig."
+ 4. shows individual pulse intensities.," \ref{fg:1326}
+ shows individual pulse intensities."
+ At intervals ο 10 min (200 rm1060 pulses) emission from the two main components ceases. typically for less than a minute. (," At intervals of 2 – 10 min (200 – 1000 pulses) emission from the two main components ceases, typically for less than a minute. ("
+Pable 1)).,Table \ref{tb:obs}) ).
+" During these ""component"" nulls. sporadic emission appears at the leading edge of the profile. forming the fourth leading component of the profile."," During these “component” nulls, sporadic emission appears at the leading edge of the profile, forming the fourth leading component of the profile."
+ Ehe pulse phase of this niocle-B emission is somewhat variable (see. for example. around: pulse 3700 3150 in the lower inset of Fig. 4))," The pulse phase of this mode-B emission is somewhat variable (see, for example, around pulse 3700 – 3750 in the lower inset of Fig. \ref{fg:1326}) )"
+ and there are frequent short intervals during this mode-D emission when this eading component is undetectable., and there are frequent short intervals during this mode-B emission when this leading component is undetectable.
+ The mode A emission is more Consistent with most pulses being visible although he relative (and absolute) amplitudes of the components vary greatly., The mode A emission is more consistent with most pulses being visible although the relative (and absolute) amplitudes of the components vary greatly.
+ Fig., Fig.
+ 5. shows the very different mean pulse xoliles for the two modes., \ref{fg:1326mode} shows the very different mean pulse profiles for the two modes.
+ There is some evidence for occasional drifting subpulses in. this pulsar. for. example. around pulses 3350 and 3560 in the lower inset.," There is some evidence for occasional drifting subpulses in this pulsar, for example, around pulses 3350 and 3560 in the lower inset."
+ The drift is oward the trailing edge with a rate of approximately 0.008 evcles/period., The drift is toward the trailing edge with a rate of approximately 0.008 cycles/period.
+ This pulsar has a narrow profile width of 10 ms, This pulsar has a narrow profile width of 10 ms
+The integrationsIn over velocities can be carried out following the same steps as in the previous section and. we finally have: ,"The integrations over velocities can be carried out following the same steps as in the previous section and we finally have: where $\alpha_{CMS} \equiv \mbox{\boldmath $ $}_{h_{S}} - \mbox{\boldmath $ $}_{j_{B}}) \over \mid \mbox{\boldmath $ $}_{h_{S}} - 
+\mbox{\boldmath $."
+In eq. (3.1)), In eq. \ref{EBSfin}) )
+" 1 is implicitlv. assumed. that the distance »etween a generic satellite particle and a generic background yarticle cannot be smaller than a seale. d,,;,. that. is. hat there exists a minimum effective impact parameter."," it is implicitly assumed that the distance between a generic satellite particle and a generic background particle cannot be smaller than a scale, $d_{min}$, that is, that there exists a minimum effective impact parameter."
+" As the main contribution to the integrals appearing in his expression comes [rom small | 7s rg| values (rs and rg are generic satellite and background: particle »ositions). an accurate determination of the d,,;, scale is a crucial point when studying dynamical friction."," As the main contribution to the integrals appearing in this expression comes from small $\mid \mbox{\boldmath $ $}_{S} - \mbox{\boldmath $ $}_{B}
+\mid$ values $\mbox{\boldmath $ $}_{S}$ and $\mbox{\boldmath $ $}_{B}$ are generic satellite and background particle positions), an accurate determination of the $d_{min}$ scale is a crucial point when studying dynamical friction."
+ This scale. rowever. cannot. be determined by the present approach to his problem.," This scale, however, cannot be determined by the present approach to this problem."
+ There exist in literature several estimates of (us for specific situations (White 1976: Bontekoe van Albaca 1987)., There exist in literature several estimates of $d_{min}$ for specific situations (White 1976; Bontekoe van Albada 1987).
+ White (1976) in fact shows that for spherical symmetric satellites. it is à eood approximation to consider them as point-like svstems with a cut-olf in their distances to background. particles. and that for the case of satellites with a Wing mass distribution (king 1966). this cut-oll is given by divincm where 22) ds the satellite tical radius.," White (1976) in fact shows that for spherical symmetric satellites, it is a good approximation to consider them as point-like systems with a cut-off in their distances to background particles, and that for the case of satellites with a King mass distribution (King 1966), this cut-off is given by $d_{min} \simeq {R_{t} \over 5}$ , where $R_{t}$ is the satellite tidal radius."
+ X pointmass satellite has a density profile. given bv ps(rs)= Mio(rs ποκ) where o(rs— verus) is the delta function in three dimensions and Tes is the position of the satellite center-of-mass.," A pointmass satellite has a density profile given by $\rho_{S} (\mbox{\boldmath $ $}_{S}) = M_{S} \delta(\mbox{\boldmath $ $}_{S}-
+\mbox{\boldmath $ $}_{CMS})$, where $\delta(\mbox{\boldmath $ $}_{S}-\mbox{\boldmath $ $}_{CMS})$ is the delta function in three dimensions and $\mbox{\boldmath $ $}_{CMS}$ is the position of the satellite center-of-mass."
+ With this value of ps(rz). eq. (3.1))," With this value of $\rho_{S} (\mbox{\boldmath $ $}_{S})$, eq. \ref{EBSfin}) )"
+ gives in the point mass limit: where now ocs= poPopeTro and O(r) QO(|rceus diis) is the step function.," gives in the point mass limit: where now $\alpha_{CMS} \equiv \mbox{\boldmath $ $}_{CMS} - \mbox{\boldmath $ $}_{h_{B}}) \over \mid \mbox{\boldmath $ $}_{CMS} 
+- \mbox{\boldmath $ and $ \Theta ( \mbox{\boldmath $ $}) \equiv 
+\Theta~(\mid~\mbox{\boldmath $ $}_{CMS} - 
+\mbox{\boldmath $ $} \mid - d_{min})$ is the step function."
+ When the background density is uniform. we can set pelts)=peo and reas=0 in eq. (50))," When the background density is uniform, we can set $\rho_{B}(\mbox{\boldmath $ $}_{B})
+= \rho_{B,0}$ and $\mbox{\boldmath $ $}_{CMS} = 0$ in eq. \ref{EBSpoint}) )"
+r and then it becomes: ‘This recovers the Chandrasekhar formula for the motion of a massive test particle in an homogeneous background. except for the factor containing the temperature. ratio (which. vanishes in the Chancrasckhar formula because msmg in this limit). and. now. op=(az|o)− at the place of the background. velocity. dispersion.," and then it becomes: This recovers the Chandrasekhar formula for the motion of a massive test particle in an homogeneous background, except for the factor containing the temperature ratio (which vanishes in the Chandrasekhar formula because $m_S \gg m_B$ in this limit), and, now, $\sigma_{T}=(\sigma_S^2 +\sigma_B^2)^{1/2}$ at the place of the background velocity dispersion."
+ In the Figure 1. we plot the ratio. £2. of satellite energy loss in the Chandrasekhar limit with a=0 and with ax#0. assuming that Je/2s=0.," In the Figure \ref{ma_chan} we plot the ratio, $R$, of satellite energy loss in the Chandrasekhar limit with $\sigma_S=0$ and with $\sigma_S \neq 0$, assuming that $T_B/T_S = 0$."
+ As can be seen in this Figure. the elect of a non zero σς increases with increasing as ancl is more importantat low roars.," As can be seen in this Figure, the effect of a non zero $\sigma_S$ increases with increasing $\sigma_S$ and is more importantat low $v_{CMS}$ ."
+ Ehe Chandrasckhar formula always overestimates the dynamical friction force. and in some situations the clleet of neglecting the satellite velocity dispersion could cause an error as high as an order of magnitudo.," The Chandrasekhar formula always overestimates the dynamical friction force, and in some situations the effect of neglecting the satellite velocity dispersion could cause an error as high as an order of magnitude."
+(1) Piggy) and σι are caleulated frou Eqs. (27)),(1) $h_{0(1)}$ and $c_{(1)}$ are calculated from Eqs. \ref{eq:cond5}) )
+" ancl (28)) as a trial value of 3. pg,, audpp, are obtained from Eqs. (23))"," and \ref{eq:cond6}) ) as (2) For a trial value of $\beta$,$p_{E_{12}}$ and $p_{P_{12}}$ are obtained from Eqs. \ref{eq:cond2}) )"
+ and (25)) auc. combined withEqs. (21))," and \ref{eq:cond4}) ) and,combined with Eqs. \ref{eq:cond1}) )"
+" aud (25)). eive figs) as ne Oror Then ei», is obtained from Eq. (21))"," and \ref{eq:cond4}) ), give $h_{0(2)}$ as Then $c_{(2)}$ is obtained from Eq. \ref{eq:cond1}) )"
+ or (26)). (3) The point that gives the largest value το og(r.7/2)|DepOvorlrtoday)ey) is searched in laver 1 and is referred to as point C1.," or \ref{eq:cond3}) (3) The point that gives the largest value to $-\phi_g(r,\pi/2) + h_{0(1)}^2 \phi_{\rm rot}(r:A_{(1)}) + c_{(1)}$ is searched in layer 1 and is referred to as point $C1$."
+ The counter part in laver 2 is then looked for and is called point C2., The counter part in layer 2 is then looked for and is called point $C2$ .
+ The point C is found by comparing CT and C2 to E45., The point $C$ is found by comparing $C1$ and $C2$ to $E_{12}$.
+ The masxinnun value thus obtained is divided by Z7;(Pras)μις With the appropriate / and the updated value of 9 is obtained., The maximum value thus obtained is divided by $H_{(i)}(\rho_{\rm max})/p_{\rm max}$ with the appropriate $i$ and the updated value of $\beta$ is obtained.
+ The steps (2) aud (3) are repeated uutil the value of ο) converges at a sufficient level., The steps (2) and (3) are repeated until the value of $\beta$ converges at a sufficient level.
+ The &ual step of the iteration is the updates of the deusitv distribtion and the location of the layer boundary., The final step of the iteration is the updates of the density distribution and the location of the layer boundary.
+ This is accomplished as follows., This is accomplished as follows.
+ Regarding the specific euthalpy as a funcion of pressure alone. we first solve the Bernoulli equation. Eq. (6)).," Regarding the specific enthalpy as a function of pressure alone, we first solve the Bernoulli equation, Eq. \ref{eq:Ber}) ),"
+ for cach laver in the absence of the other laver αιid the laver boundary as a result: where Il Is an iuverse function of the specific euthalpy for laver ἐν, for each layer in the absence of the other layer and the layer boundary as a result: where $H^{-1}_{(i)}$ is an inverse function of the specific enthalpy for layer $i$.
+ Since the pressure should be continuous at the laver boundary as discussed iu Sec. 27..," Since the pressure should be continuous at the layer boundary as discussed in Sec. \ref{sec:bs},"
+ we look for à point ou cach radial ray (a line with 0= coust.).," we look for a point on each radial ray (a line with $\theta=$ const.),"
+" at which pu, and pio, coincides with each other.", at which $p_{(1)}$ and $p_{(2)}$ coincides with each other.
+ This gives the updated laver boundiry as r=Πο]., This gives the updated layer boundary as $r=r_{b}(\theta)$.
+ Ou the other haud. piti.0) and pro(re) obtained above give the updated density distribution for cach laver.," On the other hand, $p_{(1)}(r,\theta)$ and $p_{(2)}(r,\theta)$ obtained above give the updated density distribution for each layer."
+ This closes the iteration procedure., This closes the iteration procedure.
+ We eo back to the Poisson equation aud repeat all the steps until a sufficient level of convergence is reached., We go back to the Poisson equation and repeat all the steps until a sufficient level of convergence is reached.
+ To demonstrate that the new formula described above really works. we will apply it to two-lavered configurations for sole represcutative rotation laws.," To demonstrate that the new formula described above really works, we will apply it to two-layered configurations for some representative rotation laws."
+ For simplicity. we euiploy. two polytropic EOS pj(p)=N;pl!?ην απ Ισ Jy; aud nj; are the polytropic constants aud indices for laver 7.," For simplicity, we employ two polytropic EOS's, $p_{i}(\rho)=K_i\rho^{1+1/n_i}$, in which $K_i$ and $n_i$ are the polytropic constants and indices for layer $i$."
+ We take rather arbitrarily A4s.=1.015«104? and Ive=)5.41101! in ces unit and 23ο=3., We take rather arbitrarily $K_1=1.015\times 10^{15}$ and $K_2=5.14\times 10^{14}$ in cgs unit and $n_1=n_2=3$.
+" Note that the specific eithalpy is given as LT;;;—(nj|↽↽Lip!""ο-Hels[1litte for laver 7.", Note that the specific enthalpy is given as $H_{(i)}=(n_i+1)K_i \rho^{1/n_i}=n_i K_i^{n_i/n_i+1}p^{1/1+n_i}$ for layer $i$.
+ Adoptiug pis=5.629«107efeu? we obain white dwartlike configurations., Adopting $\rho_{\rm max}=5.629\times 10^9 {\rm g/cm^3}$ we obtain white dwarf-like configurations.
+ We work with the normalized variables (Eqs. (113)-(] T))), We work with the normalized variables (Eqs. \ref{eq:nor1}) \ref{eq:nor2}) ))
+ aud the nunuerical domain covers a quadrant of the meridian section. OSprS Land 086X7/2. under the assuuption of equatorial svnunetry.," and the numerical domain covers a quadrant of the meridian section, $0 \leqq r \leqq 1$ and $0 \leqq \theta \leqq \pi/2$, under the assumption of equatorial symmetry."
+" We typically deploy 1000. mesh points on the r-coordinate aud 200 exid points ou the 0-coordiuate to obtain an acceptable accuracy. which is confirmed, by the normalized virial equation (Cowling1965) defined as where T.W. and C are the rotational. gravitational aud iutcrnal energies. respectively,"," We typically deploy 1000 mesh points on the $r$ -coordinate and 200 grid points on the $\theta$ -coordinate to obtain an acceptable accuracy, which is confirmed by the normalized virial equation \citep{cow} defined as where $T,W,$ and $U$ are the rotational, gravitational and internal energies, respectively."
+ The ratio should vanish for exact solutions., The ratio should vanish for exact solutions.
+ Note finally that the Poisson equation is solved by the Cween fuuction method eimploved by (1956).., Note finally that the Poisson equation is solved by the Green function method employed by \cite{Hachisu:1986}. .
+ As a nuunlatory step. we first construct anon-rotational configuration with two lavers according to the preseut formmla and compare it with the solution obtained by the ordinary and much simpler method. that is. the radial integration.," As a mandatory step, we first construct anon-rotational configuration with two layers according to the present formula and compare it with the solution obtained by the ordinary and much simpler method, that is, the radial integration."
+" Setting r,= Laud re,= rp. we obtain fy)=0 from Eq. (30))", Setting $r_p=1$ and $r_{E_{12}}=r_{P_{12}}$ we obtain $h_{0(1)}=0$ from Eq. \ref{eq:h01}) )
+" aud ντι,=pp, from Eqs. (23))", and $p_{E_{12}}=p_{P_{12}}$ from Eqs. \ref{eq:cond2}) )
+ and, and
+while for the same masses ancl metallicity. ? find 1.51xLO! vy and 6.06x10? wr. and ?.. alter modilving the input physics. find 1.43xLOM vr and 6.85xLO? vr.,"while for the same masses and metallicity, \citet{1996A&AS..115..339C} find $1.51\times10^{10}$ yr and $6.06\times10^9$ yr, and \citet{1999A&AS..135..405C}, after modifying the input physics, find $1.43\times10^{10}$ yr and $6.85\times10^9$ yr."
+ For Pop., For Pop.
+I models. (he spread in initial Z is larger. vel our results are compatible with those obtained by ?. for a grid of stelar models with Z=0.019.,"I models, the spread in initial $Z$ is larger, yet our results are compatible with those obtained by \citet{2007ApJS..172..649S} for a grid of stelar models with $Z=0.019$."
+ The evolutionary tracks end. with a cooling WD., The evolutionary tracks end with a cooling WD.
+ A IIe-WD is obtained for the lower initial masses. 0.25<As;0.50M...," A He-WD is obtained for the lower initial masses, $0.25 \lesssim M_i\lesssim 0.50\ \Msun$."
+ The transition to a CO-WD occurs between 0.50 and 0.80M... and the heavier ONeMg-domiunated: WDs are obtained for initial masses higher than 8M. (the (vausition mass being higher for the Pop.," The transition to a CO-WD occurs between $0.50$ and $0.80\ \Msun$, and the heavier ONeMg-dominated WDs are obtained for initial masses higher than $\sim 8\ \Msun$ (the transition mass being higher for the Pop."
+I stars).,I stars).
+ It should be noted. however. that especially for the Pop.," It should be noted, however, that especially for the Pop."
+I stars. the (ransion mass for obtaining à CO-WD rather than a IIe-WD is strongly dependent on (he mass-loss rate assumed.,"I stars, the transition mass for obtaining a CO-WD rather than a He-WD is strongly dependent on the mass-loss rate assumed."
+ For example. for an initial mass of 0.80.AZ... slightly increasing the mass-loss rate may result either in a Ile-WD. when the threshold for core helium burning is not reached. or. in an Extreme llorizontal Branch (EIIB) star. when a ‘delaved’ core IleF takes place.," For example, for an initial mass of $0.80\ \Msun$, slightly increasing the mass-loss rate may result either in a He-WD, when the threshold for core helium burning is not reached, or, in an Extreme Horizontal Branch (EHB) star, when a `delayed' core HeF takes place."
+ The production of such hot (blue) IIB stars for relatively low initial masses (from around 0.80 to slightly over 1 AL.) and for a range of metallicities will be addressed im a subsequent paper., The production of such hot (blue) HB stars for relatively low initial masses (from around $0.80$ to slightly over $1\ \Msun$ ) and for a range of metallicities will be addressed in a subsequent paper.
+ For both populations. a violent ignition of helium Cakes place in the core (but. usually oll-center. because of neutrino cooling) at the tip of the first giant branch lor masses in the range 0.80—2M...," For both populations, a violent ignition of helium takes place in the core (but usually off-center, because of neutrino cooling) at the tip of the first giant branch for masses in the range $0.80 - 2\ \Msun$."
+ This is the well-known core IleF. The transition between IIeF. and «quiet Ile ignilion occurs al an initial mass between 2 and 3M.. depending mainiv on composition and mass-loss rate.," This is the well-known core HeF. The transition between HeF and quiet He ignition occurs at an initial mass between $2$ and $3\ \Msun$, depending mainly on composition and mass-loss rate."
+ During the flash. the peak nuclear energy. generation rate is in the range 5x10*€Lineman&10L.. decreasing with increasing initial mass. due lo a corresponding decrease in (he degree of electron clegeneracy of (he core material.," During the flash, the peak nuclear energy generation rate is in the range $5\times 10^7 \lesssim L_{nuc,max} \lesssim 5\times 10^9\ \Lsun$, decreasing with increasing initial mass, due to a corresponding decrease in the degree of electron degeneracy of the core material."
+ lt is worth noting that the huninositv of the star during the flash is unallected by what is Lalàng place in the core. despite the huge nuclear luminosity. which surpasses the luminosity obtained at anv evolutionary stage.," It is worth noting that the luminosity of the star during the flash is unaffected by what is taking place in the core, despite the huge nuclear luminosity, which surpasses the luminosity obtained at any evolutionary stage."
+ The overall duration of the flash (when μις is in excess oL. sav. LO? L.) is of the order of a [ew vears.," The overall duration of the flash (when $L_{nuc}$ is in excess of, say, $10^5\ \Lsun$ ) is of the order of a few years."
+ We note that during this stage tine steps are automatically reduced down to days. then hours and minutes.," We note that during this stage time steps are automatically reduced down to days, then hours and minutes."
+ Once the flash is over. it will Lake some extra 107 to 10? vears before the star settles into stable core Ile burning. the IIB phase.," Once the flash is over, it will take some extra $10^3$ to $10^5$ years before the star settles into stable core He burning, the HB phase."
+ The well-known thermal pulses (hat arise as a result of (he double shell-burning instability. are Clearly seen in the evolutionary sequences during the final stages of the AGL.," The well-known thermal pulses that arise as a result of the double shell-burning instability, are clearly seen in the evolutionary sequences during the final stages of the AGB."
+ Fig., Fig.
+ 6 shows a typical example for a Pop., \ref{fig:tpagb} shows a typical example for a Pop.
+.. 2M. model.,"II, $2\ \Msun$ model."
+ The thermal pulses in this example span about 7xLO? vr. and clear trends are evident. such as the monotonic decrease in effective," The thermal pulses in this example span about $7\times 10^5$ yr, and clear trends are evident, such as the monotonic decrease in effective"
+"wavelength dependence of the intrinsic size becomes steeper (xVT),",wavelength dependence of the intrinsic size becomes steeper $\propto \lambda^{1.4 -1.5}$ ).
+ The analysis of the closure phases at the highest frequency available in these observations GGHz) confirmed that the apparent VLBI source structure is indeed symmetric., The analysis of the closure phases at the highest frequency available in these observations GHz) confirmed that the apparent VLBI source structure is indeed symmetric.
+ We founc o evidence of time variability of the closure phases from timescales of minutes to days., We found no evidence of time variability of the closure phases from timescales of minutes to days.
+ The lack of any visible structural variation or component ejection after an NIR flare points to a upper velocity limit in the range of of ce. First. an a-priort amplitude calibration was performed using measurements of the antenna gain and system temperature for each telescope.," The lack of any visible structural variation or component ejection after an NIR flare points to an upper velocity limit in the range of of c. First, an a-priori amplitude calibration was performed using measurements of the antenna gain and system temperature for each telescope."
+ The system temperatures and station gains were then corrected for atmospheric opacity effects., The system temperatures and station gains were then corrected for atmospheric opacity effects.
+ For each station. the zenith opacity was found to vary slightly over epochs. but not in a systematic manner that could produce the final source flux density variations.," For each station, the zenith opacity was found to vary slightly over epochs, but not in a systematic manner that could produce the final source flux density variations."
+ That the secondary calibrators showed almost no variations in total flux. further reassures that the amplitude variations seen for AA* are real.," That the secondary calibrators showed almost no variations in total flux, further reassures that the amplitude variations seen for A* are real."
+ To remove effects of pointing or focus errors. which in principle could lead to low amplitudes of individual VLBI scans. particularly at (?).. ? ," To remove effects of pointing or focus errors, which in principle could lead to low amplitudes of individual VLBI scans, particularly at \citep{2001AJ....121.2610D}, \citet{2004Sci...304..704B} "
+immediately after70:: total exposure time was 27 s. During both the CAHA and Keck observations the weather was photometric and the average seeing was 0788 (CAHA) and 0755 (Keck) in the K-band.,immediately after; total exposure time was 27 s. During both the CAHA and Keck observations the weather was photometric and the average seeing was 8 (CAHA) and 5 (Keck) in the $K$ -band.
+ Raw data were reduced in a standard fashion and included both sky subtraction. flat-fielding. and combination of individual frames.," Raw data were reduced in a standard fashion and included both sky subtraction, flat-fielding, and combination of individual frames."
+ We performed point-spread-funetion (PSF) photometry usingirar!., We performed point-spread-function (PSF) photometry using.
+. CAHA and Keck instrumental magnitudes were transformed into observed magnitudes using between | and 33 2MASS sources (depending on the filter) that were present within the fields of view., CAHA and Keck instrumental magnitudes were transformed into observed magnitudes using between 1 and 33 2MASS sources (depending on the filter) that were present within the fields of view.
+ The dispersion of the photometric zero point is +£0.05 mag in all bands., The dispersion of the photometric zero point is $\pm$ 0.05 mag in all bands.
+ T dwarfs exhibit complex spectral energy distributions within the near-infrared., T dwarfs exhibit complex spectral energy distributions within the near-infrared.
+" To account for this effect and in order to transform the CAHA and Keck magnitudes into the 2MASS photometric system. we applied the following correction factors: 0.04 mag (J). 40.02 mag (H). and +0.11 mag (K,) for CAHA data. +0.02 mag (J) and —0.03 mag (K) for Keck data."," To account for this effect and in order to transform the CAHA and Keck magnitudes into the 2MASS photometric system, we applied the following correction factors: $-$ 0.04 mag $J$ ), $+$ 0.02 mag $H$ ), and $+$ 0.11 mag $K_s$ ) for CAHA data, $+$ 0.02 mag $J$ ) and $-$ 0.03 mag $K'$ ) for Keck data."
+ These were obtamed via the integration. of observed near-infrared spectra of both main-sequence stars of different spectral types as well as T dwarfs (Pickles 1998:: Leggett et al. 2002::, These were obtained via the integration of observed near-infrared spectra of both main-sequence stars of different spectral types as well as T dwarfs (Pickles \cite{pickles98}; Leggett et al. \cite{leggett02};
+ Geballe et al. 2002).," Geballe et al. \cite{geballe02}) ),"
+ convolved with the Omega-2000. NIRSPEC and 2MASS filter passbands and atmospheric. transmissio1 curves of the observatories.," convolved with the Omega-2000, NIRSPEC and 2MASS filter passbands and atmospheric transmission curves of the observatories."
+ All filter passbands are plotted in Fig. Al.., All filter passbands are plotted in Fig. \ref{filters}.
+ Corrections are expected to be larger for the J and K filters than for the H-band., Corrections are expected to be larger for the $J$ and $K$ filters than for the $H$ -band.
+ Nevertheless. we note that color corrections are similar or slightly smaller thai the measured photometric uncertainties.," Nevertheless, we note that color corrections are similar or slightly smaller than the measured photometric uncertainties."
+ CAHA and Keck magnitudes (2MASS photometric system) of are reported in Table 2.., CAHA and Keck magnitudes (2MASS photometric system) of are reported in Table \ref{phot}.
+ As a test of consistency. we also calibratec the Keck / magnitude of using the observations of J0559—]4 (T4.5. any color effect should be minimal). deriving J + 00.05 mag. which coincides within the observed errors with the values given in Table 2..," As a test of consistency, we also calibrated the Keck $J$ magnitude of using the observations of $-$ 14 (T4.5, any color effect should be minimal), deriving $J$ $\pm$ 0.05 mag, which coincides within the observed errors with the values given in Table \ref{phot}."
+ The excellent agreement in / and K between the CAHA and Keck data indicates that no strong sistematic errors are affecting. the measurements., The excellent agreement in $J$ and $K$ between the CAHA and Keck data indicates that no strong sistematic errors are affecting the measurements.
+ Our data suggests that does not show significant photometric variability over a few days (0.1 mag)., Our data suggests that does not show significant photometric variability over a few days $\le$ 0.1 mag).
+ We derived mid-infrared photometry using public mosaic images available from the Spitzer Space Telescope Data Archive. taken on 2004 October 9 under the Spitzer Guaranteed Time Observation program #337.," We derived mid-infrared photometry using public mosaic images available from the Spitzer Space Telescope Data Archive, taken on 2004 October 9 under the Spitzer Guaranteed Time Observation program 37."
+ Hernánndez et al. (2007)), Hernánndez et al. \cite{hernandez07}) )
+ provides information on data aquisition. exposure times. and data reduction.," provides information on data aquisition, exposure times, and data reduction."
+ We downloaded mosaic images in the four channels (3.6. 4.5. 5.8. and 8.0 gam) of the IRAC instrument (Fazio et al. 2004))," We downloaded mosaic images in the four channels (3.6, 4.5, 5.8, and 8.0 $\mu$ m) of the IRAC instrument (Fazio et al. \cite{fazio04}) )"
+ using the Leopard software., using the Leopard software.
+ Is only detected in the [3.6] and [4.5] bands with noise ratios of 9 and 16 in the peak flux., is only detected in the [3.6] and [4.5] bands with signal-to-noise ratios of 9 and 16 in the peak flux.
+" We performed aperture photometry on various single. bright sources. near using deophot inigAr. with an aperture radius of and a background annulus extending from 12 to24""."," We performed aperture photometry on various single, bright sources near using $daophot$ in, with an aperture radius of and a background annulus extending from 12 to."
+. For the [3.6] and [4.5] bands we adopteddaophot zero point magnitudes of 17.30 and 16.82. respectively. as described at http:;//ssc.spitzer.caltech.edu/archanaly/quick.phot.," For the [3.6] and [4.5] bands we adopted zero point magnitudes of 17.30 and 16.82, respectively, as described at http://ssc.spitzer.caltech.edu/archanaly/quick.phot."
+ We derived the PSF from the brighter sources and applied this to determine PSF photometry for70., We derived the PSF from the brighter sources and applied this to determine PSF photometry for.
+. Results are reported in Table 2.., Results are reported in Table \ref{phot}.
+ We determined the proper motion of via the comparison of its position. with. respect. to. four common point-like sources m various near-infrared images obtained at different epochs., We determined the proper motion of via the comparison of its position with respect to four common point-like sources in various near-infrared images obtained at different epochs.
+" As a starting point for the astrometric analysis. we used the A,-band frame shown in Fig."," As a starting point for the astrometric analysis, we used the $K_s$ -band frame shown in Fig."
+ | of Zapatero Osorio et al. (2002a)).," 1 of Zapatero Osorio et al. \cite{osorio02a}) ),"
+" which was obtained on 2001 December 29 with the Keck I telescope and has an average seeing of 0/88 and a field of view of 327x32""..", which was obtained on 2001 December 29 with the Keck I telescope and has an average seeing of 8 and a field of view of $\times$.
+ The angular resolution was 07115 pix!., The angular resolution was 15 $^{-1}$.
+" The three recent CAHA JHK, images and the Keck / and A’ frames reported here act as the second epoch astrometric data.", The three recent CAHA $JHK_s$ images and the Keck $J$ and $K'$ frames reported here act as the second epoch astrometric data.
+ Hence. the time interval spanned by the available observations is 3.82 yr.," Hence, the time interval spanned by the available observations is 3.82 yr."
+ Our proper motion measurements are shown in Table 3 using the CAHA and Keck data separately as well as the combined result., Our proper motion measurements are shown in Table \ref{pm} using the CAHA and Keck data separately as well as the combined result.
+ The associated uncertainty is derived from the dispersion of the four, The associated uncertainty is derived from the dispersion of the four
+Because monitoring teams are already identifving shor(-cluration events. identification of planets. brown dwarls. ancl hieh-velocity stars can start immecdiatelv.,"Because monitoring teams are already identifying short-duration events, identification of planets, brown dwarfs, and high-velocity stars can start immediately."
+" This is demonstrated by Gould et (2009). although for an event that had a duration on the longer end of ""short-duration events."," This is demonstrated by Gould et (2009), although for an event that had a duration on the longer end of “short-duration” events."
+ Current programs could discover. every. vear. 45—69 short evenis caused by brown clwarls. 23—36 short events caused by wide-orbit planets. 8—13 short events caused by high-speed neutron stars. and perhaps additional events caused bv hypervelocity stars or their remnants.," Current programs could discover, every year, $45-69$ short events caused by brown dwarfs, $23-36$ short events caused by wide-orbit planets, $8-13$ short events caused by high-speed neutron stars, and perhaps additional events caused by hypervelocity stars or their remnants."
+ Comprehensive model tesis of a large fraction of (hese events. including the planet events. will be possible.," Comprehensive model tests of a large fraction of these events, including the planet events, will be possible."
+ Nearby planetary systems (12-18 per vear) may be susceptible to follow-up observations. including radial velocity or transit studies.," Nearby planetary systems (12-18 per year) may be susceptible to follow-up observations, including radial velocity or transit studies."
+ Pan-STARRS is about to begin taking data over a large region using several sampling strategies (Chambers et 2004)., Pan-STARRS is about to begin taking data over a large region using several sampling strategies (Chambers et 2004).
+ For example. regular monitoring will take place over roughly 70 square degrees in its Medium Deep Survey.," For example, regular monitoring will take place over roughly 70 square degrees in its Medium Deep Survey."
+ Additional wide-field coverage will eventually be provided by LSST. which will cover ~20.000 degrees every three to four nights.," Additional wide-field coverage will eventually be provided by LSST, which will cover $\sim 20,000$ degrees every three to four nights."
+ While the cadence of these programs is not ideal for the detection of short events. the efficiency [for the detection of events caused by brown dwarls. high-velocity stars. and the most massive planets will be high.," While the cadence of these programs is not ideal for the detection of short events, the efficiency for the detection of events caused by brown dwarfs, high-velocity stars, and the most massive planets will be high."
+ Once identified by wide-fiekd monitoring programs. more lrequent follow-up is required to [fully characterize the light curve.," Once identified by wide-field monitoring programs, more frequent follow-up is required to fully characterize the light curve."
+ For events of shorter duration. the efficiency of event detection and identification will be lower.," For events of shorter duration, the efficiency of event detection and identification will be lower."
+" If, however. as part of each independent scan. different filters are used in a sequence over a {ime interval curing which the magnification is changing. even very. short-duration events could be identified as targets for more frequent monitoring."," If, however, as part of each independent scan, different filters are used in a sequence over a time interval during which the magnification is changing, even very short-duration events could be identified as targets for more frequent monitoring."
+ Phe masxinnun value of the efficiency would be: f75/.M. where f is à number in the range 0.1—1. A/ is the time between distinet sequences of observations.," The maximum value of the efficiency would be: $f\, \tau_E/\Delta t,$ where $f$ is a number in the range $0.1-1,$ $\Delta t$ is the time between distinct sequences of observations."
+ In LSST and portions of the Pan-STARRs fields. A? will be 3—6 davs.," In LSST and portions of the Pan-STARRs fields, $\Delta t$ will be $3-6$ days."
+ Even efficiencies of a percent are significant. because the total area monitored will be large enough that many lensing events are expected. in front οἱ both dense and sparse fields (Di Stefano 2003b. Han 2008).," Even efficiencies of a percent are significant, because the total area monitored will be large enough that many lensing events are expected, in front of both dense and sparse fields (Di Stefano 2008b, Han 2008)."
+ It is likely Chat. within five vears. a program ideal for (he detection of short-duration events will begin.," It is likely that, within five years, a program ideal for the detection of short-duration events will begin."
+ Funding was recently approved for theNel). which plans to use 3 telescopes across the Southern IHenmisphere to provide monitoring with I0-minute cadence of approximately 16 sq.," Funding was recently approved for the, which plans to use $3$ telescopes across the Southern Hemisphere to provide monitoring with 10-minute cadence of approximately $16$ sq."
+ degrees (Ilan 2009)., degrees (Han 2009).
+ As Figure 2 shows. KAITNet can provide excellent monitoring for Earth-mass planets.," As Figure 2 shows, KMTNet can provide excellent monitoring for Earth-mass planets."
+ I1 could even discover dwarf planets with Einstein-diameter crossing (mes of less than an hour. although in such cases. follow-up bv other observers would be required to obtain eood sampling.," It could even discover dwarf planets with Einstein-diameter crossing times of less than an hour, although in such cases, follow-up by other observers would be required to obtain good sampling."
+ Below we summarize the steps that can be taken in the immediate füture by monitoring teams already in action., Below we summarize the steps that can be taken in the immediate future by monitoring teams already in action.
+ At the end of the section we discuss work that could be conducted with archival data., At the end of the section we discuss work that could be conducted with archival data.
+where J;;=Ji. and we have made use of the ollowing identity for theJj outer product: As the Stokes parameters are used. almost. universally in radio polarimetry. we now make the conversion from he 'óhereney representation. to. the Stokes visibility representation.,"where $\textbf{J}_{ij} \equiv \textbf{J}_i \otimes \textbf{J}_j^*$ and we have made use of the following identity for the outer product: As the Stokes parameters are used almost universally in radio polarimetry, we now make the conversion from the coherency representation to the Stokes visibility representation."
+ A distinction should be made at this point: he Stokes peremclers.. 1. Q. © and V. are used. when describing the intensity or brightness of radiation. whereas he Stokes are used. to represent. the complex visibilities as measured by an interferometer.," A distinction should be made at this point: the Stokes , $I$ , $Q$, $U$ and $V$, are used when describing the intensity or brightness of radiation, whereas the Stokes are used to represent the complex visibilities as measured by an interferometer."
+ The context in which these terms are used provides the means for disambiguation., The context in which these terms are used provides the means for disambiguation.
+ The Stokes visibility vector as measured by a single baseline made up of antennas ὁ and. j. δε=Qij2UiMEE eis related to the coherency vector via a simpledij coordinate transformation. S. through where the inverse of which is Note that the complex vector input signals ancl Jones matrices are two-dimensional whereas the coherence vectors and Stokes visibility vectors are four-climensional.," The Stokes visibility vector as measured by a single baseline made up of antennas $i$ and $j$, $\textbf{s}_{ij} = [\,I_{ij}\;Q_{ij}\;U_{ij}\;V_{ij}\,]^{\rm T}$ , is related to the coherency vector via a simple coordinate transformation, $\textbf{S}$, through where the inverse of which is Note that the complex vector input signals and Jones matrices are two-dimensional whereas the coherence vectors and Stokes visibility vectors are four-dimensional."
+" Operators that act in cohereney- or Stokes-space can be transformed. between representations using the same transformation matrix. S. used for the coherency and Stokes vectors. through where A""! TOMis the operator in. cohereney representation. and Ab”!is is the operator in Stokes representation (operators in the Stokes representation that are used. to. reproduce the elfect of one or more non-depolarising. transforming elements are called Mueller matrices)."," Operators that act in coherency- or Stokes-space can be transformed between representations using the same transformation matrix, $\textbf{S}$, used for the coherency and Stokes vectors, through where $\textbf{A}^{\rm (C)}$ is the operator in coherency representation and $\textbf{A}^{\rm (S)}$ is the operator in Stokes representation (operators in the Stokes representation that are used to reproduce the effect of one or more non-depolarising, transforming elements are called Mueller matrices)."
+ Therefore. we can write the full Mueller matrix for the single baseline svstem as Assuming the signal through both interferometic arms experience the same degree. of Faraday rotation Ay (a valid assumption for baselines shorter than the scale of Iluctuations in the ionosphere). applying the coherencyStokes visibility transformation on the Faraday. rotation operators of a single baseline gives This is equivalent to a rotation of the P=(Q.U]? vector through an angle of 2A\ radian.," Therefore, we can write the full Mueller matrix for the single baseline system as Assuming the signal through both interferometic arms experience the same degree of Faraday rotation $\Delta\chi$ (a valid assumption for baselines shorter than the scale of fluctuations in the ionosphere), applying the coherency--Stokes visibility transformation on the Faraday rotation operators of a single baseline gives This is equivalent to a rotation of the $P = [Q,U]^{\rm T}$ vector through an angle of $2\Delta\chi$ radian."
+ So we see that Faraday rotation in the reduced Stokes representation (6) and ( only) is the same as that in the complex vector representation (where F acts), So we see that Faraday rotation in the reduced Stokes representation $Q$ and $U$ only) is the same as that in the complex vector representation (where $\textbf{F}$ acts).
+ This is reconcilable with Equation (3)) since P? is rotated hy 2AY=20A? from an initial angle 2x0 resulting in a final angle of 24=2yu|Αλ., This is reconcilable with Equation \ref{eq:P}) ) since $P$ is rotated by $2\Delta\chi = 2\phi\lambda^2$ from an initial angle $2\chi_0$ resulting in a final angle of $2\chi = 2\chi_0 + 2\Delta\chi$.
+" Vhe clffect o£. thin Faraday screens on incident svnchrotron emission can be modelled by where Ni is the number of thin Faraday. screens. and Ve=Ok|On(At""Ma). (\o)e being the intrinsic polarisation angle at the Ath screen."," The efffect of thin Faraday screens on incident synchrotron emission can be modelled by where $N_{\rm F}$ is the number of thin Faraday screens and $\chi_{jk} = (\chi_0)_k + \phi_k(\lambda_j^2-\lambda_0^2)$, $(\chi_0)_k$ being the intrinsic polarisation angle at the $k$ th screen."
+ In practice. feeds are non-ideal and this leads to leakage οους all Stokes parameters. {νι(Q.C. and V.," In practice, feeds are non-ideal and this leads to leakage between all Stokes parameters, $I,\,Q,\,U$ and $V$."
+ The reatment of instrumental polarisation leakage is similar o that as for Faraday rotation. but. complicated. by the act that in general the feed errors corrupt all four Stokes xwameters and cannot be assumed to be the same in all antennas.," The treatment of instrumental polarisation leakage is similar to that as for Faraday rotation, but complicated by the fact that in general the feed errors corrupt all four Stokes parameters and cannot be assumed to be the same in all antennas."
+ Assuming equal weights for all Stokes visibilities gives he following Mueller. operator for the full interferometric svsteor. where both the baseline and its conjugate baseline are sununed over in order to give a real-valued result.," Assuming equal weights for all Stokes visibilities gives the following Mueller operator for the full interferometric system, where both the baseline and its conjugate baseline are summed over in order to give a real-valued result."
+ Using Equation (47)). 10 Mueller. matrix describing leakage for a baseline with parallellinear [ους(ie..," Using Equation \ref{eq:JStokes}) ), the Mueller matrix describing leakage for a baseline with parallellinear feeds(i.e.,"
+" cross-cipole {ους in antenna ὁ are approximately parallel to those in antenna j) to first order in leakage ancl gain error is given by where I, is the 4 identity matrix."," cross-dipole feeds in antenna $i$ are approximately parallel to those in antenna $j$ ) to first order in leakage and gain error is given by where $\textbf{I}_4$ is the $4 \times 4$ identity matrix,"
+Alfvóun coutinuun atf a ΤΟΝΟΙ surface. associated with the element At). av shift away from the local Doppler-shitted frequency. &e; because of the (1) geodesic curvature. involving its square wp. (2) compressibility ad finite xessure (5p). aud (3) the pressure variation Ol a fiux surface.,"Alfvénn continuum at a resonant surface, associated with the element $K(\vartheta)$, may shift away from the local Doppler-shifted frequency $kv_{z}$ because of the (1) geodesic curvature, involving its square $\kg^{2}$, (2) compressibility and finite pressure $\gamma p$ ), and (3) the pressure variation on a flux surface."
+ All these effects suggest that au Alfvénn gap can appear around the local Doppler-shifted frequency aud that Alfvénn-slow coupling should be taken iuto account for these eaps., All these effects suggest that an Alfvénn gap can appear around the local Doppler-shifted frequency and that Alfvénn-slow coupling should be taken into account for these gaps.
+ This kind of coupling is investigated in the uext section., This kind of coupling is investigated in the next section.
+ Iu previous sections. we lave derived equations for the continuous MIID spectrum of a translationally svuuucetric equilibriu.," In previous sections, we have derived equations for the continuous MHD spectrum of a translationally symmetric equilibrium."
+ These derivations were completely ecueral. and allowed us to ideutifv au imstabilitv criterion ou the one hand. as well as to identify the terms responsible for A]lfvóun-slow type coupling.," These derivations were completely general, and allowed us to identify an instability criterion on the one hand, as well as to identify the terms responsible for Alfvénn-slow type coupling."
+S In this section. we quautifv| ↑↕∐↴∖↴↸⊳≺∏∏≻↕∐∶↴∙⊾↸∖↕−↥⋅↸∖↸⊳↑∙∏↴∖↴↕∐∶↴∙⊾↑∐↸∖↴∖↴⊔⋜↧∐∶↴∙⊾↥⋅⋜↧↖⇁↕↑⋅↖⇁↻⋜∐⋅⋜⊔⊔↸∖↑↸∖↥⋅ expansion for the equilibrimu. similar to those discussed in paper L Iu the remainder of this paper. we assume a gravitational potential in which the prominence is enibedded given by where sy is the location of the ceuter of the last closed flux surface of the prominence and the role of eravitv is represented by the constaut g.," In this section, we quantify this coupling effect, using the small gravity parameter expansion for the equilibrium, similar to those discussed in paper I. In the remainder of this paper, we assume a gravitational potential in which the prominence is embedded given by where $x_{0}$ is the location of the center of the last closed flux surface of the prominence and the role of gravity is represented by the constant $g$."
+ This gravity paralcter g denotes a dimensionless quantification of the relative nuportauce of eravity «1 the equilibrimu structure. aud we explained in paper [that this parameter can meaniuefullv vary by orders of magnitude for solar pronunence coufigurations.," This gravity parameter $g$ denotes a dimensionless quantification of the relative importance of gravity on the equilibrium structure, and we explained in paper I that this parameter can meaningfully vary by orders of magnitude for solar prominence configurations."
+ For small g aud assunüug a circular cross-section of the prominence plasma. we quantified in paper I the change induced in the equilibrium (the downwarcl shift of the flux surfaces). and the corresponding straight field ine coordinates.," For small $g$ and assuming a circular cross-section of the prominence plasma, we quantified in paper I the change induced in the equilibrium (the downward shift of the flux surfaces), and the corresponding straight field line coordinates."
+ This kiud of expansion is similar to the small inverse aspect ratio expansion for tokamals plasiuas., This kind of expansion is similar to the small inverse aspect ratio expansion for tokamak plasmas.
+ Using this analytical result for the equilibrium. we can now demonstrate that gaps (or even instabilities) may appear in the continuous ATID spectrum because of uode coupling. as a direct consequence of the force of gravity.," Using this analytical result for the equilibrium, we can now demonstrate that gaps (or even instabilities) may appear in the continuous MHD spectrum because of mode coupling, as a direct consequence of the force of gravity."
+ We apply the iuall gravity expiusiou to the expressions found iu Eqs. (38)... (," We apply the small gravity expansion to the expressions found in Eqs. , ,"
+39)... C10)... and for the matrix elements Wd). M(U). N(U). aud LG). respectively.," and for the matrix elements $K(\vartheta)$, $M(\vartheta)$, $N(\vartheta)$, and $L(\vartheta)$, respectively."
+ For this expansion. we rely ou the metric elements that are valid for straight field line coordinates exploiting the Shafranov shift Αη). eiven by Eq. (," For this expansion, we rely on the metric elements that are valid for straight field line coordinates exploiting the Shafranov shift $\Delta(r)$, given by Eq. ("
+26) in our accolpanving paper (23..,26) in our accompanying paper \citep{Blokland_2011A}.
+ In addition. one can derive the quautitics where the total magnetic field BT=v7|DP.," In addition, one can derive the quantities where the total magnetic field $B_{0}^{2} = {\psi'}^{2} + I^{2}$."
+" The parameter f,=E has been introduced to track the luear and quadratic dependence on the geodesic curvature.", The parameter $t_{g}=1$ has been introduced to track the linear and quadratic dependence on the geodesic curvature.
+" Firthermore. the parameter f, allows us to distinguish the three options for the fux function dependence in the equilibrium. namely These flux functions lead to different variants of the extended Grad-Shatfranoy equation in Eq. ("," Furthermore, the parameter $t_{e}$ allows us to distinguish the three options for the flux function dependence in the equilibrium, namely These flux functions lead to different variants of the extended Grad-Shafranov equation in Eq. ("
+10) for the equilibriiuu. as presented in our accompanying paper (?,"10) for the equilibrium, as presented in our accompanying paper \citep{Blokland_2011A}."
+ We again note that the maeuetically modified Drunut-Visallà frequency NE fron Eq.," We again note that the magnetically modified Brunt-V\""aiis\""all\""a frequency $\Nmpol$ from Eq."
+ can only be uceative if the deusitv is a flux function., can only be negative if the density is a flux function.
+ This once mere coufirias that or this particular choice of tle flux fiction. the MIID contiuua may become unstable as already found to be true in general above.," This once more confirms that for this particular choice of the flux function, the MHD continua may become unstable as already found to be true in general above."
+ ApplyingC» the smallC»eravityexpansion.the eigenvalue[m] equation in Eq.," Applying the smallgravityexpansion,the eigenvalue equation in Eq."
+ cau be written as where, can be written as where
+spectra (Paerels ct al.,"spectra (Paerels et al.,"
+ 2001) ancl detection. of € ancl ο lines in the optical (Nelemans et al., 2001) and detection of C and O lines in the optical (Nelemans et al.
+ 2004) has led to the suggestion that 4U0614]091 is an ultracompact X-ray binary (ΕςΝΑ) with a bydrogen and helium. deficient donor star Le. à C/O white dwarf., 2004) has led to the suggestion that 4U0614+091 is an ultracompact X-ray binary (UCXRB) with a hydrogen and helium deficient donor star i.e. a C/O white dwarf.
+ The direct implication ofa WD mass donor is that the orbital period of the svsten has to be less than  1 hour., The direct implication of a WD mass donor is that the orbital period of the system has to be less than $\sim$ 1 hour.
+ Nelemans et al. (, Nelemans et al. (
+2006) obtained phase resolved spectra over the tentative (unpublished) 50r min period and found a weak sinusoidal radial velocity signal at 48.5miin in the line at4960A.,2006) obtained phase resolved spectra over the tentative (unpublished) 50 min period and found a weak sinusoidal radial velocity signal at 48.5min in the line at.
+. Llowever. the signal was not detected in their period analysis.," However, the signal was not detected in their period analysis."
+ Lf real. the amplitude of this variation would imply an inclination of ~ 15 degrees assuming a neutron star mass of 1.4M.," If real, the amplitude of this variation would imply an inclination of $\sim$ 15 degrees assuming a neutron star mass of $\odot$."
+* Schulz et al. (, Schulz et al. (
+2010) analyzed 200 ksec of Chandra transmission grating cata and find that relativistic elfects needed: to be included. to Lit the X-ray spectra.,2010) analyzed 200 ksec of Chandra transmission grating data and find that relativistic effects needed to be included to fit the X-ray spectra.
+ Their fits. the relativistic line profiles in particular. imply that the system inclination should be very high i.e. 86-89 degrees.," Their fits, the relativistic line profiles in particular, imply that the system inclination should be very high i.e. 86-89 degrees."
+ This is rather contradictory though. since the svstem does not show any evidence for even partial eclipses. neither in optical nor N-ravs.," This is rather contradictory though, since the system does not show any evidence for even partial eclipses, neither in optical nor X-rays."
+ Another peculiarity associated with the system is that its X-rav emission seems to be anticorrelated with the optical emission (Machin ct al., Another peculiarity associated with the system is that its X-ray emission seems to be anticorrelated with the optical emission (Machin et al.
+ 1990)., 1990).
+ Moreover. the X-rays become harder with increasing N-rav luminosity ruling out any photoclectric absorption effects as à cause for changes in the X-rav Lux.," Moreover, the X-rays become harder with increasing X-ray luminosity ruling out any photoelectric absorption effects as a cause for changes in the X-ray flux."
+ Recently Shahbaz ct al. (, Recently Shahbaz et al. (
+2008) have published the results of their photometric monitoring programme.,2008) have published the results of their photometric monitoring programme.
+ They founcl evidence for three dillerent. periods ic. 42. 51.3 and 64 minutes in their data. with the 51.3 min period showing the clearest case of periodicity in terms of signal-to-noise.," They found evidence for three different periods i.e. 42, 51.3 and 64 minutes in their data, with the 51.3 min period showing the clearest case of periodicity in terms of signal-to-noise."
+ still. the 51.8 min period is only present in some of their data. whilst other observations show evidence for periods at 42.0 and 64.1 minutes.," Still, the 51.3 min period is only present in some of their data, whilst other observations show evidence for periods at 42.0 and 64.1 minutes."
+ Clearly much more data is required., Clearly much more data is required.
+ We have obtained. fast optical photometry of 4U0614|091 over several epochs spanning over ten vears in order to search [or the orbital period. and attempt to unravel its remarkable temporal variability.," We have obtained fast optical photometry of 4U0614+091 over several epochs spanning over ten years in order to search for the orbital period, and attempt to unravel its remarkable temporal variability."
+ Phe observations have been carried out at the Nordic Optical Telescope. La Palma from. 1998-2000 and consist mainly of high time resolution white light CCD photometry time-series with a tvpical duration of 3-4 hours per epoch (see detailed observing log in Table 1).," The observations have been carried out at the Nordic Optical Telescope, La Palma from 1998-2009 and consist mainly of high time resolution white light CCD photometry time-series with a typical duration of 3-4 hours per epoch (see detailed observing log in Table 1)."
+ Our data are taken in white light using two cillerent CCD detectors making their light curves not strictly comparable with each other due to the different QL properties., Our data are taken in white light using two different CCD detectors making their light curves not strictly comparable with each other due to the different QE properties.
+ However. both the CCD’s used are blue sensitive.," However, both the CCD's used are blue sensitive."
+" Also. only cillerential photometry was obtained cach time. since we were mainly interested in looking for periodicities and also because there is no meaningful zero point calibration for white light ""instrumental magnitudes."," Also, only differential photometry was obtained each time, since we were mainly interested in looking for periodicities and also because there is no meaningful zero point calibration for white light “instrumental magnitudes""."
+ In order to achieve fast. readout the CCD images were sub-windowed and binned by a factor of 2., In order to achieve fast readout the CCD images were sub-windowed and binned by a factor of 2.
+ The images have been. bias-subtracted ancl Dat-fielded in the normal manner., The images have been bias-subtracted and flat-fielded in the normal manner.
+ Light curves were extracted: using standard aperture photometry and a nearby star located 227 WSW of the target was used. as à comparison for the differential photometry.," Light curves were extracted using standard aperture photometry and a nearby star located 22"" WSW of the target was used as a comparison for the differential photometry."
+ The resulting light curves show strong and variable IHickering and are plotted in Fig 1., The resulting light curves show strong and variable flickering and are plotted in Fig 1.
+ In addition to the optical data. we use the RATE ASAI X-ray light curves and we have also revisited. the original EXOSAT X-ray. data published. in Machin ct al. (," In addition to the optical data, we use the RXTE ASM X-ray light curves and we have also revisited the original EXOSAT X-ray data published in Machin et al. ("
+1990).,1990).
+ EXNOSAT MIS. light curve data products. were extracted clirectly from the NASA-GSEC LEASARC database. and used. as such.," EXOSAT ME light curve data products were extracted directly from the NASA-GSFC HEASARC database, and used as such."
+ The time series analysis of our data has been performed in two stages., The time series analysis of our data has been performed in two stages.
+ Firstly. we have obtained Lomb-Scarele (Scargle 1982) power spectra of all the individual time series searching for any periodicities.," Firstly, we have obtained Lomb-Scargle (Scargle 1982) power spectra of all the individual time series searching for any periodicities."
+ The time resolution of our light curves varies. but it is typically of the order of 10-20 seconds.," The time resolution of our light curves varies, but it is typically of the order of 10-20 seconds."
+ “Phus our analysis is sensitive to periods down to about 40 sec., Thus our analysis is sensitive to periods down to about 40 sec.
+ In the second phase of our analysis we have modelled the variability in each of the light curves with a stochastic time series model ancl used those as a basis for a Monte Carlo studs ofthe noise properties of the light curves., In the second phase of our analysis we have modelled the variability in each of the light curves with a stochastic time series model and used those as a basis for a Monte Carlo study of the noise properties of the light curves.
+ This enabled. us to obtain confidence limits on any peaks in the power spectra. even in the presence of red. (correlated) noise that is often present in accreting sources and evidenced by non-periodic flickering behaviour.," This enabled us to obtain confidence limits on any peaks in the power spectra, even in the presence of red (correlated) noise that is often present in accreting sources and evidenced by non-periodic flickering behaviour."
+ The standard false. alarm: probability estimate. from. the Lomb-Scarele algorithm gives the significance of the highest peak in the power spectrum assuming that all the datapoints are independent., The standard false alarm probability estimate from the Lomb-Scargle algorithm gives the significance of the highest peak in the power spectrum assuming that all the datapoints are independent.
+ However. in the presence of correlated data (i.c. red noise) we will have to take a dillerent. approach in order to properly estiniate the significance of the peaks evident in the various power spectra.," However, in the presence of correlated data (i.e. red noise) we will have to take a different approach in order to properly estimate the significance of the peaks evident in the various power spectra."
+ We have done this numerically by means of Monte. Carlo simulations., We have done this numerically by means of Monte Carlo simulations.
+ Our approach consists of estimating the red noise component in the time domain using an autoregressive time series moclel as à basis for our analysis., Our approach consists of estimating the red noise component in the time domain using an autoregressive time series model as a basis for our analysis.
+ Autoregressive (AR) mocdels (sce Chatlicld 1989 for time domain time series analysis) of different order can be fitted to equally spaced time series in order to estimate the red noise content., Autoregressive (AR) models (see Chatfield 1989 for time domain time series analysis) of different order can be fitted to equally spaced time series in order to estimate the red noise content.
+ Once the correct AR model is chosen and fitted to the data. it can be used," Once the correct AR model is chosen and fitted to the data, it can be used"
+"Assuming the steady formation of sources over time period T, the luminosity distribution can be written in the form (Note that the first factor 7/T would not appear for a LF constructed from a sample of sources produced in an instantaneous star formation burst; however, this is the case neither for the galactic bulge nor for elliptical galaxies we consider here.)","Assuming the steady formation of sources over time period $T$, the luminosity distribution can be written in the form (Note that the first factor $\tau/T$ would not appear for a LF constructed from a sample of sources produced in an instantaneous star formation burst; however, this is the case neither for the galactic bulge nor for elliptical galaxies we consider here.)"
+" Clearly, the shape of LF will be determined by the dependence of all three factors in this expression on the X-ray luminosity."," Clearly, the shape of LF will be determined by the dependence of all three factors in this expression on the X-ray luminosity."
+ Consider these factors separately for bright (with giant donors) and dim (with main-sequence donors) LMXBs., Consider these factors separately for bright (with giant donors) and dim (with main-sequence donors) LMXBs.
+" At luminosities below Lxz10°” ((M«10-? Mo/year), the majority of donor stars in our LMXB sample are main sequence stars with long life times (>Gyr)."," At luminosities below $L_{\rm x} \approx 10^{37}$ $\dot{M}<10^{-9}$ $M_\odot$ /year), the majority of donor stars in our LMXB sample are main sequence stars with long life times $>$ Gyr)."
+" The stable mass transfer episode is also very long, so the factor 7/T should not strongly depend on Lx."," The stable mass transfer episode is also very long, so the factor $\tau/T$ should not strongly depend on $L_{\rm x}$."
+" The period distribution n(P) is generally determined by the evolutionary history of binaries which includes as a minimum a supernova explosion to produce the neutron star, a common envelope phase, etc. ("," The period distribution $n(P)$ is generally determined by the evolutionary history of binaries which includes as a minimum a supernova explosion to produce the neutron star, a common envelope phase, etc. ("
+"or involves dynamical processes in dense stellar clusters), and can be found for example from population synthesis studies.","or involves dynamical processes in dense stellar clusters), and can be found for example from population synthesis studies."
+ Let us put it in the form dN/dlogP and leave it as it is for a while., Let us put it in the form $dN/d\log P$ and leave it as it is for a while.
+" The luminosity of a LMXB due to magnetic stellar wind braking depends on the orbital period as a power law L,cP?5 (seee.g.ben&Tutukov|1984) or P? (Patterson|[1984),, or LxcP*? (Cote&Pylyse1989) so written in the form dlogP/dlogL, this factor isindependent of L,."," The luminosity of a LMXB due to magnetic stellar wind braking depends on the orbital period as a power law $L_{\rm x}\propto P^{2.5}$ \citep[see e.g.][]{iben84} or $L_{\rm x}\propto P^{3.3}$ \citep{patterson84}, or $L_{\rm x}\propto P^{3.8}$ \citep{cote89} so written in the form $d\log P/d\log L_x$ this factor isindependent of $L_x$."
+" So we find This would match the observed shape of the LF of dim LMXBs if the factor dN/dlogP~ const(L,).", So we find This would match the observed shape of the LF of dim LMXBs if the factor $dN/d\log P\sim$ $L_x$ ).
+" Unfortunately, it is difficult to construct this distribution from observations due to many selection effects."," Unfortunately, it is difficult to construct this distribution from observations due to many selection effects."
+" Population synthesis calculations (e.g. produce a variable distribution, which only roughly can be considered as constant."," Population synthesis calculations \cite[e.g.][]{fragos08} produce a variable distribution, which only roughly can be considered as constant."
+" Observations suggest that initial orbital periods of binary stars do follow this dependence, ie. dN/dlogP =const and we shall assume this to hold approximately at all [1982)..,stages of binary evolution."," Observations suggest that initial orbital periods of binary stars do follow this dependence, i.e. $dN/d\log P=$ const \citep{popova82}, , and we shall assume this to hold approximately at all stages of binary evolution."
+ LMXBs with luminosities higher than L4>10°75 hhave predominantly giants companions (see above)., LMXBs with luminosities higher than $L_{\rm x}>10^{37.5}$ have predominantly giants companions (see above).
+" The duration 7 of the mass transfer (and hence the accretion stage) in these systems is significantly smaller than that for systems with main sequence donors, and to a large extent is determined by the mass transfer rate (seee.g.ebbink,Rappaport,& Savonije|[1983)."," The duration $\tau$ of the mass transfer (and hence the accretion stage) in these systems is significantly smaller than that for systems with main sequence donors, and to a large extent is determined by the mass transfer rate \cite[see e.g.][]{webbink83}."
+". Let us assume that the luminosity of LMXB is a powerlaw function of its orbital period L4οςP (which is true for the magnetic stellar wind braking dominated regime, a~2.5—3.8, see above, and also holds when nuclear evolution of the giant donor[Webbink, is responsibleRappaport, for its &RocheSavonijd lobe[[983]). overflow, a~1, see "," Let us assume that the luminosity of LMXB is a powerlaw function of its orbital period $L_{\rm x}\propto P^{\alpha}$ (which is true for the magnetic stellar wind braking dominated regime, $\alpha\sim2.5-3.8$, see above, and also holds when nuclear evolution of the giant donor is responsible for its Roche lobe overflow, $\alpha\sim1$, see \citealt{webbink83}) )."
+We also adopted that the duration of the bright LMXB stage is inversely proportional to its luminosity τoL;!., We also adopted that the duration of the bright LMXB stage is inversely proportional to its luminosity $\tau \propto L_{\rm x}^{-1}$.
+" This is approximately correct because the total mass, which might be accreted from a late-type giant donor stars does not vary much from system to system (seee.g.Rappaport,&Savonije|1983)."," This is approximately correct because the total mass, which might be accreted from a late-type giant donor stars does not vary much from system to system \cite[see e.g.][]{webbink83}."
+. Combining these we obtain Thusone Fabbianof2004).," Combining these we obtain Thus \citep{gilfanov04,kim04}."
+. The condition for this to be correct[Kim is& that the donor stars in these LMXBs are short-living giants., The condition for this to be correct is that the donor stars in these LMXBs are short-living giants.
+" We can try to estimate the longest lifetimes of such giants assuming that the maximum mass, which might be accreted from them is $0.6Mo (αμ. and their duty cycle is close to unity."," We can try to estimate the longest lifetimes of such giants assuming that the maximum mass, which might be accreted from them is $\la 0.6 M_{\odot}$ \cite[e.g.][]{webbink83} and their duty cycle is close to unity."
+ In this case TgiantsS0.6M/(2x1079Ms /year)~3x105 years., In this case $\tau_{\rm giants}\la 0.6M_\odot/(2\times10^{-9}M_\odot/$ $)\sim3\times10^{8}$ years.
+ This means that the list of bright LMXBs in the Galaxy should alter after approximately 30 Myrs., This means that the list of bright LMXBs in the Galaxy should alter after approximately 30 Myrs.
+" We would like to note here that the boundary between different types of donors at luminosities above and below Lx~10°” erg/sec was also previously noted in numerical simulations of LMXB populations in |Fragosetal.|(2008); (2009),, but the slopes of the LMXB LF above and below this luminosity was found to be not so different."," We would like to note here that the boundary between different types of donors at luminosities above and below $L_{\rm x}\sim10^{37}$ erg/sec was also previously noted in numerical simulations of LMXB populations in \cite{fragos08,kim09}, but the slopes of the LMXB LF above and below this luminosity was found to be not so different."
+" The inability of particular numerical simulations to reproduce the observed break in the X-ray luminosity function of the LMXB population might indicate that we do not (yet) understand all details of the physical processes (e.g. the mode of the angular momentum loss, the common envelope phase, etc.)"," The inability of particular numerical simulations to reproduce the observed break in the X-ray luminosity function of the LMXB population might indicate that we do not (yet) understand all details of the physical processes (e.g. the mode of the angular momentum loss, the common envelope phase, etc.)"
+ that shape the LMXB LF., that shape the LMXB LF.
+ We plan to perform more detailed calculations of the LMXB population in our future work. Acknowledgments.., We plan to perform more detailed calculations of the LMXB population in our future work. .
+" This work is partiallysupported by RFBR grants 10-02-00492-a, 10-02-00599-a, by Russian Federation President programm NSh-5069.2010.2, Federal Russian Science Agency through the research contract 02.740.11.0575, and program of Presidium RAS P-19 and OFN-16."," This work is partiallysupported by RFBR grants 10-02-00492-a, 10-02-00599-a, by Russian Federation President programm NSh-5069.2010.2, Federal Russian Science Agency through the research contract 02.740.11.0575, and program of Presidium RAS P-19 and OFN-16."
+the accretion disc) when the system is in a high optical state and is a few percent longer than the orbital period.,the accretion disc) when the system is in a high optical state and is a few percent longer than the orbital period.
+ On longer timescales. both VS03 Cen (Patterson et al.," On longer timescales, both V803 Cen (Patterson et al."
+ 2000) and CR Boo (Patterson et al., 2000) and CR Boo (Patterson et al.
+ 1997) have shown a prominent modulation (—1 mag) on a period of ~1922 hrs at. various epochs: this has been called. the ceveling state’., 1997) have shown a prominent modulation $\sim$ 1 mag) on a period of $\sim$ 19–22 hrs at various epochs: this has been called the `cycling state'.
+ Patterson ct al. (, Patterson et al. (
+2000) noted that the amplitude ancl period. of these modulations placed. these svstenis on the short. period. end. of the WKukearkin-Parenago period-amplitude relationship of normal dwarf novac outbursts (Warner 1987).,2000) noted that the amplitude and period of these modulations placed these systems on the short period end of the Kukarkin-Parenago period-amplitude relationship of normal dwarf novae outbursts (Warner 1987).
+ They concluded that these near day-long modulations were similar to hwdrogen-rich dwarl novae outbursts (in particular the short-period LR. UMa svstems), They concluded that these near day-long modulations were similar to hydrogen-rich dwarf novae outbursts (in particular the short-period ER UMa systems).
+ On even longer timescales. evidence has been found. for a period of 77 davs in VSO3 Cen (Ixato ct al.," On even longer timescales, evidence has been found for a period of 77 days in V803 Cen (Kato et al."
+ 2000a) and a period of 46 davs in Cl Boo (Ixato et al., 2000a) and a period of 46 days in CR Boo (Kato et al.
+ 20005)., 2000b).
+ Both binaries spend around half the time in a bright state., Both binaries spend around half the time in a bright state.
+ Whilst our observations of KL Dra are less extensive than those of VSO3. Cen and CR Doo. they show characteristics which are clearly dillerent to these two AM CVn systems.," Whilst our observations of KL Dra are less extensive than those of V803 Cen and CR Boo, they show characteristics which are clearly different to these two AM CVn systems."
+ Indeed. with a possible outburst evele of 60 davs and with a cuty-evele of ~1/5. WL Dra closely resembles the outbursts seen in the SU UMa or U Gom type dwarf novac.," Indeed, with a possible outburst cycle of $\sim$ 60 days and with a duty-cycle of $\sim$ 1/5, KL Dra closely resembles the outbursts seen in the SU UMa or U Gem type dwarf novae."
+ However. at this stage we cannot exclude that the outbursts seen in KL Dra are not the AM CVn equivalent of WZ Sec outhursts (dwarf novae which recur on a timescales of vears and have outburst amplitudes of 69 mag. eg Patterson ct al.," However, at this stage we cannot exclude that the outbursts seen in KL Dra are not the AM CVn equivalent of WZ Sge outbursts (dwarf novae which recur on a timescales of years and have outburst amplitudes of 6–9 mag, eg Patterson et al."
+ 2002)., 2002).
+ As such it makes KL Dra a prime target to investigate the similarities and dillerences. between hyvdrogen-rich. and. hyvdrogen-delficient accretion Iows., As such it makes KL Dra a prime target to investigate the similarities and differences between hydrogen-rich and hydrogen-deficient accretion flows.
+ The hyerogen-acereting dwarf novae have been much studied at. various. wavelengths. over several decades. and show regular outbursts on timescales ranging from a few weeks to months or even vears., The hydrogen-accreting dwarf novae have been much studied at various wavelengths over several decades and show regular outbursts on timescales ranging from a few weeks to months or even years.
+ Many systems show characteristics similar to SS (νο which showed in one outburst that the increase in the extreme UV emission was delayed by 1.52.0 cays with respect to the optical emission (Wheatley. Alauche Mattei 2003).," Many systems show characteristics similar to SS Cyg which showed in one outburst that the increase in the extreme UV emission was delayed by 1.5–2.0 days with respect to the optical emission (Wheatley, Mauche Mattei 2003)."
+ In contrast. the hard X-rays. after an initial burst. were suppressed during the remainder of the outburst.," In contrast, the hard X-rays, after an initial burst, were suppressed during the remainder of the outburst."
+ To complicate matters. U Com has shown at least one outburst where the hard X-rays follow the optical and UV tux (Mattei. Mauche Wheatley 2000).," To complicate matters, U Gem has shown at least one outburst where the hard X-rays follow the optical and UV flux (Mattei, Mauche Wheatley 2000)."
+ lt is not clear if this is due to the relatively high. binary inclination of U Gem., It is not clear if this is due to the relatively high binary inclination of U Gem.
+ More recently CAV Lib (which has only had two known outbursts) also shows the enhancement of X-rays at the start of the outburst (Byekline et al., More recently GW Lib (which has only had two known outbursts) also shows the enhancement of X-rays at the start of the outburst (Byckling et al.
+ 2009)., 2009).
+ One of the kev reasons for obtaining observations was ο determine how the X-ray (us of KL Dra varied over the course of the outburst evele., One of the key reasons for obtaining observations was to determine how the X-ray flux of KL Dra varied over the course of the outburst cycle.
+ I is clear that X-rays are detected. from KWL Dra both during the low and the high optical states., It is clear that X-rays are detected from KL Dra both during the low and the high optical states.
+ However. there is no evidence that the X-ray Ilux changed significantly curing the course of the outburst (ligure 3)).," However, there is no evidence that the X-ray flux changed significantly during the course of the outburst (Figure \ref{swift-light}) )."
+ Our observations of IKL Dra show that it undergoes frequent optical outbursts which are also seen at UV wavelengths., Our observations of KL Dra show that it undergoes frequent optical outbursts which are also seen at UV wavelengths.
+ Although our coverage is by no means complete. our data are consistent with WL Dra undergoing outbursts on a period of ~60 days.," Although our coverage is by no means complete, our data are consistent with KL Dra undergoing outbursts on a period of $\sim$ 60 days."
+ The amplitude of the outbursts (3. mag) and duration of the outburst (~2 weeks) are very similar to the typical outbursts that seen in the SU UMa or U Cem sub-class of (hydrogen) accreting chwarl novae., The amplitude of the outbursts $\sim$ 3 mag) and duration of the outburst $\sim$ 2 weeks) are very similar to the typical outbursts that seen in the SU UMa or U Gem sub-class of (hydrogen) accreting dwarf novae.
+ As such we encourage other observers. whether they are amateur astronomers who have suitable equipment or team members of survey telescopes which have suitable spatial resolution (to resolve WL Dra from its “nearby galaxy). to obtain a more complete long term coverage and hence verily the 60 day outburst period.," As such we encourage other observers, whether they are amateur astronomers who have suitable equipment or team members of survey telescopes which have suitable spatial resolution (to resolve KL Dra from its `nearby' galaxy), to obtain a more complete long term coverage and hence verify the $\sim$ 60 day outburst period."
+ The fact that KL Dra shows such regular outbursts makes it an ideal svstenm with which to investigate helium accretion Hows in detail., The fact that KL Dra shows such regular outbursts makes it an ideal system with which to investigate helium accretion flows in detail.
+ The Liverpool Telescope is operated. on the island. of La Palma by Liverpool John Moores University with financial support from the Ulx Science and. Technology. Facilities Council., The Liverpool Telescope is operated on the island of La Palma by Liverpool John Moores University with financial support from the UK Science and Technology Facilities Council.
+ We thank the LP Support Astronomer Chris Moss for his assistance in scheduling our observations and the PL Neil Gehrels. together with the science and operations teams for their support. of these observations.," We thank the LT Support Astronomer Chris Moss for his assistance in scheduling our observations and the PI, Neil Gehrels, together with the science and operations teams for their support of these observations."
+ We also thank Pasi Llakala ancl Tina Limets for obtaining several images using the Nordic Optical Telescope which is operated. on the island. of La Palma jointlv by Denmark. Finland. Iceland. Norway. and Sweden.," We also thank Pasi Hakala and Tiina Liimets for obtaining several images using the Nordic Optical Telescope which is operated on the island of La Palma jointly by Denmark, Finland, Iceland, Norway, and Sweden."
+ The Isaac Newton and William Lerschel Telescopes are operated on the island of La Palma by the Isaac Newton Group., The Isaac Newton and William Herschel Telescopes are operated on the island of La Palma by the Isaac Newton Group.
+ All three telescopes are located in the Spanish Observatorio del Roque de los Aluchachos of the Lnstituto cle Astrofisica ce Canarias., All three telescopes are located in the Spanish Observatorio del Roque de los Muchachos of the Instituto de Astrofisica de Canarias.
+ Observations were also obtained at the Gemini Observatory. which is operated. by the Association of Universities Lor Research in Astronomy. Inc.. under a cooperative agreement with the NSE on behalf of the Ciemini partnership: the National Science Foundation. (United States). the Science," Observations were also 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 Science"
+The 2ALASS dataset used in this paper is described. brielly below.,The 2MASS dataset used in this paper is described briefly below.
+ This material is common to a companion paper about attempted. detection of the ISW elfect via cross-correlation (?) and Culler details are given there., This material is common to a companion paper about attempted detection of the ISW effect via cross-correlation \citep{Francis_ISW} and fuller details are given there.
+" PAIASS is an allskv survey in the J. // and A, bands (?).."," 2MASS is an all-sky survey in the $J$ , $H$ and $K_{s}$ bands \citep{Jarrett_2MASS}."
+ This near-infrarect selection means that 24LASS is sensitive to old. stars and. hence to the most massive structures. which makes it well-suited for constructing maps of galaxy mass density.," This near-infrared selection means that 2MASS is sensitive to old stars and hence to the most massive structures, which makes it well-suited for constructing maps of galaxy mass density."
+ Vhe final extended source catalogue (Νο) contains over 1.6 million objects. over 984 of which are galaxies.," The final extended source catalogue (XSC) contains over 1.6 million objects, over $98\%$ of which are galaxies."
+ Photometric redshifts for the 2ALASS NSC have been generated by matching the 2MLASS data with optical catalogues. from SuperCOSALOS scans of major photographic sky surveys (UINST. in the south: POSS2 in the north) to ereate a S-band dataset., Photometric redshifts for the 2MASS XSC have been generated by matching the 2MASS data with optical catalogues from SuperCOSMOS scans of major photographic sky surveys (UKST in the south; POSS2 in the north) to create a 5-band dataset.
+ Details of the SuperCOSALOS catalogue construction process are given in ? and details of the photometric calibration and photo-z determination method are given in ?.., Details of the SuperCOSMOS catalogue construction process are given in \citet{Hambly_SuperCOSMOS} and details of the photometric calibration and $z$ determination method are given in \citet{Peacock_photoz}.
+ Phe final photometric redshift dataset has an overall ris in. tyne+ OF 0.033., The final photometric redshift dataset has an overall rms in $z_{\mathrm{phot}}-z$ of 0.033.
+ A magnitude cut νο<13.8 ds imposed on the extinction-corrected data to ensure. uniformity across the region outside the Galactic plane., A magnitude cut $K_{s}<13.8$ is imposed on the extinction-corrected data to ensure uniformity across the region outside the Galactic plane.
+ The Galactic plane mask is constructed. using the dust maps of ane exclucles all regions with A-band reddening sly2? 0.05. “Phis leaves approximately 674 of the sky. unmasked., The Galactic plane mask is constructed using the dust maps of \citet{Schlegel_dust} and excludes all regions with $K$ -band reddening $A_{K}>0.05$ This leaves approximately $67\%$ of the sky unmasked.
+ Whilst the choice ol Alyc0.05 is conservative. ? note that the number density within pixels of faint galaxies (18.5<foo 14.0) drops olf as a function. of cli for elas0.065.," Whilst the choice of $A_{K}>0.05$ is conservative, \citet{Afshordi} note that the number density within pixels of faint galaxies $13.5<K_{20}<14.0$ ) drops off as a function of $A_{K}$ for $A_{K}>0.065$."
+ X less conservative mask. for example with elyσε0.1. would. still only leave 79% of the sky unmasked at the price of possible incompleteness issues in the analysis.," A less conservative mask, for example with $A_{K}>0.1$, would still only leave $79\%$ of the sky unmasked – at the price of possible incompleteness issues in the analysis."
+ Lt is possible to attempt a full 3D reconstruction of the loca density field with this photometric redshift dataset. and then use this to estimate the local ISW signal (see 2)).," It is possible to attempt a full 3D reconstruction of the local density field with this photometric redshift dataset, and then use this to estimate the local ISW signal (see \citealt{Francis_PHD}) )."
+ However. this recovered. density field is subject to significant. radia smearing due to the uncertainty in the radial position of cach galaxy (approximately 100. Mpc).," However, this recovered density field is subject to significant radial smearing due to the uncertainty in the radial position of each galaxy (approximately $100\mpcoh$ )."
+ The extent of this smearing justifies a simpler and more transparent approach: to recover the projected. density feld. for 3 thick redshif slices of width As=0.1 to suas=0.3. ancl estimate the ISW signal due to each slice.," The extent of this smearing justifies a simpler and more transparent approach: to recover the projected density field for 3 thick redshift slices of width $\Delta z=0.1$ to $z_{\mathrm{max}}=0.3$, and estimate the ISW signal due to each slice."
+ The galaxy distribution in cach photometric recdshif shell is known outside the masked region., The galaxy distribution in each photometric redshift shell is known outside the masked region.
+ We use a Wiener filtering technique to recover the galaxy density field in each shell., We use a Wiener filtering technique to recover the galaxy density field in each shell.
+ When the Wiener filter is applied to all-sky data. i recovers the optimal density field in the least squares sense: a thorough cliscussion of Wiener-lilter density reconstruction is presented in e.g. ?..," When the Wiener filter is applied to all-sky data, it recovers the optimal density field in the least squares sense; a thorough discussion of Wiener-filter density reconstruction is presented in e.g. \citet{Zaroubi_1995}."
+ Jecause our dataset is masked at low Galactic latitudes. the simple all-skv Wiener approach must be mocified.," Because our dataset is masked at low Galactic latitudes, the simple all-sky Wiener approach must be modified."
+ We proceed by initially filling the masked region in cach of the redshift shells with a random Poisson sampling of galaxies according to the average number density observed. outside the plane., We proceed by initially filling the masked region in each of the redshift shells with a random Poisson sampling of galaxies according to the average number density observed outside the plane.
+ This galaxy distribution is then used to deduce the projected. galaxy density field for the redshift shell., This galaxy distribution is then used to deduce the projected galaxy density field for the redshift shell.
+ Phe Wiener filter appropriate for cach shell is then constructed as usual: where Cu.(f) is the angular power spectrum of galaxies in that redshift shell and Chou(6) describes shot noise., The Wiener filter appropriate for each shell is then constructed as usual: where $C_{\mathrm{gg}}(\ell)$ is the angular power spectrum of galaxies in that redshift shell and $C_{\mathrm{noise}}(\ell)$ describes shot noise.
+" The Wiener filter is applied. in the first instance to the τη, coellicients computed for the expansion of this initial full-skv galaxy density field into a basis of spherical harmonics.", The Wiener filter is applied in the first instance to the $a_{\ell m}$ coefficients computed for the expansion of this initial full-sky galaxy density field into a basis of spherical harmonics.
+ The resultant density. field within the masked. region. is Poisson sampled. for galaxies once more to vield a revised galaxy distribution composed of a sample from the Wiener reconstruction within the plane. and the original 2\LASS ealaxies in the unmasked regions.," The resultant density field within the masked region is Poisson sampled for galaxies once more to yield a revised galaxy distribution composed of a sample from the Wiener reconstruction within the plane, and the original 2MASS galaxies in the unmasked regions."
+ In this wav. large-scale features in surface density can be extrapolated across the plane in a consistent way.," In this way, large-scale features in surface density can be extrapolated across the plane in a consistent way."
+ “Phe process can be iterated. but this did not vield a significant improvement in the final results.," The process can be iterated, but this did not yield a significant improvement in the final results."
+ We have carried. out. tests of this. method on simulated data. and find satisfactory recovery of. the estimated ESW signal (Francis 2008).," We have carried out tests of this method on simulated data, and find satisfactory recovery of the estimated ISW signal (Francis 2008)."
+ Fig., Fig.
+ 1. shows the recovered. projected density field for each of the three shells.," \ref{fig:ISW_from_gals}
+ shows the recovered projected density field for each of the three shells."
+ We use HEALPix software (Gorski et al., We use HEALPix software (Gorski et al.
+ 2005: see also //healpix.jpl.nasa.gov)) to generate these maps and to perform harmonic analysis of them., 2005; see also ) to generate these maps and to perform harmonic analysis of them.
+" Given a projected. galaxy density. field. in a thick redshift shell. it is possible to estimate the CALB temperature Iuctuations. AY. due to the ISW οσοι,"," Given a projected galaxy density field in a thick redshift shell, it is possible to estimate the CMB temperature fluctuations, $\Delta T$, due to the ISW effect."
+ Poisson's equation tells us the time variation of the potential field in terms of the density growth factor in the linear regime. g(a): llence. using lPoisson's equation. we can integrate through a shell of thickness Ar: where 9 is the projected galaxy density Ποιά in the redshift shell under consideration and e=(1|2)! with > the redshift at the midpoint of the shell.," Poisson's equation tells us the time variation of the potential field in terms of the density growth factor in the linear regime, $g(a)$: Hence, using Poisson's equation, we can integrate through a shell of thickness $\Delta r$: where $\delta$ is the projected galaxy density field in the redshift shell under consideration and $\bar{a} = (1+\bar{z})^{-1}$ with $\bar{z}$ the redshift at the midpoint of the shell."
+" Phe bias in cach redshift shell between the galaxy. ancl matter density fields is approximated. using a linear bias relation 6,=ὑδι, and is assumed. to be independent of scale ancl redshift in each shell.", The bias in each redshift shell between the galaxy and matter density fields is approximated using a linear bias relation $\delta_{g}=b\delta_{m}$ and is assumed to be independent of scale and redshift in each shell.
+ Fo deduce the appropriate value of 6 in cach redshift shell. a maximum likelihood. fit of the predicted: galaxy angular power spectrum to the measured galaxy. angular power spectrum is performed as in ?..," To deduce the appropriate value of $b$ in each redshift shell, a maximum likelihood fit of the predicted galaxy angular power spectrum to the measured galaxy angular power spectrum is performed as in \citet{Francis_ISW}."
+ Fig., Fig.
+ 1 shows the results of estimating the local. I8W ellect from the Wiener filtered density field in redshift shells., \ref{fig:ISW_from_gals} shows the results of estimating the local ISW effect from the Wiener filtered density field in redshift shells.
+ These ISW maps are alfected by overlap effects due tothe uncertainty in the radial position of cach galaxy: galaxies associated with an overdensity near the boundary of a shell will to some extent be spread. across both the “true” shell and the adjacent redshift shell., These ISW maps are affected by overlap effects due tothe uncertainty in the radial position of each galaxy: galaxies associated with an overdensity near the boundary of a shell will to some extent be spread across both the `true' shell and the adjacent redshift shell.
+Fig.,Fig.
+ 2) shows the angular, \ref{fig:ISW_angpow} shows the angular
+The parameters used for the simulations are summarized in ‘Table 1..,The parameters used for the simulations are summarized in Table \ref{tab:mhd_sum}.
+ Phe Table columns (1)-(8) show respectively the number of a given run: the BIL radius and the Bondi radius ratio (Ro): the cireularization radius compared to the Bondi racius (Ao): the adiabatic index (5): the initial magnetic plasma parameter (53) at à: the final time (fpay) at which we stopped cach simulation in units of the Keplerian orbital time atr; (fin=5958 for Alpy=3.7.107M. the same final time fp but in units of the Ixeplerian orbital time ator. and finally the time-averaged mass accretion rate on to the BLL in units of the Bondi accretion rate CAL/A4) over LOOO νι at the end of cach run.," The Table columns (1)-(8) show respectively the number of a given run; the BH radius and the Bondi radius ratio $R'_{\rm S}$ ); the circularization radius compared to the Bondi radius $R'_{\rm C}$ ); the adiabatic index $\gamma$ ); the initial magnetic plasma parameter $\beta_{\rm in}$ ) at $r_{\rm o}$; the final time $t_{\rm f,in}$ ) at which we stopped each simulation in units of the Keplerian orbital time at $r_i$ $t_{\rm dyn} = 595 s$ for $M_{\rm BH}= 3.7\times 10^6\MSUN$ ); the same final time $t_{\rm f,out}$ but in units of the Keplerian orbital time at $r_{\rm o}$, and finally the time-averaged mass accretion rate on to the BH in units of the Bondi accretion rate $\MDOT_{\rm a}/\MDOT_{\rm B}$ ) over 1000 $t_{\rm dyn}$ at the end of each run."
+ The shorter cluration of ow simulations (compared. to the duration of PBOSh’s simulations) is caused by the high density gradients which appear in our runs for low 5. and which lead to significantly smaller time steps.," The shorter duration of our simulations (compared to the duration of PB03b's simulations) is caused by the high density gradients which appear in our runs for low $\gamma$, and which lead to significantly smaller time steps."
+ 1n our simulations. the evolution of the poloidal velocity field initially proceeds in a similar manner as in the corresponding LID simulations (see e.g. panels 3 and 4 in fig.2 in PALOS for each >).," In our simulations, the evolution of the poloidal velocity field initially proceeds in a similar manner as in the corresponding HD simulations (see e.g. panels 3 and 4 in fig.2 in PM08 for each $\gamma$ )."
+" Namely. after a short. episode of the racial inflow. the matter with relatively high-£ forms a torus in the equatorial plane around the BIL(high-£ is defined here as doc7, and low-fas |<ἐν. where ων=Diese is a critical specific angular momentum)."," Namely, after a short episode of the radial inflow, the matter with relatively forms a torus in the equatorial plane around the BH is defined here as $l > l_{cr}$ and as $l < l_{cr}$, where $l_{cr} = 2 R_{\rm S}c$ is a critical specific angular momentum)."
+ For 5? —5/3. 4/3. and 1.2. initially when the torus forms. a shock wave is proclucect.," For $\gamma$ =5/3, 4/3, and 1.2, initially when the torus forms, a shock wave is produced."
+ The shock wave forms because some of the gas forming a torus turns around and. moves outward along the equator due to the centrifugal force and gas pressure or tries to move closer to the poles anc accrete., The shock wave forms because some of the gas forming a torus turns around and moves outward along the equator due to the centrifugal force and gas pressure or tries to move closer to the poles and accrete.
+ For 7? —1.01. the gas is the most gravitationally bound. compared to other 5s. As such it accumulates near DII and does not Hlow out along the equator contrary to the higher + cases.," For $\gamma=$ 1.01, the gas is the most gravitationally bound compared to other $\gamma$ 's. As such it accumulates near BH and does not flow out along the equator contrary to the higher $\gamma$ cases."
+ In its early phase of the evolution. the torus is not able o accrete on to a DIL independently of 4. because the MIID urbulence has not vet fully developed. το transoport. the angular momentum.," In its early phase of the evolution, the torus is not able to accrete on to a BH, independently of $\gamma$, because the MHD turbulence has not yet fully developed to transoport the angular momentum."
+ The matter in the polar regions doces not feel any centrifugal barier. thus it Hows raclially toward he DII.," The matter in the polar regions does not feel any centrifugal barier, thus it flows radially toward the BH."
+ In our. MILD simulations and the HD. simulations stucied in PALOS. the sonic surfaces topology evolves in a similar manner.," In our MHD simulations and the HD simulations studied in PM08, the sonic surfaces topology evolves in a similar manner."
+ Lt changes from a circular shape into two obes where the lobes are elongated along the poles., It changes from a circular shape into two lobes where the lobes are elongated along the poles.
+ As we reported in. PMOS. the evolution of the sonic surface is the slowest for 7 —1.01. due to a lack on an equatorial outflow.," As we reported in PM08, the evolution of the sonic surface is the slowest for $\gamma$ =1.01, due to a lack on an equatorial outflow."
+ After à short time the magnetic field. effects start to dominate the aecretion Low evolution., After a short time the magnetic field effects start to dominate the accretion flow evolution.
+ The torus starts to accrete toward the BI., The torus starts to accrete toward the BH.
+ The biggest and expected dillerence between LID and. MILD. simulations is the development of bipolar outllows in. ALD simulations regardless of 5., The biggest and expected difference between HD and MHD simulations is the development of bipolar outflows in MHD simulations regardless of $\gamma$.
+ PDBOS3b's simulations well illustrate the general evolution of the magnetic fields., PB03b's simulations well illustrate the general evolution of the magnetic fields.
+ Ehe polar outllows form at small racii (r« PORS) close to the poles anc expand. to the outer boundary ofthe computational domain along the svmmetry axis., The polar outflows form at small radii $r < 20 R_{\rm S}$ ) close to the poles and expand to the outer boundary of the computational domain along the symmetry axis.
+ We find that the time at which the outllows forni depends on 5: the lower 5. the later the outflow forms.," We find that the time at which the outflows form depends on $\gamma$: the lower $\gamma$, the later the outflow forms."
+ Although initially LID and ATID Hows evolve similarly. at the end of the simulations. the geometry and character of the accretion Lows in the LID and MEIID. simulations. are significantly clilferent.," Although initially HD and MHD flows evolve similarly, at the end of the simulations, the geometry and character of the accretion flows in the HD and MHD simulations, are significantly different."
+ At the ened of the simulation. the AILLD flows are composed of a bipolar outllow propagating in the polar regions. and the inflow near the equator.," At the end of the simulation, the MHD flows are composed of a bipolar outflow propagating in the polar regions, and the inflow near the equator."
+ The equatorial region is turbulent and strongly asvmumetric [or all 5s. Fig., The equatorial region is turbulent and strongly asymmetric for all $\gamma$ 's. Fig.
+" 1. shows the AL,Aly as a function of time. for cach studied 5 as listed in Table 1.."," \ref{fig:accretion_evolution} shows the $\MDOT_{\rm a}/\MDOT_{\rm B}$ as a function of time, for each studied $\gamma$ as listed in Table \ref{tab:mhd_sum}."
+ Phe time is given in units of the orbital time at the inner boundary of the computational domain., The time is given in units of the orbital time at the inner boundary of the computational domain.
+ A pair arrows in cach panel mark the moments at which the polar outflow forms. in the Northern (N) and Southern (8) directions.," A pair arrows in each panel mark the moments at which the polar outflow forms, in the Northern (N) and Southern (S) directions."
+ The moment of the formation of the large seale polar outflow is defined as a moment when the matter stops accreting through the polar regions and the radial velocity sign changes sign from negative to positive., The moment of the formation of the large scale polar outflow is defined as a moment when the matter stops accreting through the polar regions and the radial velocity sign changes sign from negative to positive.
+ We determine this by inspecting the velocity maps and checking that the outllow velocity becomes larecr than escape velocity at the outer boundary., We determine this by inspecting the velocity maps and checking that the outflow velocity becomes larger than escape velocity at the outer boundary.
+ The upper-left. panel in Fig., The upper-left panel in Fig.
+ 1. shows M/Ma [or run 1l.," \ref{fig:accretion_evolution} shows $\dot{M}_{\rm
+a}/\dot{M}_{\rm B}$ for run 1."
+" Initially. Al,Ah, decreases [rom 1 to 0.8 and then reniains constant until {τε 2500."," Initially, $\MDOT_{\rm a}/\MDOT_{\rm B}$ decreases from 1 to 0.3 and then remains constant until $t \approx $ 2500."
+" The value for AZ,/Aly of 0.3 is consistent with that obtained in the LID counterpart (see ALPOS and. PDBO3b). ancl it corresponds to the matter accreting in the polar funels in the initial LED phase of the simulation."," The value for $\dot{M}_{\rm a} /\dot{M}_{\rm B}$ of 0.3 is consistent with that obtained in the HD counterpart (see MP08 and PB03b), and it corresponds to the matter accreting in the polar funels in the initial HD phase of the simulation."
+ However. later around t= 3700. AL/Mg suddenly decreases to a minimum value (below 0.001). ancl oscillates between the minimum valuc ancl 0.08.," However, later around t= 3700, $\dot{M}_{\rm a}/
+\dot{M}_{\rm B}$ suddenly decreases to a minimum value (below 0.001), and oscillates between the minimum value and 0.03."
+" We note that the moment of the S outllow formation does not correspond. exactly to the moment. of jo Al,Aly reduction to its minimum value.", We note that the moment of the S outflow formation does not correspond exactly to the moment of the $\MDOT_{\rm a}/\MDOT_{\rm B}$ reduction to its minimum value.
+ In general (the details will xe discussed in the next sections) the mass accretion rate reduction is caused by the mentioned bipolar outflow. which »ushes the raclially infalling material away from the polar unels.," In general (the details will be discussed in the next sections) the mass accretion rate reduction is caused by the mentioned bipolar outflow, which pushes the radially infalling material away from the polar funels."
+ Such a vacuuming of the polar regions makes the orus the only source of gas available for accretion., Such a vacuuming of the polar regions makes the torus the only source of gas available for accretion.
+" During ater phases of the evolution (i.e. />4000). AL,Aly shows narrow spikes and dips similar to those found. /bx. PBOShb. corresponding to the episodes of enhanced ancl suppressed orus accretion. respectively."," During later phases of the evolution (i.e., $t > 4000$ ), $\MDOT_{\rm a}/\MDOT_{\rm B}$ shows narrow spikes and dips similar to those found by PB03b, corresponding to the episodes of enhanced and suppressed torus accretion, respectively."
+ The upper-right panel in Fie., The upper-right panel in Fig.
+" 1 shows the Ad,/My evolution for 5 —4/3 (run 2).", \ref{fig:accretion_evolution} shows the $\MDOT_{\rm a}/\MDOT_{\rm B}$ evolution for $\gamma$ =4/3 (run 2).
+" At the beginning but after the orus formation. Al,/Adp_zz 0.2. which is consistent with miss accretion rate found previously in the LED. simulations."," At the beginning but after the torus formation, $\dot{M}_{\rm a} / \dot{M}_{\rm
+B}\approx$ 0.2, which is consistent with mass accretion rate found previously in the HD simulations."
+ At this stage. the gas is being aceretecl from the polar regions.," At this stage, the gas is being accreted from the polar regions."
+" Later when an outllow forms at /2 4000. AL,/My suddenly decreases."," Later when an outflow forms at $t\approx 4000$ , $\MDOT_{\rm
+a}/\MDOT_{\rm B}$ suddenly decreases."
+" For f74000. the AL,/My evolution is characterized. by narrow spikes. dips (similarly to run 1). ancl additionally broader features with higher amplitudes."," For $t >4000$, the $\MDOT_{\rm
+a}/\MDOT_{\rm B}$ evolution is characterized by narrow spikes, dips (similarly to run 1), and additionally broader features with higher amplitudes."
+ The mass accretion rate in the spikes can reach a level of about 0.02 and corresponds to the torus aceretion., The mass accretion rate in the spikes can reach a level of about 0.02 and corresponds to the torus accretion.
+" During he ips. the A,Mg bas a minimum value below 0.001. when the torus stops accreting."," During the dips, the $\dot{M}_{\rm a}/ \dot{M}_{\rm B}$ has a minimum value below 0.001, when the torus stops accreting."
+ The rates seen in the two xoad features (around (25500 and (7000) reach the level ol ~ 0.1., The rates seen in the two broad features (around t=5500 and t=7000) reach the level of $\sim$ 0.1.
+" At even later times (/> 9000). AL,Aly is almost constant around 0.01."," At even later times $t > 9000$ ), $\MDOT_{\rm a}/\MDOT_{\rm B}$ is almost constant around 0.01."
+ Phe broad features have been found xloreby PBO3b. ancl correspond to the non-rotating gas accretion that reenter the BLL vicinity in the equatorial glance.," The broad features have been found beforeby PB03b, and correspond to the non-rotating gas accretion that reenter the BH vicinity in the equatorial plane."
+ Near the enc of the simulation. the minima (dips) reappear.," Near the end of the simulation, the minima (dips) reappear."
+ For run 3(the lower-left panel in Fig. 1)).," For run 3(the lower-left panel in Fig. \ref{fig:accretion_evolution}) ),"
+ the evolution, the evolution
+Assume a quark star is born with a temperature 7/2T; (T. is the critical temperature below which superconcductivity sets in). ancl enters a superconducting-superíIuid phase in the core as it cools by neutrino emission (Ouved et al.,"Assume a quark star is born with a temperature $T > T_{\rm c}$ $T_{\rm c}$ is the critical temperature below which superconductivity sets in), and enters a superconducting-superfluid phase in the core as it cools by neutrino emission (Ouyed et al."
+ 2002: Nerannen οἱ al., 2002; Kerännen et al.
+ 2005). and contracts due to spin-down.," 2005), and contracts due to spin-down."
+" The front quickly expands to the entire star followed by the formation of rotationally induced vortices. analogous to rotating superfIuid ""Ile (the vortex lines are parallel to the rotation axis: Tillev&TTillev 1990)."," The front quickly expands to the entire star followed by the formation of rotationally induced vortices, analogous to rotating superfluid $^3$ He (the vortex lines are parallel to the rotation axis; Tilley 1990)."
+ Via the Meissner elfect. the magnetic field is partially screened [rom the regions outside the vortex cores.," Via the Meissner effect, the magnetic field is partially screened from the regions outside the vortex cores."
+ Now the svstem will consist ol. alternating regions of superconducting material with a screened magnetic field. ancl the vortices where most of the magnetic field resides.," Now the system will consist of, alternating regions of superconducting material with a screened magnetic field, and the vortices where most of the magnetic field resides."
+ As discussed in Quved et al. (, As discussed in Ouyed et al. (
+2004). this has interesting consequences on how the surface magnetic [ield adjusts to the interior field which is confined in the vortices.,"2004), this has interesting consequences on how the surface magnetic field adjusts to the interior field which is confined in the vortices."
+ In Quved et al. (, In Ouyed et al. (
+2005) we performed nunerical simulations of the alignment. of a quark stars exterior field. and. found (hat the plivsies involved was indicative of SGRactivity.,"2005) we performed numerical simulations of the alignment of a quark star's exterior field, and, found that the physics involved was indicative of SGR."
+ It has been shown that pure CFL matter is rigorously electrically neutral despite the unequal quark masses (Rajagopal Wilezek 2001)., It has been shown that pure CFL matter is rigorously electrically neutral despite the unequal quark masses (Rajagopal Wilczek 2001).
+ ILowever other work (Usov 2004: and references therein) indicates a thin crust (auus=LO 7M.) is allowed around a quark star clue to surface depletion of strange quarks.," However other work (Usov 2004; and references therein) indicates a thin crust $M_{\rm crust, max} = 10^{-5}M_{\odot}$ ) is allowed around a quark star due to surface depletion of strange quarks."
+ In our model we have assumed no depletion ol strange quarks which implies a bare quark star., In our model we have assumed no depletion of strange quarks which implies a bare quark star.
+ Another simplicity of our model resides in the fact that we have a single superconducting fluid (the CFL phase)., Another simplicity of our model resides in the fact that we have a single superconducting fluid (the CFL phase).
+ In the case of neutron stars (e.g. IXonenkov Geppert 2000). one has to deal with the neutron superlIuid inducing (he neutron vortices parallel to the rotation axis and (he proton superconductor inducing the [Iuxoids (the magnetic field concentrated into euantized proton vortex lines) in a direction parallel to the magnetic field.," In the case of neutron stars (e.g. Konenkov Geppert 2000), one has to deal with the neutron superfluid inducing the neutron vortices parallel to the rotation axis and the proton superconductor inducing the fluxoids (the magnetic field concentrated into quantized proton vortex lines) in a direction parallel to the magnetic field."
+ In our case. having a single superconducting-superfhiud implies (hat [Iuxoids are contained inside the vortex cores.," In our case, having a single superconducting-superfluid implies that fluxoids are contained inside the vortex cores."
+ Thus. the forces at plav. are quite different for the CFL quark star than for a neutron star: 1) the drag force induced by electron scattering is non-existent in our model since no electrons are admitted in pure CFL matter: ii) in (he case of neutron stus. there exists a force on the vortices due to the variation of the neutron or proton superfhuid gaps (A) with density which can expel or trap the vortices (Hsu 1999).," Thus, the forces at play are quite different for the CFL quark star than for a neutron star: i) the drag force induced by electron scattering is non-existent in our model since no electrons are admitted in pure CFL matter; ii) in the case of neutron stars, there exists a force on the vortices due to the variation of the neutron or proton superfluid gaps $\Delta$ ) with density which can expel or trap the vortices (Hsu 1999)."
+ In quark matter. (he variation of the gap is not well constrained and it is common to assume the BCS relation AxV1—(T/T.)*.," In quark matter, the variation of the gap is not well constrained and it is common to assume the BCS relation $\Delta \propto \sqrt{1- (T/T_c)^2}$."
+ The nearly uniform temperature implies a nearly uniform gap inside the quark star. thus.," The nearly uniform temperature implies a nearly uniform gap inside the quark star, thus,"
+"e.g. in ?,, ? or ?..","e.g. in \cite{1981A&A....97..373S}, , \cite{2008IJNMF..56...37V} or \cite{2009arXiv0906.0774R}."
+ These studies suggest that it is the shape of the kernel function that causes clumping instability., These studies suggest that it is the shape of the kernel function that causes clumping instability.
+ Given thatrpSPH does not show any spurious velocities in the uniform density test given above while standard SPH develops clumps within a few sound crossing times already shows that it cannot be the shape of the kernel alone that is relevant here., Given that does not show any spurious velocities in the uniform density test given above while standard SPH develops clumps within a few sound crossing times already shows that it cannot be the shape of the kernel alone that is relevant here.
+ In the following tests we have looked for any sign of the clumping instability but have not found any evidence for it independent of the number of neighbors we used., In the following tests we have looked for any sign of the clumping instability but have not found any evidence for it independent of the number of neighbors we used.
+ This fact is certainly one of the factors that makes our formulation have radically reduced errors in general., This fact is certainly one of the factors that makes our formulation have radically reduced errors in general.
+ The clumping instability stems from particles pushing each other closer to other particles., The clumping instability stems from particles pushing each other closer to other particles.
+ With a smaller distance to the other particles the gradient of the kernel becomes smaller and in the next time step the particle gets pushed further away from its initial position., With a smaller distance to the other particles the gradient of the kernel becomes smaller and in the next time step the particle gets pushed further away from its initial position.
+ This way particles can pile up in the flat central part of the smoothing kernel., This way particles can pile up in the flat central part of the smoothing kernel.
+ ThatrpSPH does not show the clumping instability makes one hopeful that even higher order kernels could now be used to further improve the accuracy of the obtained solution., That does not show the clumping instability makes one hopeful that even higher order kernels could now be used to further improve the accuracy of the obtained solution.
+" An easy two dimensional setup uses anadiabatic index of y=1.4 in a unit square domain x€{0,1}, y€{0,1} with periodic boundary conditions."," An easy two dimensional setup uses anadiabatic index of $\gamma=1.4$ in a unit square domain $x\in \{0,1\}$, $y\in \{0,1\}$ with periodic boundary conditions."
+" Particles are initialised on an exactly square lattice with a density of p(x,y)=1 so that the initial density estimate from the SPH kernel gives a density estimate of unity to better than four parts in one thousand."," Particles are initialised on an exactly square lattice with a density of $\rho(x,y)=1$ so that the initial density estimate from the SPH kernel gives a density estimate of unity to better than four parts in one thousand."
+ We then add a sinusoidal velocity perturbations to this uniform distribution., We then add a sinusoidal velocity perturbations to this uniform distribution.
+ We set the pressure to Po=p/* to have a sound speed of unity., We set the pressure to $P_0=\rho/\gamma$ to have a sound speed of unity.
+ For the first tests here we only use 50? particles as there are no features to resolve., For the first tests here we only use $50^2$ particles as there are no features to resolve.
+ In all cases we evolve to time t=4., In all cases we evolve to time $t=4$.
+" We choose vz(y)=óvycos(2»xy) with óv,=1/2."," We choose $v_x(y) = \delta v_{\rm y}\cos(2\pi\, y)$ with $\delta
+v_{\rm y}=1/2$."
+ This shear flow setup gives an average kinetic energy of 1/8., This shear flow setup gives an average kinetic energy of $1/8$.
+" A detailed discussion of how SPH behaves on this test for different neighbor numbers, particle numbers and viscosity prescriptions is given in the Appendix."," A detailed discussion of how SPH behaves on this test for different neighbor numbers, particle numbers and viscosity prescriptions is given in the Appendix."
+" In summary, it does very poorly and transfers kinetic energy into heat very rapidly loosing tens of percent in only four sound crossing times (two crossing times of the fastest particles)."," In summary, it does very poorly and transfers kinetic energy into heat very rapidly loosing tens of percent in only four sound crossing times (two crossing times of the fastest particles)."
+ It also givesmore dissipation when using more particles which makes a convergence study at fixed neighbour number impossible., It also gives dissipation when using more particles which makes a convergence study at fixed neighbour number impossible.
+" Below, when we discussrpSPH for viscous flow, we show that the effective numerical viscosity of standard SPH is non-Newtonian and very large which explains why standard SPH is inadequate to modeling fluids in general."," Below, when we discuss for viscous flow, we show that the effective numerical viscosity of standard SPH is non-Newtonian and very large which explains why standard SPH is inadequate to modeling fluids in general."
+" The results forrpSPH are summarized in Figure 2 which is to be compared to the bottom panel for SPH, plotting the kinetic energy in the box as a function of sound crossing times."," The results for are summarized in Figure \ref{fig:reduced-S} which is to be compared to the bottom panel for SPH, plotting the kinetic energy in the box as a function of sound crossing times."
+ Note that the y-axis in the two panels of Figure 2 is different by a factor of 30., Note that the $y$ -axis in the two panels of Figure \ref{fig:reduced-S} is different by a factor of 30.
+ We should note that this test problem when used in typical Cartesian grid codes will give zero numerical dissipation to machine precision because the uniform nature of the flow along the grid axes., We should note that this test problem when used in typical Cartesian grid codes will give zero numerical dissipation to machine precision because the uniform nature of the flow along the grid axes.
+ ? considered seemingly similar tests in his thesis., \cite{PriceThesis2004} considered seemingly similar tests in his thesis.
+" However, note that there an isothermal equation of state was used."," However, note that there an isothermal equation of state was used."
+" He shows one case with c,—0, i.e. a pressure less fluid and another case with c,—0.05."," He shows one case with $c_s=0$, i.e. a pressure less fluid and another case with $c_s=0.05$."
+ His initial shear profile is like the one we choose here but with twice the amplitude so the fastest particles move with unit speed., His initial shear profile is like the one we choose here but with twice the amplitude so the fastest particles move with unit speed.
+ The pressure-less case is irrelevant here since without pressure forces the discretization of the momentum equation cannot make a difference., The pressure-less case is irrelevant here since without pressure forces the discretization of the momentum equation cannot make a difference.
+" For the second case with c,—0.05 he gives the result only after one sound crossing time at t— 20.", For the second case with $c_s=0.05$ he gives the result only after one sound crossing time at $t=20$ .
+ 'The density fluctuations have grown to order of a percent of the initial value after that single soundcrossing time., The density fluctuations have grown to order of a percent of the initial value after that single soundcrossing time.
+"After estimating A; and A4 for the 70 ancl LOO Cll channels by using the fit with the ALLIWE (see the previous subsection). we can obtain the estimated values of. the amplitude. οἱ, and the slope. a for each half sky by applving eqs. (","After estimating $K_3$ and $K_4$ for the 70 and 100 GHz channels by using the fit with the MHWF (see the previous subsection), we can obtain the estimated values of the amplitude, $A$, and the slope, $\alpha$ for each half sky by applying eqs. ("
+82). (33) ancl (34).,"32), (33) and (34)."
+ LE we consider ten half skies. we can eive an average value and the rms deviation.," If we consider ten half skies, we can give an average value and the rms deviation."
+ The results can be seen in. Table 5. where they are compared with those obtained from the point source only simulations: in this last case we have used the values AI7 and caleulated previously from the point source maps.," The results can be seen in Table 5, where they are compared with those obtained from the point source only simulations; in this last case we have used the values $K_{3}^{ps}$ and $K_{4}^{ps}$, calculated previously from the point source maps."
+ The resultsIUD. at 70 CGllz for a dl Jv Dux limit are: à=2.19dx0.43. A=24344.3. whereas we obtain from the point source simulations: àP=242E0406. A=22.1c 1.5.," The results at 70 GHz for a 1 Jy flux limit are: $\alpha=2.19\pm 0.43$, $A=24.3\pm
+4.3$, whereas we obtain from the point source simulations: $\alpha^{ps}=2.32\pm 0.06$, $A^{ps}=22.1\pm 1.5$ ."
+ For this lux limit. the average relative error of a (1) with respect to al CAP) is below 1054 (15%).," For this flux limit, the average relative error of $\alpha$ $A$ ) with respect to $\alpha^{ps}$ $A^{ps}$ ) is below $10\%$ $15\%$ )."
+ Lhe estimation is not bad. out. the rms deviation is high.," The estimation is not bad, but the rms deviation is high."
+ We can obtain reasonable estimations down to 0.7 Jv., We can obtain reasonable estimations down to 0.7 Jy.
+ Leis clear that small errors in he determination of the cumulants produce high errors in he unplitude and slope., It is clear that small errors in the determination of the cumulants produce high errors in the amplitude and slope.
+ The values obtained from the fit for a flux limit of 1 Jv at 100 Cillz are: à= 3.200.190. 24=21.024.2. whereas from he point source simulations we estimate: a=2.350.10. AP=20.0€3.7.," The values obtained from the fit for a flux limit of 1 Jy at 100 GHz are: $\alpha=2.26\pm 0.19$ , $A=21.0\pm 4.2$, whereas from the point source simulations we estimate: $\alpha^{ps}=2.35\pm 0.16$, $A^{ps}=20.0\pm 3.7$."
+ The average relative error of a (Ενα respect to a7 CAP?) is below 5% (670)., The average relative error of $\alpha$ $A$ )with respect to $\alpha^{ps}$ $A^{ps}$ ) is below $5\%$ $6\%$ ).
+ We obtain reasonable values down to 0.7 Jy. with less dispersion than for the 70 Cillz channel.," We obtain reasonable values down to 0.7 Jy, with less dispersion than for the 70 GHz channel."
+ In Figure 3. the fit for both channels at 1 Jy is compared to the dillerential number counts of the De Zotti mocel. which was used for the simulations and with the estimation from the point source simulations.," In Figure 3, the fit for both channels at 1 Jy is compared to the differential number counts of the De Zotti model, which was used for the simulations and with the estimation from the point source simulations."
+ As can be seen in the plots. the fit is quite good for fluxes between 0.1 Jy. and 1 Jv.," As can be seen in the plots, the fit is quite good for fluxes between 0.1 Jy and 1 Jy."
+ We have considered a natural ecneralization of the ALIWI on the plane. #7. based on the iterative application of the Laplacian operator to the MIIWI.," We have considered a natural generalization of the MHW1 on the plane, $R^2$, based on the iterative application of the Laplacian operator to the MHW1."
+ We have called this group of wavelets the Mexican. Lat Wavelet) Family (MIIN (González-Nuevoctal.2006:López-Caniegoal.E) 2006).," We have called this group of wavelets the Mexican Hat Wavelet Family (MHWF) \citep{gon06,lop06}."
+. We have applied the MEINE to the study of the non-Gaussianity produced in CAIB maps by extragalactic point sources. a very important issue for the separation an analysis of cillerent non-Gaussian components in the CMD.," We have applied the MHWF to the study of the non-Gaussianity produced in CMB maps by extragalactic point sources, a very important issue for the separation and analysis of different non-Gaussian components in the CMB."
+ We have first. derived a formula. cq. (," We have first derived a formula, eq. ("
+14). which relates the cumulants of 2D point. source maps filtere with the MIIWE with the cumulants of the original poin source maps.,"14), which relates the cumulants of 2D point source maps filtered with the MHWF with the cumulants of the original point source maps."
+ From this formula. we have obtained severa equations. eqs. (," From this formula, we have obtained several equations, eqs. ("
+18-23). which relate the contribution of point sources to the skewness and. kurtosis of wavele (MEINE) filtered 2D CMD maps to their contribution to the skewness and kurtosis of beam filtered CMD maps.,"18-23), which relate the contribution of point sources to the skewness and kurtosis of wavelet (MHWF) filtered 2D CMB maps to their contribution to the skewness and kurtosis of beam filtered CMB maps."
+ According to these formulas. the skewness of a wavele filtered map is proportional to A7. with A the amplification. and the kurtosis is proportional to AT.," According to these formulas, the skewness of a wavelet filtered map is proportional to $\lambda^{3}$, with $\lambda$ the amplification, and the kurtosis is proportional to $\lambda^{4}$."
+ This means that the higher the amplification the higher the non-Ciaussian signa will be., This means that the higher the amplification the higher the non-Gaussian signal will be.
+ Vhereforce. the MEIIW2 ancl MEIN at their optima scales can be very suitable filters to detect the non-Gaussian signal due to point sources. since they produce a higher amplification than the MEINI López-Caniegoetal... 2006).," Therefore, the MHW2 and MHW3 at their optimal scales can be very suitable filters to detect the non-Gaussian signal due to point sources, since they produce a higher amplification than the MHW1 \citep{gon06,lop06}."
+. We have then considered. Ciaussian simulations of CAB plus anisotropic noise and compared them. with simulations including all the relevant components: CALB. anisotropic noise. point sources ancl Galactic foregrouncds.," We have then considered Gaussian simulations of CMB plus anisotropic noise and compared them with simulations including all the relevant components: CMB, anisotropic noise, point sources and Galactic foregrounds."
+ The simulations are 2D maps of 12.8«. square degrees with the characteristics of the Planck. experiment at 70 Cllz ancl 100 Cllz., The simulations are 2D maps of $12.8\times12.8$ square degrees with the characteristics of the Planck experiment at 70 GHz and 100 GHz.
+ We perform a Caussianity test. based on the skewness and kurtosis of the maps filtered with the ΔΗΛ. ΛΗΛ and. MIINW at the optimal scales (those which produce the highest amplifications). we compare 1000 simulations of CIB plus noise with 1000 simulations including point sources up to different [ux limits from 0.1 Jv to 1 Jy.," We perform a Gaussianity test based on the skewness and kurtosis of the maps filtered with the MHW1, MHW2 and MHW3 at the optimal scales (those which produce the highest amplifications), we compare 1000 simulations of CMB plus noise with 1000 simulations including point sources up to different flux limits from 0.1 Jy to 1 Jy."
+ The power of the test is higher for the MIIN2 and the MIIW3 than for the standard MEINI. the kurtosis test is more powerful than the skewness test.," The power of the test is higher for the MHW2 and the MHW3 than for the standard MHW1, the kurtosis test is more powerful than the skewness test."
+ We could. detect the non-Gaussianity due to point sources down to a flux limit of 0.4 Jv (70 Gllz) and 0.3 Jy (100 CGLIz). the powers are higher for 100 €illz than for 70 Cllz.," We could detect the non-Gaussianity due to point sources down to a flux limit of 0.4 Jy (70 GHz) and 0.3 Jy (100 GHz), the powers are higher for 100 GHz than for 70 GHz."
+ The results can be seen in Figure 2 and. Table 1., The results can be seen in Figure 2 and Table 1.
+ We would like to point out that according to López-Caniegoetal.(2006). point sources could be detected in these channels. assuming a 5% of spurious detections. down to 0.47 Jv (70 Cllz)and 0.25 Jv (100 11).," We would like to point out that according to \cite{lop06}
+ point sources could be detected in these channels, assuming a $5\%$ of spurious detections, down to $0.47$ Jy (70 GHz)and $0.25$ Jy (100 GHz)."
+ This means that. using the MIIWE. we can see à non-Gaussian signal due to weak point sources. which are in the limit of detection of current. source. extraction techniques.," This means that, using the MHWF, we can see a non-Gaussian signal due to weak point sources, which are in the limit of detection of current source extraction techniques."
+ Therefore. it is very important to characterise the non-Ciaussianity due to unsubtracted: point sources. in order to avoid its possible misidentification as intrinsic non-," Therefore, it is very important to characterise the non-Gaussianity due to unsubtracted point sources, in order to avoid its possible misidentification as intrinsic non-Gaussianity."
+ The influence ofGalactic foregrouncds is negligible. since when we consider simulations without point sources. but keeping the Galactic foregrouncds. the powers go down below 6%.," The influence of Galactic foregrounds is negligible, since when we consider simulations without point sources, but keeping the Galactic foregrounds, the powers go down below $6\%$."
+ We have also studied the performance of the Daubechies 4 wavelet. obtaining powers below 50% even for tux limits as high as 1 Jw.," We have also studied the performance of the Daubechies 4 wavelet, obtaining powers below $50\%$ even for flux limits as high as 1 Jy."
+ Phe MIIWE performs much better than this wavelet (see Figure 2). since it is adapted to deal with point SOULCCS.," The MHWF performs much better than this wavelet (see Figure 2), since it is adapted to deal with point sources."
+ This method could. be used. for all the frequencies of the Planck. experiment and could also be extended: to spherical CMD images. as the ones provided by the Planck experiment.," This method could be used for all the frequencies of the Planck experiment and could also be extended to spherical CMB images, as the ones provided by the Planck experiment."
+ We have also tried to estimate the third order cumulant dys ancl the fourth order cumulant fo. produced. by point sources at the 70 and 100 CGllz Planck channels.," We have also tried to estimate the third order cumulant $K_3$ and the fourth order cumulant $K_4$, produced by point sources at the 70 and 100 GHz Planck channels."
+ We filter 12.8 square degrees. CAIB maps with the characteristics of the Planck experiment with the MEINI. ΔΗΛ ancl MIIWA2 at. three. scales. for. each wavelet (including the optimal ones).," We filter $12.8\times 12.8$ square degrees CMB maps with the characteristics of the Planck experiment with the MHW1, MHW2 and MHW3 at three scales for each wavelet (including the optimal ones)."
+ The maps include C'MD anisotropic noise. point sources and other foregrounds.," The maps include CMB, anisotropic noise, point sources and other foregrounds."
+ We calculate the third. and fourth. order. cumulants of these wavelet filtered maps and write chi-squared functions. eqs. (," We calculate the third and fourth order cumulants of these wavelet filtered maps and write chi-squared functions, eqs. ("
+24) and (25). which allow us to estimate the cumulants of the original point source maps: we have used. eq. (,"24) and (25), which allow us to estimate the cumulants of the original point source maps; we have used eq. ("
+14) to relate the eumulants of the wavelet filtered maps with those of the original maps.,14) to relate the cumulants of the wavelet filtered maps with those of the original maps.
+ We have compared our estimated values with those. A and A17. obtained directly [rom simulations of the same tvpe but including only unfiltered point sources without any other component.," We have compared our estimated values with those, $K_3^{ps}$ and $K_4^{ps}$, obtained directly from simulations of the same type but including only unfiltered point sources without any other component."