source,target
 Seeing couditious varied from 0.5 to 0.6 arcsec between suele exposures. with an average airniass of 1.06.," Seeing conditions varied from 0.5 to 0.6 arcsec between single exposures, with an average airmass of 1.06."
 Data reduction las been performed in Carching bv ↕⋯∖∐∐⋈∖↥⋅↴∖↴∪↕≯∏∐∖⊏≋↻≋∖⊽↑↸∖, Data reduction has been performed in Garching by members of the ESO SV team.
⋜⋯↕∙↕⋟⋜∐⋅↑↕↸⊳∏↕⋜∐⋅↸⊳⋜∐⋅↸∖↖↖↽⋜↧↴∖↴∏↴∖↴↸∖≼↧ for the basic reduction steps since the CCD suffered for color dependent blemishes. with the largest one located practically at the ceuter.," Particular care was used for the basic reduction steps since the CCD suffered for color dependent blemishes, with the largest one located practically at the center."
 Variations in the CCD sensitivity caused by dust erains were also noted by coluparing flat field nuages taken iu different nights, Variations in the CCD sensitivity caused by dust grains were also noted by comparing flat field images taken in different nights.
 All these problems made dathelding werv tricky and requested a non-stancdad mrocedure., All these problems made flatfielding very tricky and requested a non-standard procedure.
 Data were thus flatiiclded using flats obtained directle from the sky by inediau-combining several science exposures taken im different niehts., Data were thus flatfielded using flats obtained directly from the sky by median-combining several science exposures taken in different nights.
 The use of these lead finally to a flatfieldine accuracy for wide-baud filters down to I much higher than achievable either with dome or twilight fats.," The use of these lead finally to a flatfielding accuracy for wide-band filters down to $1\%$, much higher than achievable either with dome or twilight flats."
 Bias frames were obtained nearly every dav and show no noticeable structure or changes with iue., Bias frames were obtained nearly every day and show no noticeable structure or changes with time.
 Bias subtraction and flatfielding were porfonnued usus the ΠΑΕ package., Bias subtraction and flatfielding were performed using the IRAF package.
 huages were then corrected for the CCD dithering and combined using he IRAF tasks audincombine., Images were then corrected for the CCD dithering and combined using the IRAF tasks and.
 Photometric calibrations have been performed by imaging the staudard star PO1633009 from a Landolt field., Photometric calibrations have been performed by imaging the standard star PG1633+099 from a Landolt field.
 The zero-poiut was computed applving the IRAF photometry packagedigiphot. with a final accuracy of 0.03 inagnuitudes.," The zero-point was computed applying the IRAF photometry package, with a final accuracy of 0.03 magnitudes."
 Astrometry on the image has been performed using as a reference the coordinates of a few field stars extracted from the USNO catalogue using the ESO Skycat tool., Astrometry on the image has been performed using as a reference the coordinates of a few field stars extracted from the USNO catalogue using the ESO Skycat tool.
 The pixel-to-sky coordinates transformation has been thus computed using the UN STARLINK software ASTROM (Wallace. 1990). leadius to a final accuracy of &0.5 arcsec on our absolute astrometry after takiug in the consideration both the raus.," The pixel-to-sky coordinates transformation has been thus computed using the UK STARLINK software ASTROM (Wallace, 1990), leading to a final accuracy of $\simeq 0.5$ arcsec on our absolute astrometry after taking in the consideration both the r.m.s."
 of the astrometric fit (+0.3 aresec] and the average uncertainty on the USNO coordinates (270.25 arcsec)., of the astrometric fit $(\simeq 0.3$ arcsec) and the average uncertainty on the USNO coordinates $(\simeq 0.25$ arcsec).
 The most recent radio position of Ll. also reported in Tab.2 of the paper of Chakrabarty Kaspi," The most recent radio position of $-$ 44, also reported in Tab.2 of the paper of Chakrabarty Kaspi"
"Jandthe —module Jaregivenbyto ZT El= Furthermore, let r:=00)..(4.1) A grading function is defined by delta((Z!):= and thus, Define for any AA:— Here is a graded vector space according to 17 and denotes the set of all formal sums T;),where and[10&,.","and the -module are given by Z^m_n E^l_k:=. Furthermore, let r:=. A grading function is defined by (Z^l_k):= and thus, Define for any A:= Here is a graded vector space according to \cite{GazorYuFormal,GazorYuSpec} and denotes the set of all formal sums where and."
". The action of on Jisgivenby((S, for any ((S,T) and We first remove the linear part from the system with a linear change of state variable5."," The action of on is given by for any (S, and We first remove the linear part from the system with a linear change of state variable."
"3.6].. Therefore, we may assume Since d??(0,Z1)= d??(0,Z9) —(,—,89* both terms of and are eliminated from the second level extended partial orbital normal form, = 8*P* Now let min{kk|, 00, (In case is not finite, an arbitrary choice for trivially completes the proof.)"," Therefore, we may assume Since (0,Z^1_1)= (0,Z^0_1) ^0_1) both terms of and are eliminated from the second level extended partial orbital normal form, = ^0_k Now let k| 0, (In case is not finite, an arbitrary choice for trivially completes the proof.)"
" Denote = ((0,...,,0,ZB,m where the number of zeros is 22k--times for1."," Denote := ,0, Z^0_m, where the number of zeros is -times for."
". Then, = ((0,0, 20, Επι) pr(4.2)dimer)=z—Lm-—1PER for any By Theorem there exist transformations that send the vector field into the infinite level orbital normal form πο Μο... For any 1, there exists a nonzero state solution JandatimesolutionTT )suchthat((T, S)*u=—Zlu."," Then, = (0,0, Z^0_m, ^0_r }, for any By Theorem \ref{LNF} there exist transformations that send the vector field into the infinite level orbital normal form ^0_r for some 2, ^0_1, For any 1, there exists a nonzero state solution and a time solution such that (T, S)*u=-Z^l_ku."
" Therefore, the symmetry group for is not finitely generated."," Therefore, the symmetry group for is not finitely generated."
 Now the linear change of state variable )inEquation(4.2) )isobtained., Now the linear change of state variable in Equation \ref{inftONF}) ) is obtained.
T hecompositionofthismapwiththesymmetrygroupassociatedwithuu results in a symmetry group for The rest of the proof follows Theorem ∙, The composition of this map with the symmetry group associated with results in a symmetry group for The rest of the proof follows Theorem \ref{SNF}. .
Iu the following models of several extragalactic SSC aud ΑΠΟΤΟ sources are discussed.,In the following models of several extragalactic SSC and CMB-IC sources are discussed.
 The exact definitions of the model parameters can be found in the Appendix., The exact definitions of the model parameters can be found in the Appendix.
 We eive lints on the detectability of theses sources by the iustruineuts listed iu Tab., We give hints on the detectability of theses sources by the instruments listed in Tab.
 d. based on rough detection sienificance estimates as described in Sect. 3..," \ref{tab:sensitivity}
 based on rough detection significance estimates as described in Sect. \ref{sec:newinstruments}."
" Our high huuinositv radio galaxy is assuned to have spherical lobes with a diameter of LO kpc. ae ainagnctic field strength of 35 pC. The clectrou population is a single power-law with spectral iudex of s= 2.1. which extends frou py=10 to po=3105 aud which is normalized by C=6.63-10temὉ, "," Our high luminosity radio galaxy is assumed to have spherical lobes with a diameter of 40 kpc, and a magnetic field strength of 35 $\mu$ G. The electron population is a single power-law with spectral index of $s = 2.4$ , which extends from $p_{1} = 10$ to $p_{2} =
3\,\cdot 10^4$ and which is normalized by $C = 6.63\cdot10^{-4}\, {\rm
cm}^{-3}$."
This corresponds to an cherey content iu relativistic electronsof 19-1p? cre and in maenetic fields of L8:1077 ere., This corresponds to an energy content in relativistic electronsof $4.9\cdot 10^{59}$ erg and in magnetic fields of $4.8\cdot 10^{58}$ erg.
 The energy ratio of electronic to magnetic euergv is LO. aud therefore by onlv a factor of a few above what racio-astronomers call equipartition. for which the energy of clectrous racliating above LO MIIz equals that of the magnetic fields.," The energy ratio of electronic to magnetic energy is 10, and therefore by only a factor of a few above what radio-astronomers call equipartition, for which the energy of electrons radiating above 10 MHz equals that of the magnetic fields."
 The Iuuinositv at 1 GIIz is 3.9-10°ore/TIz/s.," The luminosity at 1 GHz is $3.9\cdot10^{34}\,{\rm
erg/Hz/s}$."
 τος A. for exaluple. has a flux of 211 Jv at 5 GIIz (Stull.1971).. as two of our model cocoons if they are placed at the Iuimuinositv distance of 265 Mpc. similar to that of Cveuus A for Hy=ODkan/s/AIpe and gy=0.5 (Owenetal.1997.forthe redshift)...," Cygnus A, for example, has a flux of 214 Jy at 8 GHz \cite{1971AJ.....76....1S}, as two of our model cocoons if they are placed at the luminosity distance of 265 Mpc, similar to that of Cygnus A for $H_0 = 65\,{\rm km/s/Mpc}$ and $q_0 = 0.5$ \cite[for the
redshift]{1997ApJ...488L..15O}."
 The svuchrotrou aud SSC ceutra surface brightuess of our radio cocoons are shown in Fie. 0.., The synchrotron and SSC central surface brightness of our radio cocoons are shown in Fig. \ref{fig:CygAexp}.
 Iu this figure possible later stages in the evolution of the radio source. after the supplv with fresh electrons was shut down. are also shown.," In this figure possible later stages in the evolution of the radio source, after the supply with fresh electrons was shut down, are also shown."
 In order to model these stages we asstuue that the radio cocoons will expat adiabatically while they vise buovautlv in the cluster atmosphere (Clurazovetal.2001:Briigecu&Ixaiscr. 2001)..," In order to model these stages we assume that the radio cocoons will expand adiabatically while they rise buoyantly in the cluster atmosphere \cite{2001ApJ...554..261C,2001MNRAS.325..676B}."
 Iu our illustrating example we assume that the jet supply of fresh radio plasima shuts down at the preseut evolutionary stage of Cyeguus A. The two radio cocoons are assuned to expand with hb=3/2 aud ty=75 Myr (see Eq. A3))., In our illustrating example we assume that the jet supply of fresh radio plasma shuts down at the present evolutionary stage of Cygnus A. The two radio cocoons are assumed to expand with $b = 3/2$ and $t_0= 75$ Myr (see Eq. \ref{eq:expansion}) ).
" This should serve as a rough model for a rise with half sonic velocity in a sineular isothermal cluster atmosphere (cluster density * cluster 3], ", This should serve as a rough model for a rise with half sonic velocity in a singular isothermal cluster atmosphere (cluster density $\propto$ cluster $^{-2}$ ).
We note that he evolution of the electron spectra is not corrected. or the reduced svuchrotron cooling at self absorbed frequencies., We note that the evolution of the electron spectra is not corrected for the reduced synchrotron cooling at self absorbed frequencies.
 An investigation of this more complicated problem was doue by Rees (1967) ]t is obvious. that a radio cocoon of a powerful. conipact radio galaxy produces a detectable SSC flux for a long time after the the svuchrotrou radio emission vanished from the radio observable frequency band.," An investigation of this more complicated problem was done by Rees \cite*{1967MNRAS.136..279R}
 It is obvious, that a radio cocoon of a powerful, compact radio galaxy produces a detectable SSC flux for a long time after the the synchrotron radio emission vanished from the radio observable frequency band."
 At low radio frequencies the SSC fux remains even for a few Cor. due to the slow cuerev losses by sclfabsorbed svuchrotrou chussion.," At low radio frequencies the SSC flux remains even for a few Gyr, due to the slow energy losses by self-absorbed synchrotron emission."
 But adiabatic losses due the expansion of the source are able to reduce the SSC fiux substautiallv., But adiabatic losses due the expansion of the source are able to reduce the SSC flux substantially.
 The typical angular area is of the order of arcani., The typical angular area is of the order of $^2$.
 Therefore all of the instruments listed in Tab., Therefore all of the instruments listed in Tab.
 1. should be able to detect a Cyeuus-Á like radio cocoon., \ref{tab:sensitivity} should be able to detect a Cygnus-A like radio cocoon.
 All but LOFAR aud CAIRT see it best shortly after he decay of the svuchrotron emission., All but LOFAR and GMRT see it best shortly after the decay of the synchrotron emission.
 The latter is because at below CGIIz frequencies the source clission is risiue with time during the first few 100 Myr., The latter is because at below GHz frequencies the source emission is rising with time during the first few 100 Myr.
 PLANCTS and HERSCTIEL should eive oulv relative week detection of carly stages., PLANCK and HERSCHEL should give only relative week detection of early stages.
" Iowever. if such a very powerful radio ghost would be much: closer. e.g. located iu the Virgo cluster. even later PAages may be detectable with these ιβαποτς,"," However, if such a very powerful radio ghost would be much closer, e.g. located in the Virgo cluster, even later stages may be detectable with these instruments."
 GPS are believed to be the τον carly stage of radio ealaxies. im which a small (< kpc) radio cocoon is working its wav through the dense interstellar medimu of a radioealaxy (Suellenetal.2000:DeZottiab.2000.and therem)..," GPS are believed to be the very early stage of radio galaxies, in which a small $<$ kpc) radio cocoon is working its way through the dense interstellar medium of a radiogalaxy \cite[and references
therein]{2000MNRAS.319..445S,2000A&A...354..467D}. ."
 The spectral peak at typically a CIIz is likely due to svuchrotron-solf-absorptiou at lower frequencies., The spectral peak at typically a GHz is likely due to synchrotron-self-absorption at lower frequencies.
 Although it turns out that SSC emission, Although it turns out that SSC emission
Solar flare iupuldve clectrou events present an alternative to the more traditional hard N-ray diagnostics of poorly understood acceleration and transport of solar cuereetic electrons.,Solar flare impulsive electron events present an alternative to the more traditional hard X-ray diagnostics of poorly understood acceleration and transport of solar energetic electrons.
 While hard X-rav observations provide insight mto energetic electrous ij the lower deuse solar atinosphere (es.—777). nupulsive solar olectrou eveuts (e.g.77) provide crucial information about escapiug electrons from the acceleration region.," While hard X-ray observations provide insight into energetic electrons in the lower dense solar atmosphere \citep[e.g.][]{Arnoldy_etal1968,DennisSchwartz1989,BrownKontar05}, impulsive solar electron events \citep[e.g.][]{Lin1985,Krucker_etal07} provide crucial information about escaping electrons from the acceleration region."
 Because of the rather Buited spatial resolution of past and current hard X-ray. observatious. even the spatially resolved lard ταν spectrum of energetic electrons with RITESSI(c.g.2) electronisa convolution of transport effects aud possibly acceleration (?)..," Because of the rather limited spatial resolution of past and current hard X-ray observations, even the spatially resolved hard X-ray spectrum of energetic electrons with RHESSI \citep[e.g.][]{Emslie_etal2003} is a convolution of transport effects and possibly electron acceleration \citep{Brown_etal2009}."
" Ποσο, de-convolution of the electron accelerated spectrum aud accelerator properties from hard N-ravs is a non-trivial task using both forwarcd-anodelling (?7?) and modoel-àiudepeudenut techniques (??).."," Hence, de-convolution of the electron accelerated spectrum and accelerator properties from hard X-rays \citep{Brown_etal2006} is a non-trivial task using both forward-modelling \citep{Holman_etal2003,Kasparova_etal2005} and model-independent techniques \citep{Piana_etal2003,Kontar_etal2005}."
 Solar electron iupulsve events propagate outward through the aliiost collisionless plasina of the solar corona and solar wind (?).., Solar electron impulsive events propagate outward through the almost collisionless plasma of the solar corona and solar wind \citep{Lin1985}.
 Even with this collixiouless regie the energetic electrous cau interact with plasma via generation aud isorption of electrostatic plasima waves., Even with this collisionless regime the energetic electrons can interact with plasma via generation and absorption of electrostatic plasma waves.
 Iu the standard scenario (?).. the non-linear interaction of beam-driven plasma waves leads to the appearance of rather strong radio enüsson - type 1 solar/iuterplauctary radio bursts.," In the standard scenario \citep{GinzburgZhelezniakov1958}, the non-linear interaction of beam-driven plasma waves leads to the appearance of rather strong radio emission - type III solar/interplanetary radio bursts."
 The observatious of type III solar bursts aud energetic particles (2777). as well as theoretical (2???) iid nuncerical investieatious (2°222222?) provide strong support to the standard type ITE model.," The observations of type III solar bursts and energetic particles \citep{Linetal1981,Ergun_etal98,Gosling_etal2003,Krucker_etal07} as well as theoretical \citep{ZheleznyakovZaitsev1970,Zaitsev_etal72,Melnik1995} and numerical investigations \citep{MagelssenSmith1977,Grognard1982,Kontar_etal1998,Yoon_etal2000,Kontar2001_a,Li_etal2006a,Ledenev04,Krasnoselskikh_etal2007} provide strong support to the standard type III model."
 The plaua of the solar corona aud the solar wind is a ποπο turbulent medium with density perturbations at various leugth scales., The plasma of the solar corona and the solar wind is a non-uniform turbulent medium with density perturbations at various length scales.
 The structure of the solar wind density fluctuations have been analyses using sciutillatious of siaLsize radio sources (6.8.TT). ," The structure of the solar wind density fluctuations have been analysed using scintillations of small-size radio sources \citep[e.g.][]{Hollweg1970,Young1971}. ."
Iu-situ measurements have :uso been used to determine the deusitv spectrum near 1e Earth aud between 0.3 ane LAU withMelos (2).., In-situ measurements have also been used to determine the density spectrum near the Earth and between $0.3$ and $1$ AU with \citep{MarschTu1990}.
 While the detailed structre of the density turbulence iu the àmer lieliosplere is not wel established. the density fluctuation spectrum rear the Earth seciis close to a power-law spectrum with spectra index about 5/3. sinular to earlier observations.," While the detailed structure of the density turbulence in the inner heliosphere is not well established, the density fluctuation spectrum near the Earth seems close to a power-law spectrum with spectral index about $5/3$, similar to earlier observations."
 It has been recognized {7T) that beanicdriven Langulr Waves cal be effectively altered by even weak density eradicuts.," It has been recognized \citep{Ryutov1969,KarpmanIstomin1974} that beam-driven Langmuir waves can be effectively altered by even weak density gradients."
 Therefore the effect of density fluctuations ο1 beali-driven plasma waves responsible for type III radio bursts has been considered both uuuercallv. arc analytically (0?7)..," Therefore the effect of density fluctuations on beam-driven plasma waves responsible for type III radio bursts has been considered both numerically and analytically \citep{Melroseetal1986,Robinsonetal1992,Kontar2001_solphys}."
" Density fluctuations are believed. to SUppress plasma wave growth (77) aud|IC τοspousidle for 1C chuupy plasiua. wave distrinition observec ""msiti near the Earth(?7).."," Density fluctuations are believed to suppress plasma wave growth \citep{SmithSime1979,Muschiettietal1985} and be responsible for the clumpy plasma wave distribution observed in-situ near the Earth \citep{Gurnettetal1978,Linetal1981}."
 The fluctuations. whils changing 1C distribution of plasma waves significantly. ive a rather weak modulation effect on t10 lustantancous «istribution of electrous. (?)..," The fluctuations, whilst changing the distribution of plasma waves significantly, have a rather weak modulation effect on the instantaneous distribution of electrons \citep{Kontar2001_a}."
 Receutly. ? have shown that 1C electron beam plasma interaction via Laugniuir waves in the non-uniform solar corona leacs to he appeara1ος of a break energy in t16 obscved spectruu a the Earth and can explain the oISCTVCL Lappareut carlyinjection of ow-enerev electrons.," Recently, \citet{KontarReid2009} have shown that the electron beam plasma interaction via Langmuir waves in the non-uniform solar corona leads to the appearance of a break energy in the observed spectrum at the Earth and can explain the observed apparent early injection of low-energy electrons."
 Towever. the net effec of deusitv fluctuations in the solar wind on the electro) spectrin detected near 1 AU has not been addressed before.," However, the net effect of density fluctuations in the solar wind on the electron spectrum detected near 1 AU has not been addressed before."
 Iu this paper. we investigate the effects of ckerounud ράσα deusitv Πιοuation ou the generation ancl absorption of plasma waves from a high ejergw solar electron. beam travelling fixnu the Sun to the Earth.," In this paper, we investigate the effects of background plasma density fluctuation on the generation and absorption of plasma waves from a high energy solar electron beam travelling from the Sun to the Earth."
 We demoustrate the ccpendence of plasina waves on the evelof deusitv fluctuations. with hieh levels dampine εινα waves too much to be πι accordance wih detected ype II radio ciission.," We demonstrate the dependence of plasma waves on the levelof density fluctuations, with high levels damping plasma waves too much to be in accordance with detected type III radio emission."
 We also show how he level of, We also show how the level of
the flux calibration we extracted a 1-D spectrum of the standare star to find the calibration function: then we extracted a set of I-D spectra of the galaxy summing up a number of lines corresponding to the slit width.,the flux calibration we extracted a 1-D spectrum of the standard star to find the calibration function; then we extracted a set of 1-D spectra of the galaxy summing up a number of lines corresponding to the slit width.
" Since the slit width was 2 anc the scale of the instrument was /0.252""pix. we collapsed eight lines to obtain each. 1-D spectrum."," Since the slit width was $2 ''$ and the scale of the instrument was $0.252 ''/pix$, we collapsed eight lines to obtain each 1-D spectrum."
 Finally we applied the flux calibration to this collection of spectra., Finally we applied the flux calibration to this collection of spectra.
 The wavelength and flux-calibrated spectra are shown in Fig., The wavelength and flux-calibrated spectra are shown in Fig.
 4. and Fig. 5.., \ref{spec1} and Fig. \ref{spec2}.
 The fluxes of the above mentioned emission lines were measured using the IRAF routine. that provides an interactive facility to display and analyze spectra.," The fluxes of the above mentioned emission lines were measured using the IRAF routine, that provides an interactive facility to display and analyze spectra."
 We evaluated flux and equivalent width by marking two continuum points around the line to be measured., We evaluated flux and equivalent width by marking two continuum points around the line to be measured.
 The linear continuum ts, The linear continuum is
any of the other bolometer-based LDBs. such as Boomerang (Lange 1997) or NLANINLA. (Richarels 1997) or even HENE-based LDBs. such as BEAST (Lubin 1997) and derived similar error forecasts.,"any of the other bolometer-based LDBs, such as Boomerang (Lange 1997) or MAXIMA (Richards 1997) or even HEMT-based LDBs, such as BEAST (Lubin 1997) and derived similar error forecasts."
 For most parameters the inclusion of prior constraints on their variation have no effect. particularly for a high precision experiment like Planck.," For most parameters the inclusion of prior constraints on their variation have no effect, particularly for a high precision experiment like Planck."
 Even for MAP the inclusion of priors has little impact. except for variables such as Va. which are poorly constrained. [rom the experiment ione.," Even for MAP the inclusion of priors has little impact, except for variables such as $Y_{He}$ which are poorly constrained from the experiment alone."
 I£ the helium abundance is allowed to float. freely. it has a substantial effect on the other parameters: however. limiting its value to be 0.2300.02 results in little impact on the other numbers.," If the helium abundance is allowed to float freely, it has a substantial effect on the other parameters; however, limiting its value to be $0.23 \pm 0.02$ results in little impact on the other numbers."
 For the LDB|DMIEE case. the errors are susbstantially larger with no controlling priors.," For the LDB+DMR case, the errors are susbstantially larger with no controlling priors."
 As expected. the parameters have correlations among rwemsclyes that range from weak to very strong in all models. anc can dilfer from experiment to experiment as well as model to model.," As expected, the parameters have correlations among themselves that range from weak to very strong in all models, and can differ from experiment to experiment as well as model to model."
 The power amplitude: £C1/2p3 oaud τι. have à correlation coellicient about for SCDAL lor Planck and ALAR., The power amplitude $\avrg{{\cal C}_\ell}_B^{1/2}$ and $\tau_C$ have a correlation coefficient about for SCDM for Planck and MAP.
 The most highly. correlated. are oj and wa. as expected from. the anele-cistance near-degeneracy.," The most highly correlated are $\omega_{k}$ and $\omega_{\Lambda}$, as expected from the angle-distance near-degeneracy."
" In the ENTables. the O4h7.2 numbers are. determined.. with. O.h fixed. and the £,h numbers are determined. with OXh fixed: the other 1parameters are relatively insensitive o fixing either. or neither."," In the Tables, the $\Omega_{\Lambda}{\rm h}^2$ numbers are determined with $\Omega_k{\rm
h}^2$ fixed, and the $\Omega_k{\rm h}^2$ numbers are determined with $\Omega_{\Lambda}{\rm h}^2$ fixed; the other parameters are relatively insensitive to fixing either, or neither."
 Thus. although. our. estimates of errors after. mareinalization are gratifvinglv small for many parameters. especially for the specifications of Planck. hey are large in other cases δν].," Thus, although our estimates of errors after marginalization are gratifyingly small for many parameters, especially for the specifications of Planck, they are large in other cases $\delta \Omega_{\Lambda}{\rm
h}^2$ )."
 Error estimates in square brackets are those obtained when the most correlated component for that. variable is constrained. to of the target value., Error estimates in square brackets are those obtained when the most correlated component for that variable is constrained to be the target value.
 A more natural way to deal with strong correlations between variables is. to perform a xinciple component analysis in parameter space. rank-ordering linear combinations of parameters. as described in Section ??:: some linear combinations are determined exquisitely well and some are less well determined because of near-degeneracies as is illustrated in the Tables.," A more natural way to deal with strong correlations between variables is to perform a principle component analysis in parameter space, rank-ordering linear combinations of parameters, as described in Section \ref{sec:prior}: some linear combinations are determined exquisitely well and some are less well determined because of near-degeneracies as is illustrated in the Tables."
 The Tables also show values obtained in round brackets when positivity constraints on parameters such as 7c: are used., The Tables also show values obtained in round brackets when positivity constraints on parameters such as $\tau_C$ are used.
 The àh/h shown are determined [rom h?=Mw hence it is a derived. rather than fundamental quantity.," The $\delta {\rm h}/ {\rm h}$ shown are determined from ${\rm h}^2 =
\sum_j \omega_j$, hence it is a derived rather than fundamental quantity."
 However. b errors depend upon what is kept fixed and what is varied.," However, ${\rm h}$ errors depend upon what is kept fixed and what is varied."
 Thus we can use h to replace one of d. c. wy. with the other two to be marginalized.," Thus we can use ${\rm h}$ to replace one of $\omega_{m}$, $\omega_{\Lambda}$, $\omega_k$, with the other two to be marginalized."
" In that case. the error estimate would be ób/h=0.50,fh?"," In that case, the error estimate would be $\delta {\rm h}/{\rm h} = 0.5
\delta \omega_j /{\rm h}^2$."
 We find that the estimated. errors on. parameters. are sensitive to their input target-mocdel values., We find that the estimated errors on parameters are sensitive to their input target-model values.
" Table 3. shows results for two other @,=O mocels. a ACDAL model and an LICDAL model."," Table \ref{tab:satparams2} shows results for two other $\Omega_k=0$ models, a $\Lambda$ CDM model and an HCDM model."
 This illustrates the sensitivity of parameter error estimation to relatively modest. changes in the target Cy., This illustrates the sensitivity of parameter error estimation to relatively modest changes in the target ${\rm C}_\ell$.
 In interpreting these tables it is also important to take into account the restrictions that we have imposed on the models., In interpreting these tables it is also important to take into account the restrictions that we have imposed on the models.
 Phe OCDAL model error estimates are derived assuming there is no tensor component., The OCDM model error estimates are derived assuming there is no tensor component.
 Including it has little clleet on the results: even the most correlated. the amplitude. zc: and mi. are only slightly allected.," Including it has little effect on the results: even the most correlated, the amplitude, $\tau_C$ and $n_s$, are only slightly affected."
" Phe angle-distance scaling ensures that the tensor power spectrum does not fall off until hisher £ than in the ο=0 cases. and this leads to a substantial improvement in 07,."," The angle-distance scaling ensures that the tensor power spectrum does not fall off until higher $\ell$ than in the $\Omega_k=0$ cases, and this leads to a substantial improvement in $\delta r_{ts}$."
 We have also found that the errors on rm; and rj; are extremely dependent on the input εν if they are allowed to vary independentIy (ox and TFurner 1994. Ixnox. 1995. Efstathiou 1997).," We have also found that the errors on $ n_t$ and $ r_{ts}$ are extremely dependent on the input $r_{ts}$ if they are allowed to vary independently (Knox and Turner 1994, Knox 1995, Efstathiou 1997)."
" However. hand the various matter densities. &x,5,,etc... are insensitive to the tensor spectrum for reasonable values of ες£;2."," However, ${\rm h}$ and the various matter densities, $\omega_{cdm}$, are insensitive to the tensor spectrum for reasonable values of $r_{ts} \simlt 2$."
 In open universes. features in the power-speetrum shift to larger multipoles according to the angle-distance relation. roughly as (0)σος2 thus. for low Qo. high resolution is required. to determine parameters which allect the Doppler peak structure h ancl the various v).," In open universes, features in the power-spectrum shift to larger multipoles according to the angle-distance relation, roughly as $C(\ell)
\rightarrow C(\ell/\Omega_0^{1/2})$; thus, for low $\Omega_0$, high resolution is required to determine parameters which affect the Doppler peak structure ${\rm h}$ and the various $\omega$ )."
 Phe relative accuracies of the parameters are less sensitive to variations of £2; and bh., The relative accuracies of the parameters are less sensitive to variations of $\Omega_b$ and ${\rm h}$.
 In summary. we have described how το compute the errors in the estimation of cosmological parameters from measurements of the CAIB power spectrum at a number of frequencies with dillerent angular resolutions are sensitivities.," In summary, we have described how to compute the errors in the estimation of cosmological parameters from measurements of the CMB power spectrum at a number of frequencies with different angular resolutions and sensitivities."
 We have also shown how prior information on the values of parameters can be incorporated. into the analysis and. deseribed some of the pitfalls of this type of analysis that can arise if inaccurate cderivatives of the C's are used and if poor parameter choices are adopted., We have also shown how prior information on the values of parameters can be incorporated into the analysis and described some of the pitfalls of this type of analysis that can arise if inaccurate derivatives of the ${\rm C}_\ell$ 's are used and if poor parameter choices are adopted.
 We have applied. our machinery to the MAD ane Planck satellites and find that these missions are capable of determining fundamental cosmological parameters to an accuracy (that far exceeds that from conventiona astronomical techniques., We have applied our machinery to the MAP and Planck satellites and find that these missions are capable of determining fundamental cosmological parameters to an accuracy that far exceeds that from conventional astronomical techniques.
" In particular. Planck is capable of determining the Hubble constant anc the baryon density parameter O, to a precision of a few percent or better for each of the target models listed in Tables 2 and 3.."," In particular, Planck is capable of determining the Hubble constant and the baryon density parameter $\Omega_b$ to a precision of a few percent or better for each of the target models listed in Tables \ref{tab:satparams1} and \ref{tab:satparams2}."
 However. some parameter combinations are poorly determined: by CAB observations alone as described. in Section ?? and Section 3...," However, some parameter combinations are poorly determined by CMB observations alone as described in Section \ref{sec:targetmod} and Section \ref{sec:satparams}."
 Nevertheless. despite this caveat. it is evident from this work that accurate CAIB observations have the potential to revolutionize our knowledge of the kev cosmological parameters describing our Universe.," Nevertheless, despite this caveat, it is evident from this work that accurate CMB observations have the potential to revolutionize our knowledge of the key cosmological parameters describing our Universe."
 We would. like to thank Llovel Knox for. useful discussions., We would like to thank Lloyd Knox for useful discussions.
 JLB was supported by the Canadian Institute for Advanced. Research and. NSERC., JRB was supported by the Canadian Institute for Advanced Research and NSERC.
 GPE acknowledges the award of a PPARC Senior Research Fellowship., GPE acknowledges the award of a PPARC Senior Research Fellowship.
 MT was supported by NASA through a Hubble: Fellowship. 3ELIE-O01084.01-96.X. awareled by the Space Telescope Science Institute. which is operated. by AURA. Inc. uncer NASA contract NAS5-26555.," MT was supported by NASA through a Hubble Fellowship, HF-01084.01-96A, awarded by the Space Telescope Science Institute, which is operated by AURA, Inc. under NASA contract NAS5-26555."
concentrated in between the PAH cations and VSGs.,concentrated in between the PAH cations and VSGs.
 Both the spectra and spatial maps of each signal show high correlation to the results using SpitzerOs IRS-SL mode (??).. allowing us to use these results to verify our conclusions.," Both the spectra and spatial maps of each signal show high correlation to the results using SpitzerÕs IRS-SL mode \citep{Olivier, Olivier10}, allowing us to use these results to verify our conclusions."
 To further explore the origin. of the three resolved signals. we employed the NASA Ames PAH IR Spectroscopic Database.," To further explore the origin of the three resolved signals, we employed the NASA Ames PAH IR Spectroscopic Database."
 The fit shows that the component spectra resolved by BSS could be recreated by an appropriate combinatior of specific classes of PAH spectra from the database., The fit shows that the component spectra resolved by BSS could be recreated by an appropriate combination of specific classes of PAH spectra from the database.
" Then. we used a database to fit an observed spectrum and groupec the individual molecules into charge class. then compared the spectra of the combined charge classes to the BSS extractec PAH™ and PAH"" signals."," Then, we used a database to fit an observed spectrum and grouped the individual molecules into charge class, then compared the spectra of the combined charge classes to the BSS extracted $^+$ and $^0$ signals."
" The components were found to be very similar to the BSS extracted PAH™ and PAH"".", The components were found to be very similar to the BSS extracted $^+$ and $^0$.
 Specific spectral properties are found for each population: By using the BSS method and the PAH Database fit. we arrived at the above conclusions.," Specific spectral properties are found for each population: By using the BSS method and the PAH Database fit, we arrived at the above conclusions."
 Each method has unique yet complementary strengths and weaknesses., Each method has unique yet complementary strengths and weaknesses.
 The BSS method is blind. i.e. has no intrinsic assumptions about the emitting components. however since the statistical properties of the emitting components are unknown. the unmixing is not perfectly efficient.," The BSS method is blind, i.e. has no intrinsic assumptions about the emitting components, however since the statistical properties of the emitting components are unknown, the unmixing is not perfectly efficient."
 Additionally. the BSS method separates 3 mathematically distinct signals. but gives no intuition about the molecular properties of these signals.," Additionally, the BSS method separates 3 mathematically distinct signals, but gives no intuition about the molecular properties of these signals."
 The PAH Database allows for direct physical interpretation of the fit yet is biased towards smaller molecules and lacks spectral information for PAH clusters or other possible carriers of the VSG signal., The PAH Database allows for direct physical interpretation of the fit yet is biased towards smaller molecules and lacks spectral information for PAH clusters or other possible carriers of the VSG signal.
 Although both methods suffer limitations. the strengths of one compensate for the weaknesses of the other.," Although both methods suffer limitations, the strengths of one compensate for the weaknesses of the other."
" Although the database fit may be degenerated in some cases. the interpretation of the y database fit results. in terms of classes. is in agreement with the result of the BSS for PAH and PAH""."," Although the database fit may be degenerated in some cases, the interpretation of the $\chi^2$ database fit results, in terms of classes, is in agreement with the result of the BSS for $^+$ and $^0$."
 The VSG spectrum can only be obtained by BSS. and the database fit of this spectrum provide additional information on the possible chemical nature of this component.," The VSG spectrum can only be obtained by BSS, and the database fit of this spectrum provide additional information on the possible chemical nature of this component."
 Both methods. te. BSS and Database Fitting. are powerful tools. but they be used with an understanding of their limitations (described in details in Appendices A and B).," Both methods, i.e. BSS and Database Fitting, are powerful tools, but they be used with an understanding of their limitations (described in details in Appendices A and B)."
 The efficiency of NMF ts subject to two main limitations: 1) the possible non-unicity of solutions 2) the inaccuracy of the unmixing in the presence of noise., The efficiency of NMF is subject to two main limitations: 1) the possible non-unicity of solutions 2) the inaccuracy of the unmixing in the presence of noise.
 These two problems are the subject of intensive theoretical research in the field of signal processing (see e.g. ? and ?))., These two problems are the subject of intensive theoretical research in the field of signal processing (see e.g. \citealt{donoho} and \citealt{hans}) ).
 If the statistical properties of the source matrix are known. then one can determine whether NMP will function properly.," If the statistical properties of the source matrix are known, then one can determine whether NMF will function properly."
" In a ""real life” case like here however. since the matrix is not knownpriori.. we cannot verify these statistical properties."," In a “real life"" case like here however, since the matrix is not known, we cannot verify these statistical properties."
 We therefore rely on some empirical considerations., We therefore rely on some empirical considerations.
 Procuring the. component. spectra (Figure 3)). the most prominent unmixing artifact is located at 11.0 um and reveals itself with a sharp drop in Signal 3.," Procuring the component spectra (Figure \ref{fig:weightfactors}) ), the most prominent unmixing artifact is located at 11.0 $\mu$ m and reveals itself with a sharp drop in Signal 3."
 Some less discernible artifacts can be found at 11.2 jm as an added peak of Signal 2 (seen in the grey envelope of Figure 3)) and at 12.7um as a dip in intensity of Signal 3. roughly resembling an absorption feature.," Some less discernible artifacts can be found at 11.2 $\mu$ m as an added peak of Signal 2 (seen in the grey envelope of Figure \ref{fig:weightfactors}) ) and at $\mu$ m as a dip in intensity of Signal 3, roughly resembling an absorption feature."
 We suggest that the unmixing artifacts seen in some signals are compensated for by an increase or decrease of intensity. at the same wavelength. in other signals.," We suggest that the unmixing artifacts seen in some signals are compensated for by an increase or decrease of intensity, at the same wavelength, in other signals."
 This is most clearly seen at the 11.0 gam wavelength., This is most clearly seen at the 11.0 $\mu$ m wavelength.
 The sharp drop to 0 of Signal 3 is compensated by the weak satellite feature of Signal |., The sharp drop to 0 of Signal 3 is compensated by the weak satellite feature of Signal 1.
 Similarly. the absorption-type feature seen at 12.7 is compensated by a slightly increased intensity of Signal 2 at 12.7 im. In order to better understand the unmixing efficiency. two tests were conceived:," Similarly, the absorption-type feature seen at 12.7 is compensated by a slightly increased intensity of Signal 2 at 12.7 $\mu$ m. In order to better understand the unmixing efficiency, two tests were conceived:"
both DE and DEDI,both DE and DFDI.
 Ins L.. we compare the RV uncertainties of DE aud DEDI under various conditious.," In \ref{sec:CompDeDFDI}, we compare the RV uncertainties of DE and DFDI under various conditions."
 In 5.. we predict the photou-Iimited performance ofa DEDI-based Doppler instrument iu the NIB aud discuss its potential coutributious to exoplauets search around. M cdewarfs.," In \ref{sec:Performance}, we predict the photon-limited performance of a DFDI-based Doppler instrument in the NIR and discuss its potential contributions to exoplanets search around M dwarfs."
 Iu 86... we discuss the iullueuce of tellurie lines on precision RV measurements aud differeut methods of removing telluric line effects.," In \ref{sec:Telluric}, we discuss the influence of telluric lines on precision RV measurements and different methods of removing telluric line effects."
 A stummary of the study is given iu 37.., A summary of the study is given in \ref{sec:Conclusion}.
 The theory of DEDI has been discussed by several papers (22?)..," The theory of DFDI has been discussed by several papers \citep{Ge2002, Erskine2003, VanEyken2010}."
 We briefly introduce the principle of DEDI in this section. readers may refer to previous references for more detailed discussion.," We briefly introduce the principle of DFDI in this section, readers may refer to previous references for more detailed discussion."
 DEDI is realized by coupling a fixed delayinterferometer with a post-disperser (Fig. 1))., DFDI is realized by coupling a fixed delayinterferometer with a post-disperser (Fig. \ref{fig:DFDI_setup}) ).
 The resulting [ringing spectrum is recorded ou a 2-D detector., The resulting fringing spectrum is recorded on a 2-D detector.
 The formation of the final (riugiug spectrum is illustrated in Fig. 2.., The formation of the final fringing spectrum is illustrated in Fig. \ref{fig:DFDI_illus}.
 Boy.y) is a mathematical representation of the final image formed at the 2-D detector aud it is described by the following equation: where So(7) is the intrinsic stellar spectrum aud 7 is optical frequency.," $B(\nu,y)$ is a mathematical representation of the final image formed at the 2-D detector and it is described by the following equation: where $S_0(\nu)$ is the intrinsic stellar spectrum and $\nu$ is optical frequency."
 Sp is divided by fv to convert energy flux into photon flux., $S_0$ is divided by $h\nu$ to convert energy flux into photon flux.
 JF is the intensity trausinissiou Luuction (Equation (2))). 9 is the coordinate aloug the slit direction which is trausverse to dispersion direction. ο) represents convolution and LSF is the line spread function of the post-disperser which is a funcetiou of 7 aud spectral resolution A.," $IT$ is the intensity transmission function (Equation \ref{eq:IT}) )), $y$ is the coordinate along the slit direction which is transverse to dispersion direction, $\otimes$ represents convolution and $LSF$ is the line spread function of the post-disperser which is a function of $\nu$ and spectral resolution $R$."
 In Equation (2)): 5 is visibility lor a given frequency chanuel. the ratio of half ol the peak-valley ziuplitude and the DC offset. which is determiued by stellar flux Sy): c is the speed of light: aud 7 is the optical path difference (OPD) of the interferometer which is designed to be tilled along the slit direction such that several [riuges are formed aloug each 7 channel (Middle. Fig. 2)).," In Equation \ref{eq:IT}) ): $\gamma$ is visibility for a given frequency channel, the ratio of half of the peak-valley amplitude and the DC offset, which is determined by stellar flux $S_0(\nu)$; $c$ is the speed of light; and $\tau$ is the optical path difference (OPD) of the interferometer which is designed to be tilted along the slit direction such that several fringes are formed along each $\nu$ channel (Middle, Fig. \ref{fig:DFDI_illus}) )."
 We assume the LSF is a gaussian function (Equation (3))). Av=v/R/2.35 because we asstuue that oue resolution element is equal to the FWHM of a spectral liue.," We assume the LSF is a gaussian function (Equation \ref{eq:LSF}) )), $\Delta\nu=\nu/R/2.35$ because we assume that one resolution element is equal to the FWHM of a spectral line."
 Fig., Fig.
 3. shows high-resolution (0.005 sspacing) svuthetic spectra of M. dwarfs with solar metallicity (22)...," \ref{fig:Wav_Flux} shows high-resolution (0.005 spacing) synthetic spectra of M dwarfs with solar metallicity \citep{Hauschildt1999,Allard2001}."
 Teg ranges from 21001. to 91001. aud logg is 1.5.," $T_{\rm{eff}}$ ranges from 2400K to 3100K, and $\log g$ is 4.5."
 No rotational broacdeniug is added in the spectrum., No rotational broadening is added in the spectrum.
 Most absorption lines are shallow with FWHAIs of several tenths of an A., Most absorption lines are shallow with FWHMs of several tenths of an .
. Since RV information is embecdecd in the slope, Since RV information is embedded in the slope
was detected on 2003 229 during a scan of the Galactic plane by the IBIS/ISGRI soft gamma-ray detector onboard the International Gamma Ray Laboratory2003).,was detected on 2003 29 during a scan of the Galactic plane by the IBIS/ISGRI soft gamma-ray detector onboard the International Gamma Ray Laboratory.
 The source was the first and most extreme example of a number of highly absorbed Galactic X-ray binaries discovered withINTEGRAL., The source was the first and most extreme example of a number of highly absorbed Galactic X-ray binaries discovered with.
. Due to the strong absorption. which can exceed an equivalent hydrogen column of 107em7*. these sources are extremely faint in the soft X-rays and had not been detected by earlier missions2005).," Due to the strong absorption, which can exceed an equivalent hydrogen column of $10^{24}\,\mathrm{cm}^{-2}$, these sources are extremely faint in the soft X-rays and had not been detected by earlier missions."
" Right after its discovery. a re-analysis of archival data by revealed a highly photoabsorbed source (Vy=4x107 em"") coincident with the position given byINTEGRAL."," Right after its discovery, a re-analysis of archival data by revealed a highly photoabsorbed source $N_\mathrm{H} = 4 \times 10^{23}\,\mathrm{cm}^{-2}$ ) coincident with the position given by."
. The data also suggested an iron emission line at kKKeV. These results were confirmed by various subsequent studies 2003)., The data also suggested an iron emission line at keV. These results were confirmed by various subsequent studies .
 detected intense Fe Κα. Fe Kf. and Ni Ke emission lines in the spectrum.," detected intense Fe $\alpha$ , Fe $\beta$, and Ni $\mathrm{K\alpha}$ emission lines in the spectrum."
 Based on the interstellar absorption toward the system. which is two orders of magnitude lower than the measured Ny.(2003)..(2004).. and also suggested that much of the X-ray absorption is intrinsic to the compact object.," Based on the interstellar absorption toward the system, which is two orders of magnitude lower than the measured $N_\mathrm{H}$, and also suggested that much of the X-ray absorption is intrinsic to the compact object."
 In an optical study of the system. proposed that IGR 116318—4848 is a High Mass X-ray Binary (HMXB) with an sgB[e] star as the mass donor surrounded by a dense and absorbing circumstellar material2007).," In an optical study of the system, proposed that IGR $-$ 4848 is a High Mass X-ray Binary (HMXB) with an sgB[e] star as the mass donor surrounded by a dense and absorbing circumstellar material."
. This dense stellar wind results in significant photoabsorption within the binary system., This dense stellar wind results in significant photoabsorption within the binary system.
 Based on the optical data. suggest a distance between 0.9 and kkpe for the system.," Based on the optical data, suggest a distance between 0.9 and kpc for the system."
 A likely location for the source is in the Norma-Cygnus arm2004).. which would place it at a distance of kkpe2004).," A likely location for the source is in the Norma-Cygnus arm, which would place it at a distance of kpc."
 In thisPaper. we describe the results of follow-up observations of IGR 116318—4848 obtained with the satellite. the instruments on which are uniquely suited to study Compton-thick absorption.," In this, we describe the results of follow-up observations of IGR $-$ 4848 obtained with the satellite, the instruments on which are uniquely suited to study Compton-thick absorption."
 In Sect., In Sect.
 ?? we describe the data reduction., \ref{sec:data} we describe the data reduction.
 Section ?? is devoted to a presentation of the results of the spectral and temporal analysis., Section \ref{sec:obs} is devoted to a presentation of the results of the spectral and temporal analysis.
 We discuss our results in Sect. ?2.., We discuss our results in Sect. \ref{sec:conclusions}.
 We observed IGR J16318—4848 with from 2006 August ]4 until 2006 August 17 for a total net exposure of 97kks sequence number 401094010)., We observed IGR $-$ 4848 with from 2006 August 14 until 2006 August 17 for a total net exposure of ks sequence number 401094010).
 We used the standard. procedures to reduce the data from the X-Ray Imaging Spectrometer and the Hard X-Ray Detector2007)., We used the standard procedures to reduce the data from the X-Ray Imaging Spectrometer and the Hard X-Ray Detector.
.. For the XIS in particular we barycentered the data with (version 2008-03-03) and then extracted source events. images. spectra. and lighteurves with XSELECT v2.4.," For the XIS in particular we barycentered the data with (version 2008-03-03) and then extracted source events, images, spectra, and lightcurves with XSELECT v2.4."
 A circular source extractior region of 3/223 radius was applied., A circular source extraction region of 23 radius was applied.
 The background spectrum was extracted from a circular region having the same area as the source extraction region., The background spectrum was extracted from a circular region having the same area as the source extraction region.
 This process was done for every XIS., This process was done for every XIS.
 Response matrices and ancillary response files were generated using XISRMFGEN (version 2009-02-28) and XISSIMARFGEN (version 2009-02-28). taking into account the hydrocarbon contamination on the optical blocking filter2007).," Response matrices and ancillary response files were generated using XISRMFGEN (version 2009-02-28) and XISSIMARFGEN (version 2009-02-28), taking into account the hydrocarbon contamination on the optical blocking filter."
. As recommended by the team. the spectra of the three front illuminated CCDs (XISO. XIS2. and XIS3) were then combined with (version 1.30).," As recommended by the team, the spectra of the three front illuminated CCDs (XIS0, XIS2, and XIS3) were then combined with (version 1.30)."
 Although the XIS1 was operational when the observation was made. it is not used in the present study due to cross calibration Issues.," Although the XIS1 was operational when the observation was made, it is not used in the present study due to cross calibration issues."
 To extract the HXD PINspectrum. weagain followed the standard procedure of barycentric correction. eti-filtered," To extract the HXD PINspectrum, weagain followed the standard procedure of barycentric correction, gti-filtered"
spectrum as [lat as the observed. one.,spectrum as flat as the observed one.
 Therefore. an cllicient particle acceleration. mechanism is requested: to boost electrons toward. higher energies and to flatten the emitted svichrotron spectrum., Therefore an efficient particle acceleration mechanism is requested to boost electrons toward higher energies and to flatten the emitted synchrotron spectrum.
 In order to avoid to exceed the observed. brightness. a relatively small injection rate of secondary electrons. anc positrons i8. required.," In order to avoid to exceed the observed brightness, a relatively small injection rate of secondary electrons and positrons is required."
" More quantitatively, we find that the parameter space with C,f/Eyye&10.7 is excluded for this strongly magnetized case."," More quantitatively, we find that the parameter space with ${\cal E}_p/{\cal E}_{th} > 10^{-3}$ is excluded for this strongly magnetized case."
 In this Section we illustrate our calculations of the volume integrated. fluxes of radiation generated by. reaccelerated electrons and. positrons through svynchrotron emission. ancl 1C., In this Section we illustrate our calculations of the volume integrated fluxes of radiation generated by reaccelerated electrons and positrons through synchrotron emission and IC.
 The central gas density. ΕΞ 0). the 7 parameter and the core radius r7 are chosen as the representative values of the Coma cluster (Briel et al.," The central gas density, $n_{th}(r=0)$ , the $\beta$ parameter and the core radius $r_c$ are chosen as the representative values of the Coma cluster (Briel et al."
 1992)., 1992).
 In Fig., In Fig.
 5 we plot our results for the svnchrotron spectra (left panel) and the IC spectra (right panel)., \ref{fig:broad} we plot our results for the synchrotron spectra (left panel) and the IC spectra (right panel).
 The data points refer to the radio. hard X-ray and. gamma ray bands.," The data points refer to the radio, hard X-ray and gamma ray bands."
 All curves are obtained in the assumption that the cosmic rav energy density at the beginning of the reacceleration stage is proportional o the thermal energy density at any point., All curves are obtained in the assumption that the cosmic ray energy density at the beginning of the reacceleration stage is proportional to the thermal energy density at any point.
" The values of Pi(rΞ0) and the ratio £,/£u,e are not chosen to obtain a best fit to the data. they are only fixed in order to provide a viable representation of the data."," The values of $P_A(r=0)$ and the ratio ${\cal E}_p/{\cal E}_{\rm th}$ are not chosen to obtain a best fit to the data, they are only fixed in order to provide a viable representation of the data."
 Some general remarks emerge [rom the inspection. of lig. 5 , Some general remarks emerge from the inspection of Fig. \ref{fig:broad} :
The very broad extension of the svachrotron emission from. giant radio halos is among the properties which are cillicult to be fitted. by secondary models (o... Brunetti 2004 ancl rof.," The very broad extension of the synchrotron emission from giant radio halos is among the properties which are difficult to be fitted by secondary models (e.g., Brunetti 2004 and ref."
 therein)., therein).
 Although this Section is not devoted. to a detailedcomparison between observed. svnchrotron profiles, Although this Section is not devoted to a detailedcomparison between observed synchrotron profiles
galaxies. so the simulated dispersion is likely overestimated by à [actor of ~5.,"galaxies, so the simulated dispersion is likely overestimated by a factor of $\sim5$."
 While the data show a slight trend of σ increasing with ορ. the observed relation is inconsistent with simulations.," While the data show a slight trend of $\sigma$ increasing with $\Sigma_{SFR}$, the observed relation is inconsistent with simulations."
 We therefore conclude that. supernova feedback is insullicient to explain. the observed: velocity dispersions., We therefore conclude that supernova feedback is insufficient to explain the observed velocity dispersions.
 Figure 7 shows the relation between Vm Mere. clearly demonstrating that Vf decreases with sry.," Figure \ref{fig:sfrsigma} shows the relation between $V/\sigma$ $\Sigma_{SFR}$ , clearly demonstrating that $V/\sigma$ decreases with $\Sigma_{SFR}$."
 Llowever. this trend is likely ultimately due to the velocity-size relation: larger galaxies tend to have larger rotation velocities (Figure 6)) and smaller κ (Figure 7)).," However, this trend is likely ultimately due to the velocity-size relation: larger galaxies tend to have larger rotation velocities (Figure \ref{fig:dynamics}) ) and smaller $\Sigma_{SFR}$ (Figure \ref{fig:sfrsigma}) )."
 This is explained by dilferent. sensitivities of the data. as deeper spectra (lower ορ) reveal more extended: structures at larger radius with larger rotation speed.," This is explained by different sensitivities of the data, as deeper spectra (lower $\Sigma_{SFR}$ ) reveal more extended structures at larger radius with larger rotation speed."
 The velocity-size correlation contributes much more to the observed. Vo Nerr relation than any correlation between ο and the velocity. dispersion., The velocity-size correlation contributes much more to the observed $V/\sigma$ $\Sigma_{SFR}$ relation than any correlation between $\Sigma_{SFR}$ and the velocity dispersion.
 Fhese data thus do not support. the hypothesis that Ve is strongly allected by the density of star formation: e increases by less than a factor of 2 over two orders of magnitude in spp., These data thus do not support the hypothesis that $V/\sigma$ is strongly affected by the density of star formation: $\sigma$ increases by less than a factor of 2 over two orders of magnitude in $\Sigma_{SFR}$.
" The small Vi,fomLs and clumpy morphology of all ealaxies suggest that the rotating disks are highly turbulent and mav be dynamically unstable.", The small $V_{max}/\sigma \leq 1.8$ and clumpy morphology of all galaxies suggest that the rotating disks are highly turbulent and may be dynamically unstable.
 We therefore explore the scale lengths for gravitational collapse within the high redshift disk galaxies., We therefore explore the scale lengths for gravitational collapse within the high redshift disk galaxies.
 Evidence is accumulating that the mode of. star formation may be very. different in early systems compared to that seen locally (22)..," Evidence is accumulating that the mode of star formation may be very different in early systems compared to that seen locally \citep{Bournaud08,Elmegreen05}."
 Rather than forming stars within elant molecular clouds which condense out of a stable ealaxy. star formation mav be triggered by fragmentation of a dynamically unstable system.," Rather than forming stars within giant molecular clouds which condense out of a stable galaxy, star formation may be triggered by fragmentation of a dynamically unstable system."
" Driellv. in a rotating disk of gas ancl stars. perturbations smaller than a critical wavelength: £,,,; are stabilized. against the inward pull of gravity by velocity dispersion while those larger than some £L, are stabilized. by centrifugal force."," Briefly, in a rotating disk of gas and stars, perturbations smaller than a critical wavelength $L_{max}$ are stabilized against the inward pull of gravity by velocity dispersion while those larger than some $L_{min}$ are stabilized by centrifugal force."
 1£ the dispersion. and. rotation velocity are too low. LivinL4; and perturbations of intermediate wavelength grow exponentially.," If the dispersion and rotation velocity are too low, $L_{min} > L_{max}$ and perturbations of intermediate wavelength grow exponentially."
 This interplay is summarized by the Toomre parameter Qo—Linesτων which is calculated from the velocity dispersion. rotation curve. ancl mass. distribution (?)..," This interplay is summarized by the Toomre parameter $Q=L_{max}/L_{min}$ which is calculated from the velocity dispersion, rotation curve, and mass distribution \citep{Toomre64}."
" Galaxies with Q«I are therefore unstable on scales between £,,4; and L,,;, and will fragment into giant dense clumps.", Galaxies with $Q < 1$ are therefore unstable on scales between $L_{max}$ and $L_{min}$ and will fragment into giant dense clumps.
 This could trigger star formation in clouds of much higher mass and radius than GMCSs in local spiral galaxies with Q71. and can explain the clump-cluster and chain morphologies observed in many high-redshift galaxies.," This could trigger star formation in clouds of much higher mass and radius than GMCs in local spiral galaxies with $Q > 1$, and can explain the clump-cluster and chain morphologies observed in many high-redshift galaxies."
 Dynamical friction. viscosity anc tidal interactions may cause these clumps to migrate toward the center of the galaxy potential. forming a bulge which stabilizes the svstem against further fragmentation.," Dynamical friction, viscosity and tidal interactions may cause these clumps to migrate toward the center of the galaxy potential, forming a bulge which stabilizes the system against further fragmentation."
 From the galaxies whose velocity fields. can be reasonably well described by rotating svstems. we calculate the Toomre parameter via: which describes the stability of a rotating disk of gas.," From the galaxies whose velocity fields can be reasonably well described by rotating systems, we calculate the Toomre parameter via: which describes the stability of a rotating disk of gas."
 IQ-1 the system. is unstable to localgravitational collapse and will fragment into overdense clumps., If $Q<1$ the system is unstable to localgravitational collapse and will fragment into overdense clumps.
" The value of & is somewhat uncertain as it depends on the unknown mass distribution: our observations are consistent with a range ντf» oh25 correspondingm: to constant V,t and V.txJi respectively.", The value of $\kappa$ is somewhat uncertain as it depends on the unknown mass distribution; our observations are consistent with a range $\sqrt{2}\frac{V_c}{R}$ $2\frac{V_c}{R}$ corresponding to constant $V_c$ and $V_c\propto R$ respectively.
 Adopting &=νο1 appropriate for a uniform clisk and using dynamical mass to estimate the surface mass density. X. we find an inclination-corrected €x;0.6 for all galaxies in our sample.," Adopting $\kappa=\sqrt{3}V_{max}/R$ appropriate for a uniform disk and using dynamical mass to estimate the surface mass density $\Sigma$ , we find an inclination-corrected $Q\lsim0.6$ for all galaxies in our sample."
 We estimate that the uncertainty in € is dominated by à factor of 2 error in the dynamical mass., We estimate that the uncertainty in $Q$ is dominated by a factor of $\simeq2$ error in the dynamical mass.
 Phe assumed & introduces a negligible uncertainty. with an additional random error of ~30% [from the input. parameters.," The assumed $\kappa$ introduces a negligible uncertainty, with an additional random error of $\sim30$ from the input parameters."
 Disk thickness and stellar abundance also alfect the value of Q., Disk thickness and stellar abundance also affect the value of $Q$.
 Combined. these elfects result in roughly a factor of 2 uncertainty.," Combined, these effects result in roughly a factor of 2 uncertainty."
 Even so. these galaxies all appear to be dynamically unstable since Q<I.," Even so, these galaxies all appear to be dynamically unstable since $Q<1$."
 Hence we expect them to fragment into massive clumps on scales of order the Jeans length for dispersion support., Hence we expect them to fragment into massive clumps on scales of order the Jeans length for dispersion support.
 In a uniform disk. the largest scale for whieh velocity dispersion stabilizes against gravitational collapse is which is readily estimated from. the dispersion. and dynamical mass density.," In a uniform disk, the largest scale for which velocity dispersion stabilizes against gravitational collapse is which is readily estimated from the dispersion and dynamical mass density."
 As with Q. the uncertainty in Lj is à factor of &2 dominatedby the dynamical mass with additional uncertainty from the unknown mass cistribution. disk height. stellar content. and directional dependence of m.," As with $Q$, the uncertainty in $L_J$ is a factor of $\simeq2$ dominatedby the dynamical mass with additional uncertainty from the unknown mass distribution, disk height, stellar content, and directional dependence of $\sigma$ ."
 The resulting instability scale is 13 kpefor CI00024|1709 and 0.1.1.5 kpe for all other objects. consistent with the observed. clump sizes.," The resulting instability scale is 1–3 kpcfor 0024+1709 and 0.1–1.5 kpc for all other objects, consistent with the observed clump sizes."
"in their targeted observations of eSZ clusters, which resulted in detections with peak S/Ns ranging from 6.3 to 13.8 (Storyetal.2011).","in their targeted observations of eSZ clusters, which resulted in detections with peak S/Ns ranging from 6.3 to 13.8 \citep{story11}."
" 'These data were reduced according to the procedures described in detail in Sayersetal.(2011), with only minor changes."," These data were reduced according to the procedures described in detail in \citet{sayers11}, with only minor changes."
 We briefly summarize the data reduction here., We briefly summarize the data reduction here.
" First, we used frequent observations of bright quasars to obtain pointing corrections accurate to 5 arcsec, and we used observations of Uranus and Neptune to obtain flux calibration accurate to596."," First, we used frequent observations of bright quasars to obtain pointing corrections accurate to 5 arcsec, and we used observations of Uranus and Neptune to obtain flux calibration accurate to."
". In particular, the flux calibration has been updated based on recent WMAP results as described in detail in Sayersetal. (2012)."," In particular, the flux calibration has been updated based on recent WMAP results as described in detail in \citet{sayers12}."
. The FOV-average signal is then subtracted at each time sample and the data timestreams are high-pass filtered at 250 mHz in order to remove atmospheric brightness fluctuations., The FOV-average signal is then subtracted at each time sample and the data timestreams are high-pass filtered at 250 mHz in order to remove atmospheric brightness fluctuations.
" Consequently, the cluster images are high-pass filtered in a way which we characterize via simulation to obtain a signal transfer function for each cluster."," Consequently, the cluster images are high-pass filtered in a way which we characterize via simulation to obtain a signal transfer function for each cluster."
 We searched for clusters in these images using a matched filter according to the formalism described in Vanderlindeetal.(2010) (see also Haehnelt&TegmarkMarriage(1996);etal. (2011))).," We searched for clusters in these images using a matched filter according to the formalism described in \citet{vanderlinde10}
 (see also \citet{haehnelt96, herranz02a, herranz02b,
melin06, marriage11}) )."
" We (2002a,b);approximate the cluster (2006);SZ surface brightness profile S(0) as where 0 is an angular distance and 0, is the core radius.", We approximate the cluster SZ surface brightness profile $S(\theta)$ as where $\theta$ is an angular distance and $\theta_c$ is the core radius.
" We then form a filter according to where Wu) is the signal transfer function in Fourier space, S(u) is the cluster profile in Fourier space, B(u) is the point-spread function in Fourier space, is the map-noise power spectrum, u is an angular N(u)frequency, and 77-! denotes the Fourier transform back to map space."," We then form a filter according to where $\tilde{W}(u)$ is the signal transfer function in Fourier space, $\tilde{S}(u)$ is the cluster profile in Fourier space, $\tilde{B}(u)$ is the point-spread function in Fourier space, $\tilde{N}(u)$ is the map-noise power spectrum, $u$ is an angular frequency, and $\mathcal{FT}^{-1}$ denotes the Fourier transform back to map space."
 We filter our images in map-space by convolving them with after truncating at a radius of 4 arcmin.," We filter our images in map-space by convolving them with $\psi(\theta)$, after truncating $\psi(\theta)$ at a radius of 4 arcmin."
" This Φ(θ),truncation is required to v(0)prevent the filter from introducing noise in the central region of the images from the low-coverage outer regions.", This truncation is required to prevent the filter from introducing noise in the central region of the images from the low-coverage outer regions.
" Analogous to Vanderlindeetal. each image is filtered using 12 different filters, with 0, (2010),,varied from 0.25 to 3.00 arcmin in 0.25 arcmin increments (note that our PSF has a FWHM of ~1 arcmin), and we define our detection significance ¢ as the peak found in any of these filtered images within an apertureS/N of 4 arcmin radius centered on the Planck eSZ coordinates."," Analogous to \citet{vanderlinde10}, each image is filtered using 12 different filters, with $\theta_c$ varied from 0.25 to 3.00 arcmin in 0.25 arcmin increments (note that our PSF has a FWHM of $\simeq 1$ arcmin), and we define our detection significance $\zeta$ as the peak S/N found in any of these filtered images within an aperture of 4 arcmin radius centered on the Planck eSZ coordinates."
 This search radius was chosen because it corresponds to the upper limit on the Planck astrometry for all but the lowest redshift clusters (PlanckCollaboration2011)., This search radius was chosen because it corresponds to the upper limit on the Planck astrometry for all but the lowest redshift clusters \citep{planck_esz}.
. We find ¢=5.7 for PLCKESZ G115.71 (with θ.=0.25) and ¢=3.8 for PLCKESZ G189.84 (with θ.=0.25)., We find $\zeta = 5.7$ for PLCKESZ G115.71 (with $\theta_c = 0.25$ ) and $\zeta = 3.8$ for PLCKESZ G189.84 (with $\theta_c = 0.25$ ).
" The centroids of these maxima are given in Table 1,, and each is within 1 arcmin of the coordinates given in the Planck eSZ."," The centroids of these maxima are given in Table \ref{tab:one}, and each is within 1 arcmin of the coordinates given in the Planck eSZ."
 Figure 1 shows thumbnails of each target filtered with the filter that produces the maximum S/N. We characterize our detection significance via simulation by generating noise maps from the combination of jackknife realizations of our data and a model for the astronomical fluctuations due to the cosmic microwave background and point sources based on SPT measurements of the mm-wave power spectrum (Sayersetal.2011;Reichardt2011).," Figure \ref{fig:one} shows thumbnails of each target filtered with the filter that produces the maximum S/N. We characterize our detection significance via simulation by generating noise maps from the combination of jackknife realizations of our data and a model for the astronomical fluctuations due to the cosmic microwave background and point sources based on SPT measurements of the mm-wave power spectrum \citep{sayers11, reichardt11}."
". We generated 1000 such realizations for each target, along with the 14 other targets observed with Bolocam in October 2011 to similar depths."," We generated 1000 such realizations for each target, along with the 14 other targets observed with Bolocam in October 2011 to similar depths."
" We then searched for the peak S/N in each of these noise realizations, the distributions of which were statistically identical for all 16 targets."," We then searched for the peak S/N in each of these noise realizations, the distributions of which were statistically identical for all 16 targets."
" We therefore used all 16000 measurements to characterize the probability of measuring a given peak S/N, with the results shown in Figure 2.."," We therefore used all 16000 measurements to characterize the probability of measuring a given peak S/N, with the results shown in Figure \ref{fig:two}. ."
" We find the probability of a false detection of PLCKESZ G115.71 is 5.3x107°, and we therefore confirm it as a cluster."," We find the probability of a false detection of PLCKESZ G115.71 is $5.3 \times 10^{-5}$, and we therefore confirm it as a cluster."
" Note that PLCKESZ G115.71 has ¢=5.7, which exceeds the value of ¢ obtained in all 16000 of our noise realizations, we have therefore extrapolated our simulation results using a parametric fit."," Note that PLCKESZ G115.71 has $\zeta = 5.7$, which exceeds the value of $\zeta$ obtained in all 16000 of our noise realizations, we have therefore extrapolated our simulation results using a parametric fit."
" The probability of a false detection of PLCKESZ G189.84 is 0.027, and we therefore do not claim a detection of this candidate."," The probability of a false detection of PLCKESZ G189.84 is 0.027, and we therefore do not claim a detection of this candidate."
" At a given value of C, our false detection rate per given map area is approximately three times higher than the SPT false detection rate described in Vanderlindeetal."," At a given value of $\zeta$, our false detection rate per given map area is approximately three times higher than the SPT false detection rate described in \citet{vanderlinde10}."
 The false detection probability is similar when ¢(Bolocam)(2010)..~¢(SPT)+0.3.," The false detection probability is similar when $\zeta(\textrm{Bolocam}) \simeq 
\zeta(\textrm{SPT}) + 0.3$."
" Although SPT survey data are similar to our Bolocam data in many respects, there are subtle differences that contribute to this increase in our false detection probability."," Although SPT survey data are similar to our Bolocam data in many respects, there are subtle differences that contribute to this increase in our false detection probability."
" In particular, the Bolocam images have tapered coverage and consequently non-uniform noise properties."," In particular, the Bolocam images have tapered coverage and consequently non-uniform noise properties."
" This means that for a given map pixel our filter will introduce noise from both lower- and higher-coverage regions of the map, and the filter is therefore not as optimal as the filter used by SPT."," This means that for a given map pixel our filter will introduce noise from both lower- and higher-coverage regions of the map, and the filter is therefore not as optimal as the filter used by SPT."
 We have also analyzed the Wide-field Infrared Survey Explorer (WISE) public-release data (Wrightetal.2010) at the location of each eSZ candidate., We have also analyzed the Wide-field Infrared Survey Explorer (WISE) public-release data \citep{wright10} at the location of each eSZ candidate.
" The central wavelengths of the WISE bands, which areat 3.4, 4.6, 12, and 22 um, are too long to search for these candidates using a red sequence technique (Gladders&Yee 2000),,"," The central wavelengths of the WISE bands, which areat 3.4, 4.6, 12, and 22 $\mu$ m, are too long to search for these candidates using a red sequence technique \citep{gladders00}, ,"
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"Observations reveal that some faint filamentary structures are very likely associated with the high-velocity bbullets (Peng et al, in prep.)","observations reveal that some faint filamentary structures are very likely associated with the high-velocity bullets (Peng et al, in prep.)"
 reported by Zapataetal.(2009)., reported by \citet{Zapata2009}.
". In Figures 5-—7,, we show the line intensity ratio maps between the CO isotopologues and the different rotational transitions."," In Figures \ref{ratio3}- \ref{ratio2}, we show the line intensity ratio maps between the CO isotopologues and the different rotational transitions."
 From the image of the tto J-6-5ratio(Fig., From the image of the to $J=6-5$ ratio (Fig.
"Supperpanel), wecanseethespatialdistributiono pidrreguopauedififletleelv6locities."," \ref{ratio3} upper panel), we can see the spatial distribution of high optical depth regions at different velocities."
T southdenseridge f romOrionBN/K LtoOrionS , The north-south dense ridge from Orion BN/KL to Orion South is clearly seen at 9.
"This high optical depth ridge is similar to the filamentary structures of the NH3 emission (Wiseman&Ho,1996,1998),, but has some morphological differences in the north of the Orion BN/KL region."," This high optical depth ridge is similar to the filamentary structures of the $_{3}$ emission \citep{Wiseman1996,Wiseman1998}, but has some morphological differences in the north of the Orion BN/KL region."
" Besides, the straight shape of the Orion Bar is pronounced at sgz110-11s!,, where three high optical depth regions are seen at the two ends and center of the Bar at Visp=110 and 11s!,, respectively."," Besides, the straight shape of the Orion Bar is pronounced at 10–11, where three high optical depth regions are seen at the two ends and center of the Bar at 10 and 11, respectively."
" In the lower panel of Figure 5,, the J=6-5toJ=3-2ratiosaround10——11km sshow a gradient in the Orion Bar which goes in the direction from the Trapezium stars."," In the lower panel of Figure \ref{ratio3}, the $J=6-5$ to $J=3-2$ ratios around 10–11 show a gradient in the Orion Bar which goes in the direction from the Trapezium stars."
 It indicates that J-6-5isstronglyexcitedattheedgeo ftheBarbyU V photons fromtheT rapez , It indicates that $J=6-5$ is strongly excited at the edge of the Bar by UV photons from the Trapezium stars.
"Figure 6 shows the ratio between J=7-6andJ=6-Satthreedif ferentradialvelocitieso f5, 10, andl5km |."," Figure \ref{ratio1} shows the ratio between $J=7-6$ and $J=6-5$ at three different radial velocities of 5, 10, and 15."
 The three panels show clear variations across the integrated line emission., The three panels show clear variations across the integrated line emission.
" It is interesting to note that at the cloud velocity of 10 km s, there are high 7—6/6—5 ratios located very close to the position of the Trapezium stars toward Orion BN/KL."," It is interesting to note that at the cloud velocity of 10 km $^{-1}$, there are high 7–6/6–5 ratios located very close to the position of the Trapezium stars toward Orion BN/KL."
" These gradients are likely produced by those massive stars that heat the molecular cloud, causing the stronger fleghobmeaide outhisclearlyseenaf$ine"," These gradients are likely produced by those massive stars that heat the molecular cloud, causing the stronger $J=7-6$ line emission compared with the $J=6-5$ line."
stbradientisalsoseenintheOrionBarwhichgoesintheperpendiculard henorth— s, A gradient is also seen in the Orion Bar which goes in the perpendicular direction to the Trapezium stars (Fig.
"~!)), and some patches with higher J-7-6brightnesstemperaturesareseeninsidetheBarorbehindtheionizatior"," \ref{ratio1} at $V_{\rm LSR}=10$ ), and some patches with higher $J=7-6$ brightness temperatures are seen inside the Bar or behind the ionization front."
 gi," Some horizontal and vertical strips, artifacts from the OTF mapping, arealso seen in the Orion Bar and South regions."
ven the calibration errors of 10% for both J=6-SandJ=7-6., These artifacts can affect the intensity ratio by $\sim 14 \%$ given the calibration errors of $10 \%$ for both $J=6-5$ and $J=7-6$.
"However, somestructuresseeninthe J=7-6toJ=6-5ratiomaps, e.g., thedashed linesinFigure6,, areunlikelyduetothecalibrationerroror scanning strip sdi "," However, some structures seen in the $J=7-6$ to $J=6-5$ ratio maps, e.g., the dashed-lines in Figure \ref{ratio1}, are unlikely due to the calibration error or scanning strips during theobservation, and seem to be footprints of filaments or outflows."
, Some similar filament structures have been noticed in the dust continuum emission observed by \citet{Johnstone1999}. .
"In the ratio map of J23-210CO J=3-2showninF igure7, , aclearelongatedanddenseridgeisseeninthenorth southdirection, wherethe"," In the ratio map of $J=3-2$ to $J=3-2$ shown in Figure \ref{ratio2}, , a clear elongated and dense ridge is seen in the north-south direction, where the"
"In the ratio map of J23-210CO J=3-2showninF igure7, , aclearelongatedanddenseridgeisseeninthenorth southdirection, wherethe'"," In the ratio map of $J=3-2$ to $J=3-2$ shown in Figure \ref{ratio2}, , a clear elongated and dense ridge is seen in the north-south direction, where the"
"In the ratio map of J23-210CO J=3-2showninF igure7, , aclearelongatedanddenseridgeisseeninthenorth southdirection, wherethe'-"," In the ratio map of $J=3-2$ to $J=3-2$ shown in Figure \ref{ratio2}, , a clear elongated and dense ridge is seen in the north-south direction, where the"
"In the ratio map of J23-210CO J=3-2showninF igure7, , aclearelongatedanddenseridgeisseeninthenorth southdirection, wherethe'-?"," In the ratio map of $J=3-2$ to $J=3-2$ shown in Figure \ref{ratio2}, , a clear elongated and dense ridge is seen in the north-south direction, where the"
"In the ratio map of J23-210CO J=3-2showninF igure7, , aclearelongatedanddenseridgeisseeninthenorth southdirection, wherethe'-?C"," In the ratio map of $J=3-2$ to $J=3-2$ shown in Figure \ref{ratio2}, , a clear elongated and dense ridge is seen in the north-south direction, where the"
"In the ratio map of J23-210CO J=3-2showninF igure7, , aclearelongatedanddenseridgeisseeninthenorth southdirection, wherethe'-?CO"," In the ratio map of $J=3-2$ to $J=3-2$ shown in Figure \ref{ratio2}, , a clear elongated and dense ridge is seen in the north-south direction, where the"
we use f=0 (i.e. no overshooting). whereas for niasses equal to or ligher than 2.0M... we set f=0.018.,"we use $f=0$ (i.e. no overshooting), whereas for masses equal to or higher than $2.0\, M_\odot$ we set $f= 0.018$."
 In the intermediate mass range the cfitcicucy of overshooting varies lhuearlv with mass The πια]. discoutinuitv iu the overshooting efficiency is of no consequence for our study., In the intermediate mass range the efficiency of overshooting varies linearly with mass The small discontinuity in the overshooting efficiency is of no consequence for our study.
 1.1M. models with f£=0.005. as would result frou application of the above equation to this mass. lead. eiven the low overshooting efficieucv. to the same evolutionary tracks as models without overshooting.," $1.1\, M_\odot$ models with $f=0.005$, as would result from application of the above equation to this mass, lead, given the low overshooting efficiency, to the same evolutionary tracks as models without overshooting."
 This prescription for the overshooting eficiency leads to evolutionary tracks that reproduce well those by Pietiuferuietal.(2001)., This prescription for the overshooting efficiency leads to evolutionary tracks that reproduce well those by \citet{pietr:04}.
. Atomic diffusion is treated within the same diffusive uunercal scheme., Atomic diffusion is treated within the same diffusive numerical scheme.
 In this paper. hydrogen. helium aud heavicr clemeuts (Gucliding won) are diffusing.," In this paper, hydrogen, helium and heavier elements (including iron) are diffusing."
 Radiative levitation is not taken iuto account., Radiative levitation is not taken into account.
 To agree with the physical assuuptious in VCO. the default set of nuclear reaction rates are those of the NACRE collaboration (Auguloetal.1999).," To agree with the physical assumptions in VG07, the default set of nuclear reaction rates are those of the NACRE collaboration \citep{ang:99}."
 To investigate their influence. for individual reactions alternative rates were used.," To investigate their influence, for individual reactions alternative rates were used."
" For the UNGp,3350 reaction. we also used the rate from Maurtaetal.(2008).. the newest result from the LUNA collaboration. which is lower bv about a factor of two with respect to the NACRE rate at the relevant temperatures."," For the $^{14} \mathrm
N(p,\gamma)^{15} \mathrm O$ reaction, we also used the rate from \citet{mfn14:2008}, the newest result from the LUNA collaboration, which is lower by about a factor of two with respect to the NACRE rate at the relevant temperatures."
 This has a considerable effect on the TO morphology., This has a considerable effect on the TO morphology.
" Another reaction that turned out to be of surprisinely strong influence is the ITO(p,oHN yeaction."," Another reaction that turned out to be of surprisingly strong influence is the $^{17} \mathrm O(p,\alpha)^{14} \mathrm N$ reaction."
 We tested this by cluploving either the recent nieasureimeuts by Moazeuetal. (2007). which for T>τς105 K is very similar to the NACRE vrecouunendation. or the one given bv Caughlan&Fowler(1988).," We tested this by employing either the recent measurements by \citet{moaz:07}, which for $T > 4\times 10^8$ K is very similar to the NACRE recommendation, or the one given by \citet{cf:1988}."
.. A detailed discussion ou tle reaction rates can be found in$ L.2.., A detailed discussion on the reaction rates can be found in \ref{s:nucleo}.
 For all solar composition choices (ACSOS. CSOs) consistent Rosselaud mean opacity tables were prepared following the procedure described in Weiss&Schlattl (2008).," For all solar composition choices (AGS05, GS98) consistent Rosseland mean opacity tables were prepared following the procedure described in \citet{wsch:2008}."
. The final step to compare isochroues with observed CNIDs is the transformation to colors., The final step to compare isochrones with observed CMDs is the transformation to colors.
 For this we used hat bv VandenBere&Clem(2003.Χο)— which. ogether with a choice for the distance modulus and reddening of MGT. results in satisfring CMD fits on the nai sequence and subeiant brauch.," For this we used that by \citet[VC03]{vc:03}
 which, together with a choice for the distance modulus and reddening of M67, results in satisfying CMD fits on the main sequence and subgiant branch."
" Alternatively. we used the transformations by Cassisietal.(2001)/ for esting purposes,"," Alternatively, we used the transformations by \citet{cas:04} for testing purposes."
 The initial stellaz parameters (Yi. Zi. mixing length xmanieter) are obtained from solar model calibrations.," The initial stellar parameters $Y_\mathrm{in}$, $Z_\mathrm{in}$, mixing length parameter) are obtained from solar model calibrations."
 This will be discussed. in more detail in 3.2.., This will be discussed in more detail in \ref{s:calib}.
 With hese. we computed the evolution from the zero-age uaiu-sequeuce (ZAMS) to the tip of the Red Ciaut Brauch (ROB) for mass values from 0.6 to 1.5AL. in steps of 0.1AL...," With these, we computed the evolution from the zero-age main-sequence (ZAMS) to the tip of the Red Giant Branch (RGB) for mass values from $0.6$ to $1.5\ M_\odot$ in steps of $0.1\ M_\odot$."
 We recall here that. as in VGUOT. we are not interested im considering the ROB of M67. where deficiencies in our treatment of the outer lavers of the star could have au inpact on the location of the moclels in the IIRD.," We recall here that, as in VG07, we are not interested in considering the RGB of M67, where deficiencies in our treatment of the outer layers of the star could have an impact on the location of the models in the HRD."
 We confirmed by tests starting from the pre-inain sequence that for stars with mass below the critical mass for the occurrence of a convective core. Meee. the trausieut convective core at the eud of the pre-nian sequence phase also appears here.," We confirmed by tests starting from the pre-main sequence that for stars with mass below the critical mass for the occurrence of a convective core, $M_\mathrm{ccc}$, the transient convective core at the end of the pre-main sequence phase also appears here."
 This couvective core is the result of the short phase of CN-couversion. but may be sustained if couvective overshooting is iucluded.," This convective core is the result of the short phase of CN-conversion, but may be sustained if convective overshooting is included."
 Finally. for coustructing isochrones. the tracks are normalized to the so-called equivalent poiuts (Berebuscli&Vaudeuberg1992:Pietriuferuietal.2001) aud the interpolation to an isochroue is done between— the normalized tracks.," Finally, for constructing isochrones, the tracks are normalized to the so-called equivalent points \citep{vb:92,pietr:04}
 and the interpolation to an isochrone is done between the normalized tracks."
 After an isochrone age was fixed. we calculated an additional model with the TO-nass and recaleulated the isochrouc to make sure that the 'TOAanorpholosv does not depend on the isochrone interpolation sclicie.," After an isochrone age was fixed, we calculated an additional model with the TO-mass and recalculated the isochrone to make sure that the TO-morphology does not depend on the isochrone interpolation scheme."
 To compare our inodels with M67 we used the photometric data by Saudquist(2001)., To compare our models with M67 we used the photometric data by \citet{esm67:2004}.
.. These are more accurate than the older data from Aloutgomeryctal.(1993).. which were used by ΧιθἘν aud coutain bona fide suele stars only.," These are more accurate than the older data from \citet{mmm67:93}, , which were used by VG07, and contain bona fide single stars only."
 Thus. this CAID has a narrower 1iain-sequence baud.," Thus, this CMD has a narrower main-sequence band."
 However. we also tested some of our isochrones with the data by Montgomeryetal.19900).," However, we also tested some of our isochrones with the data by \citet{mmm67:93}."
 We used both (BG.V) aud (V.P) colors., We used both $(B-V)$ and $(V-I)$ colors.
 The dereddened distance modulus to AIG? is (nAM)p=9.70 according to VGOT. and 9.72+0.05 (Saudquist2001) based on subdwuf fitting to the lower imain-sequence following Percivaletal.(2003).," The dereddened distance modulus to M67 is $(m-M)_V = 9.70$ according to VG07, and $9.72\pm 0.05$ \citep{esm67:2004} based on subdwarf fitting to the lower main-sequence following \citet{psk:03}."
. The reddeniug is £(2BVy)=Οδ (WCF) iu eood agreenient with Sarajedinictal.(1999).. who eave (0.01+0.02.," The reddening is $E(B-V)=0.038$ (VG07) in good agreement with \citet{shkd:99}, who gave $0.04\pm0.02$."
 Simular values have been recently obtained by Twarogetal.(2009)., Similar values have been recently obtained by \citet{twarog:2009}.
. For the metallicity we assumed a value of ΓοΠΠ=0.0. which is well witlin spectroscopically determined errors. for example by Catton (20003.. who gave [FeTT]=0.02+0.06.," For the metallicity we assumed a value of $\mathrm{[Fe/H]}=0.0$, which is well within spectroscopically determined errors, for example by \citet{gm67:2000}, , who gave $\mathrm{[Fe/H]}=0.02\pm0.06$."
 To be cousistent with VOOT. we used for the old respectively new solar abundances those by CSO8 aud ACSü5. although the former are simply au update of GN93 taking iuto account additional literature values. and ACSO9 would be the most recent aud complete re-analysis.," To be consistent with VG07, we used for the old respectively new solar abundances those by GS98 and AGS05, although the former are simply an update of GN93 taking into account additional literature values, and AGS09 would be the most recent and complete re-analysis."
 Since ACGS09 abundances are slightly higher than ACSOS. our choice is testing the more extreme case.," Since AGS09 abundances are slightly higher than AGS05, our choice is testing the more extreme case."
 The initial abuudances of the stellar models for the M67 isochrones are taken to be the same as those resulting from solar model calibrations., The initial abundances of the stellar models for the M67 isochrones are taken to be the same as those resulting from solar model calibrations.
 The idea of WOOT to use MT for testing the effect of the new solar abundances rests ou the fact that the turn-off mass of this open cluster is very close to the critical mass for the onset of core convection. A4. or “transition mass” (VandeuBerectal.2007).," The idea of VG07 to use M67 for testing the effect of the new solar abundances rests on the fact that the turn-off mass of this open cluster is very close to the critical mass for the onset of core convection, $M_\mathrm{ccc}$, or ”transition mass” \citep{vgeef:2007}."
.. This transition mass depends on the average exponeut of the euecrgev ecueration rate in the core (Kippeuhlahnu&Weigert1990.Chap. 22).. that increases with the contribution from. the CNO-evele. for which eaxg~ T. in contrast to €upwT for the pp-chaius at the relevaut temperatures.," This transition mass depends on the average exponent of the energy generation rate in the core \citep[Chap.~22]{kw:90}, that increases with the contribution from the CNO-cycle, for which $\epsilon_\mathrm{CNO} \sim T^{15}$ , in contrast to $\epsilon_\mathrm{pp} \sim T^5$ for the pp-chains at the relevant temperatures."
 The importance of the CNO-cvcle depends on the nuclear reaction rates (see L2)) as well as on the amount of “catalysts”. tthe suu of the CNQO-abuudauces.," The importance of the CNO-cycle depends on the nuclear reaction rates (see \ref{s:nucleo}) ) as well as on the amount of “catalysts”, the sum of the CNO-abundances."
" These are. incidentally, the clemeuts with the largest reduction in their abundance according to AGSOS,"," These are, incidentally, the elements with the largest reduction in their abundance according to AGS05."
 Therefore. the morphological chanec in the CMD. ddisplaving the characteristichook at the TO. ora gap instardensity. cau be used to determine whether the TO-iass is below or above AM.," Therefore, the morphological change in the CMD, displaying the characteristichook at the TO, or a gap in stardensity, can be used to determine whether the TO-mass is below or above $M_\mathrm{ccc}$ ."
de ., .
. Here is the mass e is the fluid velocity. B is the magneticp field. is the density.gravitational acceleration. and ~ is the adiabatic index.," Here $\rho$ is the massdensity, $\bb{v}$ is the fluid velocity, $\bb{B}$ is the magnetic field, $\bb{g}$ is the gravitational acceleration, and $\gamma$ is the adiabatic index."
 gThe Lagrangian andEulerian derivatives are related via d/dt=OfOf|aV.," The Lagrangian andEulerian derivatives are related via $d/dt \equiv \pa/\pa t + \bb{v}
\bcdot \nabla$."
 (2))-(23)) describe the dynamics of a binary mixture in Equationsthe low-collisionality regime and they differ from standard MHD in three important respects: (0) The pressure tensor P.—pil|(pjpajbb. is where the symbols L and refer to the directions anisotropic:perpendicular and parallel to the magnetic| field. whose direction ts given by the versor 6=B/D. (ii))," Equations \ref{rho}) \ref{c}) ) describe the dynamics of a binary mixture in the low-collisionality regime and they differ from standard MHD in three important respects: ) The pressure tensor $\mathsf{P} \equiv p_\bot \mathsf{I} +
(p_\parallel - p_\bot) \hat{\bb{b}} \hat{\bb{b}}$, is anisotropic; where the symbols $\bot$ and $\parallel$ refer to the directions perpendicular and parallel to the magnetic field, whose direction is given by the versor $\hat{\bb{b}}\equiv \bb{B}/B$. )"
 Heat flows field lines. because the electron mean free mainly is alonglarge magneticcompared to its Larmor radius.," Heat flows mainly along magnetic field lines, because the electron mean free path is large compared to its Larmor radius."
 This process is pathmodeled by the second term on the hand side of (20)) via Q..εδ). rightwhere T is the plasma Equationtemperature. assumed to be the same for ions and electrons. and y.z6«10‘T°’? eres em1 + Lis the thermal conductivity (Spitzer1962:Bragin-skit 19635).. (iii))," This process is modeled by the second term on the right hand side of Equation \ref{S}) ) via $\bb{Q}_{\rm s}
\equiv-\chi\b(\b\bcdot\na)T$, where $T$ is the plasma temperature, assumed to be the same for ions and electrons, and $\chi \approx 6
\times 10^{-7} T^{5/2}$ ergs $^{-1}$ $^{-1}$ $^{-1}$ is the thermal conductivity \citep{1962pfig.book.....S, 1965RvPP....1..205B}. . )"
 The composition of fluid elements can change due to particle fluxes., The composition of fluid elements can change due to particle fluxes.
" Q,= Db(b-W)c on the right hand side of Equation Considering (23)) ensures that the diffusion of ions is field lines.", Considering $\bb{Q}_{\rm c}\equiv-D\b(\b\bcdot\na)c$ on the right hand side of Equation \ref{c}) ) ensures that the diffusion of ions is mainly along magnetic field lines.
 This is à good approximation mainlywhen alongthe magneticis dilute for the ion meanfree path to be large plasmacompared to the ion enoughLarmor radius., This is a good approximation when the plasma is dilute enough for the ion meanfree path to be large compared to the ion Larmor radius.
" Note that the concentration c is related to the mean molecular weight vial/jg—(1.011Zu)fpr,|60 where and Z;. with ¢=1.2. are the molecular weightsZo) fpr.and the ii;atomic numbers for the two ton "," Note that the concentration $c$ is related to the mean molecular weight via $1/\mu \equiv (1-c)(1+Z_1)/\mu_1 + c(1+Z_2)/\mu_2$, where $\mu_i$ and $Z_i$ , with $i=1,2$, are the molecular weights and the atomic numbers for the two ion species."
The of the pressure tensor is thus P?=2p.species./3|p)/3isotropicphypartpm. where &p is the Boltzmann constant and 0 is the atomic mass unit.," The isotropic partof the pressure tensor is thus $P\equiv 2p_\bot/3 + p_\parallel/3 = \rho k_{\rm B}
T/\mu m_{\rm H}$, where $k_{\rm B}$ is the Boltzmann constant and $m_{\rm H}$ is the atomic mass unit."
 We assume a plane-parallel in a constant gravitational field g=—yz which is atmospherestratified in. both and the vertical direction., We assume a plane-parallel atmosphere in a constant gravitational field $\bb{g}\equiv-g\hat{\bb{z}}$ which is stratified in both temperature and composition along the vertical direction.
 We temperatureconsider a backgroundcompositionalong field which is weak that the mechanical magnetic of the atmosphere. with enough /7. is maintained via equilibriumοον=," We consider a backgroundmagnetic field which is weak enough that the mechanical equilibrium of the atmosphere, with scaleheight $H$, is maintained via $dP/dz=-g\rho$ ."
 In general. the scaleheightbackground heat and particle fluxes do not vanish.αρ. Le.. bVT=0 and &Ve=0. unless the magnetic field and the gradients are ," In general, the background heat and particle fluxes do not vanish, i.e., $\hat{\bb{b}}\bcdot \na T \ne 0$ and $\hat{\bb{b}}\bcdot \na c \ne 0$, unless the magnetic field and the background gradients are orthogonal."
"The existence of a well defined backgroundsteady state. Le. orthogonal.V«Q.=VQ,0. demands that the gradients should be linear functions of the distance the backgrounddirection of the field."," The existence of a well defined steady state, i.e., $\na\bcdot\bb{Q}_{\rm s}=\na\bcdot\bb{Q}_{\rm c}=0$, demands that the background gradients should be linear functions of the distance along the direction of the magnetic field."
 However. even If this alongcondition is not magneticsatisfied. the dynamics of the modes that we consider is strictlyunlikely to be significantly affected if the local dynamical timescale is short compared to the timescale in which the entire system evolves (see also Quataert2008)).," However, even if this condition is not strictly satisfied, the dynamics of the modes that we consider is unlikely to be significantly affected if the local dynamical timescale is short compared to the timescale in which the entire system evolves (see also \citealt{2008ApJ...673..758Q}) )."
 The modes of interest have associated timescales that are long compared to the sound crossing time and it thus suffices to work in the Boussinesg approximation., The modes of interest have associated timescales that are long compared to the sound crossing time and it thus suffices to work in the Boussinesq approximation.
 In this limit. the equations for the linear perturbations 6—cf!eX become .. Αιde. | ," In this limit, the equations for the linear perturbations $\delta \sim e^{\sigma t + i \bb{k}\bcdot\bb{x}}$ become ), v_z + ) ."
"In agreement with the Boussinesq approximation, the velocity perturbations satisfy μυ=0 and the fluctuations in density. temperature. and mean molecular weight are related via"," In agreement with the Boussinesq approximation, the velocity perturbations satisfy $\bb{k}\bcdot \delta \bb{v} = 0$ and the fluctuations in density, temperature, and mean molecular weight are related via."
" Here. we have introduced the Alfvénn speed. v4 B/\/Tzp. the thermal speed. ey,= \/2P/p. the plasma $= ο the viscosity i of the binary mixture. the thermal diffusion coefficient «=κ. ο. and the Váusállá frequencyN? ~= -- I"," Here, we have introduced the Alfvénn speed, $\bb{v}_{\rm A}\equiv \bb{B}/\sqrt{4\pi\rho}$ , the thermal speed, $v_{\rm th}\equiv \sqrt{2P/\rho}$ , the plasma $\beta\equiv v_{\rm th}^2/v_{\rm A}^2$ , the viscosity $\nu$ of the binary mixture, the thermal diffusion coefficient $\kappa\equiv \chi T/P$ , and the $-$ Väiisällä frequencyN^2 = g ."
O For we define here several that play an important completeness.role inthe stability analysis.," For completeness,we define here severalquantities that play an important role inthe stability analysis."
 We quantitiesdenote, We denote
IMF tend to have lower values.,IMF tend to have lower values.
 This difference is induced by the larger metal content associated to the top-heavy IMF. which makes cooling more efficient within halos near the resolution limit.," This difference is induced by the larger metal content associated to the top–heavy IMF, which makes cooling more efficient within halos near the resolution limit."
 Finally. we show in Figure 8 the effect of increasing the feedback efficiency on the LF.," Finally, we show in Figure \ref{fi:lf_sw} the effect of increasing the feedback efficiency on the LF."
 In this case. we use the same normalization for the two IMFs. in order to directly see the effect of changing the feedback strength.," In this case, we use the same normalization for the two IMFs, in order to directly see the effect of changing the feedback strength."
 Quite interestingly. the effect is that of suppressing the bright end of the LF. while leaving the faint end almost unaffected.," Quite interestingly, the effect is that of suppressing the bright end of the LF, while leaving the faint end almost unaffected."
 A number of observations have established that the galaxy population in clusters is characterized by the presence of color gradients. with bluer galaxies preferentially avoiding to reside in the innermost cluster regions (2)..," A number of observations have established that the galaxy population in clusters is characterized by the presence of color gradients, with bluer galaxies preferentially avoiding to reside in the innermost cluster regions \citep{1984ApJ...285..426B}."
 For instance. ?.— found a decreasingtrend ofthe ./? color with cluster-centric distance for the galaxies lying on the CMR of nearby optically selected clusters.," For instance, \cite{2006MNRAS.366..645P} found a decreasingtrend of the $B-R$ color with cluster-centric distance for the galaxies lying on the CMR of nearby optically selected clusters."
 Similar results have also been found by 2.. ? and ? for moderately distant X-ray selected clusters.," Similar results have also been found by \cite{1996ApJ...471..694A}, \cite{1997ApJ...478..462C} and \cite{2005ApJ...627..186W} for moderately distant X–ray selected clusters."
 Quite consistently. outer cluster regions are populated by a larger fraction of blue galaxies (e.g. ?).. thus confirming. that more external galaxies are generally," Quite consistently, outer cluster regions are populated by a larger fraction of blue galaxies \citep[e.g.,
][]{2004MNRAS.351..125D}, , thus confirming that more external galaxies are generally"
gas.,gas.
" Like the 8 o'clock arc, all three are relatively massive (M.~101 ΜΜΟ) and they also show dynamics that suggest rotational support rather than merger-driven star formation (which may be more dominant in lower-mass UV-luminous systems)."," Like the 8 o'clock arc, all three are relatively massive $_\ast\sim10^{11}$ $_\odot$ ) and they also show dynamics that suggest rotational support rather than merger-driven star formation (which may be more dominant in lower-mass UV-luminous systems)."
 A cold-accreting system of this kind is likely to display the same discrepancy between metallicities determined from the youngest stars and from other methods., A cold-accreting system of this kind is likely to display the same discrepancy between metallicities determined from the youngest stars and from other methods.
" Cold-mode accretion has been hypothesised as a dominant galaxy formation mechanism at the highest redshifts (Dekeletal. 2009),, although this hypothesis remains controversial (seee.g.Steideletal.2010).."," Cold-mode accretion has been hypothesised as a dominant galaxy formation mechanism at the highest redshifts \citep{2009Natur.457..451D}, , although this hypothesis remains controversial \citep[see e.g.][]{2010ApJ...717..289S}."
" While such a case has not been proven in the 8 o’clock arc, it may prove an interesting target for detailed analysis with integral field spectroscopy, to determine whether there is any evidence for this scenario."," While such a case has not been proven in the 8 o'clock arc, it may prove an interesting target for detailed analysis with integral field spectroscopy, to determine whether there is any evidence for this scenario."
" The remaining three galaxies differ to those already discussed in that they do not require the simultaneous invocation of low metallicity, reduced carbon abundance and QHE in their synthetic stellar population to explain the spectral features under consideration, but rather display some or none of these features."," The remaining three galaxies differ to those already discussed in that they do not require the simultaneous invocation of low metallicity, reduced carbon abundance and QHE in their synthetic stellar population to explain the spectral features under consideration, but rather display some or none of these features."
 The z=3.1 Cosmic Eye spectrum (Quideretal. is illustrated in Figure 9 and is notably devoid of strong emission., The $z=3.1$ Cosmic Eye spectrum \citep{quid1} is illustrated in Figure \ref{eye} and is notably devoid of strong emission.
 This can be explained if the galaxy lies above the metallicity limit at which QHE no longer occurs., This can be explained if the galaxy lies above the metallicity limit at which QHE no longer occurs.
" The absorption section of the P Cygni profile also implies a relatively strong metallicity somewhere between Z—0.004 and 0.008, with a reduced carbon abundance required at the higher metallicity mass fraction."," The absorption section of the P Cygni profile also implies a relatively strong metallicity somewhere between $Z=0.004$ and $0.008$, with a reduced carbon abundance required at the higher metallicity mass fraction."
 Higher metallicities are ruled out by the relatively shallow absorption., Higher metallicities are ruled out by the relatively shallow absorption.
" The remaining two examples, cB58 al.2002) and the Cosmic Horseshoe (z=2.4,Quider 2009),, in Figures 10 and 11 respectively, tell a similarstory."," The remaining two examples, cB58 \citep[$z=2.7$,][]{cb58paper} and the Cosmic Horseshoe \citep[$z=2.4$,][]{horse}, in Figures \ref{cb58} and \ref{horse} respectively, tell a similarstory."
" The lack of a strong line, yet apparent P-Cygni"," The lack of a strong line, yet apparent P-Cygni"
"large fraction of pathological objects and spectra of low SNR, have redshifts derived via manual inspection of the spectra.","large fraction of pathological objects and spectra of low SNR, have redshifts derived via manual inspection of the spectra."
" Independent spectrum classifications and redshift determinations, based on direct ?-fitting of template spectra to the data, have been made at Princeton using thecode?."," Independent spectrum classifications and redshift determinations, based on direct $\chi^2$ -fitting of template spectra to the data, have been made at Princeton using the."
". The redshift determination, essentially via cross-correlation, differs from the implementation employed in thespectrold pipeline but the same composite quasar template from was used."," The redshift determination, essentially via cross-correlation, differs from the implementation employed in the pipeline but the same composite quasar template from was used."
 Fig., Fig.
 1 shows a comparison of the SDSS final-redshifts and Princeton redshifts as a function of quasarredshift?., \ref{sdss_prince} shows a comparison of the SDSS final-redshifts and Princeton redshifts as a function of quasar.
". The selection of the sub-sample of more than 700000 spectra is conservative in that only spectra with high-confidence SDSS redshifts, where there is also no inconsistency between the cross-correlation and emission line redshift determinations, are used."," The selection of the sub-sample of more than 000 spectra is conservative in that only spectra with high-confidence SDSS redshifts, where there is also no inconsistency between the cross-correlation and emission line redshift determinations, are used."
 The data in Fig., The data in Fig.
" 1 should essentially represent an internal consistency check and the large differences between redshifts, extending to +5x10~°, or!,, are surprising."," \ref{sdss_prince} should essentially represent an internal consistency check and the large differences between redshifts, extending to $\pm$ $\times$ $^{-3}$, or, are surprising."
 Perhaps even more striking is the sequence of apparent discontinuities in the behaviour as a function of redshift., Perhaps even more striking is the sequence of apparent discontinuities in the behaviour as a function of redshift.
" A second illustration of the extent of redshift-dependent systematics comes from comparing the redshift derived from the location of the AA2796,2803 emission in each quasar spectrum with the SDSS redshift."," A second illustration of the extent of redshift-dependent systematics comes from comparing the redshift derived from the location of the $\lambda\lambda$ 2796,2803 emission in each quasar spectrum with the SDSS redshift."
 Fig., Fig.
 2 presents the data for more than 600000 spectra with SNR>10 emission line locations (from the SDSS spectroscopic pipeline’))., \ref{sdss_mgii} presents the data for more than 000 spectra with $\ge$ 10 emission line locations (from the SDSS spectroscopic ).
" The rest-frame location of the emission line has been shown by many studies over the decades to be well-behaved and there is no reason to expect 7-500 !shifts over small redshift intervals, or, indeed, an apparent systematic 2x 10? s!) change in the location of the emission with increasing redshift of the quasars."," The rest-frame location of the emission line has been shown by many studies over the decades to be well-behaved and there is no reason to expect $\simeq$ shifts over small redshift intervals, or, indeed, an apparent systematic $\times$ $^{-3}$ ) change in the location of the emission with increasing redshift of the quasars."
 The systematic redshift differences show similar patterns over the redshift range common to both Fig., The systematic redshift differences show similar patterns over the redshift range common to both Fig.
 1 and Fig. 2.., \ref{sdss_prince} and Fig. \ref{sdss_mgii}.
" Although somewhat more complex to interpret (Section ??)), the equivalent plot for the emission, Fig. 3,,"," Although somewhat more complex to interpret (Section \ref{sec:em_shifts}) ), the equivalent plot for the emission, Fig. \ref{sdss_ciii},"
 also shows strong systematic effects as a function of redshift., also shows strong systematic effects as a function of redshift.
 The form and substantial amplitude of the systematic and random differences in Figs. 1-, The form and substantial amplitude of the systematic and random differences in Figs. \ref{sdss_prince}-
-3 led to the initiation of the investigation presented here.," \ref{sdss_ciii}
 led to the initiation of the investigation presented here."
 The generation of the high-SNR quasar template to be used to calculate cross-correlation redshifts begins with a sample of quasars at low redshifts that possess emission line-determined redshifts., The generation of the high-SNR quasar template to be used to calculate cross-correlation redshifts begins with a sample of quasars at low redshifts that possess emission line-determined redshifts.
 A somewhat more involved procedure is then necessary to incorporate additional quasars at higher redshifts into the master template., A somewhat more involved procedure is then necessary to incorporate additional quasars at higher redshifts into the master template.
 In this section the recipe for each element of the master template construction are outlined., In this section the recipe for each element of the master template construction are outlined.
" The narrow forbidden emission lines of A44960,5008 are prominent in many quasar spectra with redshifts z<0.8 and a composite spectrum based on the combination of quasars with redshifts determined via the location of emission forms the starting point for the construction of the master quasar template."," The narrow forbidden emission lines of $\lambda\lambda$ 4960,5008 are prominent in many quasar spectra with redshifts $z$$<$ 0.8 and a composite spectrum based on the combination of quasars with redshifts determined via the location of emission forms the starting point for the construction of the master quasar template."
In simulations we follow the procedure used by. Bednarz Ostrowski (1996) with a hvbrid approach used in Dednarz Ostrowski (1908).,In simulations we follow the procedure used by Bednarz Ostrowski (1996) with a hybrid approach used in Bednarz Ostrowski (1998).
 Mlonoenergetic seed particles are injected at the shock and then their trajectories are derived in the perturbed magnetic field., Monoenergetic seed particles are injected at the shock and then their trajectories are derived in the perturbed magnetic field.
 Phe inhomogeneities are simulated bv small amplitude particle momentum scattering within a cone with angular opening AY less than the particle anisotropy L/5 (cL., The inhomogeneities are simulated by small amplitude particle momentum scattering within a cone with angular opening $\Delta \vartheta$ less than the particle anisotropy $\sim 1/\gamma$ (cf.
 Ostrowski 1991)., Ostrowski 1991).
 A particle is excluded: from. simulations if dt escapes through the frec-escape boundary. placed far olf the shock or reaches the energy larger than the assumed. upper limit., A particle is excluded from simulations if it escapes through the free-escape boundary placed far off the shock or reaches the energy larger than the assumed upper limit.
 These particles are replaced with the ones arising from splitting the remaining high-weight particles. preserving their physical parameters.," These particles are replaced with the ones arising from splitting the remaining high-weight particles, preserving their physical parameters."
 Particles that exist longer than the time upper limit for simulations are excluded: from. simulations without replacing., Particles that exist longer than the time upper limit for simulations are excluded from simulations without replacing.
 All computations are performed. in the respective upstream or downstream plasma rest frame., All computations are performed in the respective upstream or downstream plasma rest frame.
 When particles cross the shock their parameters are. transformed. to the current plasma rest frame and the weighted: contribution divided by the particle velocity component normal to the shock (= particle density) is. added: to the time anc momentum bin depending on particle parameters. as measured in the shock normal rest framoe.," When particles cross the shock their parameters are transformed to the current plasma rest frame and the weighted contribution divided by the particle velocity component normal to the shock $\equiv$ particle density) is added to the time and momentum bin depending on particle parameters, as measured in the shock normal rest frame."
 For the considered continuous injection after initial time. the energy. cut-olf of the formed spectrum shifts toward higher energies with time.," For the considered continuous injection after initial time, the energy cut-off of the formed spectrum shifts toward higher energies with time."
 The resulting spectra allows one to fit spectral indices aud derive acceleration time in the shock normal rest. frame in units of downstream ryfe (ry - particle gvroradius in the homogeneous magnetic field component. e - speed of light: for details see Bednarz Ostrowski 1996).," The resulting spectra allows one to fit spectral indices and derive acceleration time in the shock normal rest frame in units of downstream $r_{g}/c$ ( $r_{g}$ - particle gyroradius in the homogeneous magnetic field component, $c$ - speed of light; for details see Bednarz Ostrowski 1996)."
 We transform the acceleration. time foe to the downstream: plasma rest. frame., We transform the acceleration time $t_{acc}$ to the downstream plasma rest frame.
 Llerealter. subscripts U or D mean that a parameter is measured in the upstream or downstream plasma rest frame respectively.," Hereafter, subscripts U or D mean that a parameter is measured in the upstream or downstream plasma rest frame respectively."
" We will use downstream ry, as a distance and ryὁ as à time units.", We will use downstream $r_{g}$ as a distance and $r_{g}/c$ as a time units.
 The magnetic Ποιά inclination to the shock normal upstream of the shock. c. is measured. in the upstream plasma rest [rame.," The magnetic field inclination to the shock normal upstream of the shock, $\psi$, is measured in the upstream plasma rest frame."
 Let us denote the ratio of the cross-lield diffusion covllicient #y to the parallel. diffusion cocllicicnt &q as T (the value is measured in the plasma rest. frame)., Let us denote the ratio of the cross-field diffusion coefficient $\kappa_\perp$ to the parallel diffusion coefficient $\kappa_\|$ as $\tau$ (the value is measured in the plasma rest frame).
 Simulations prove that Uuctuations upstream of the shock (measured by 7!) and downstream of the shock (measured wor?) influence the acceleration process. independently., Simulations prove that fluctuations upstream of the shock (measured by $\tau^{U}$ ) and downstream of the shock (measured by $\tau^{D}$ ) influence the acceleration process independently.
 The minimum fluctuations upstream of the shock needed o run the acceleration process cllicienthy tend to zero when ~ox., The minimum fluctuations upstream of the shock needed to run the acceleration process efficiently tend to zero when $\gamma \rightarrow \infty$.
 We. checked by simulations. with. different- 7?) that its value does not influence the spectral index considerably Or à given Tp, We checked by simulations with different $\tau^{D}$ that its value does not influence the spectral index considerably for a given $\tau^{U}$.
 Our scattering model is very simple but also universal., Our scattering model is very simple but also universal.
 In the model we are not able to cliscuss gxroresonat scattering., In the model we are not able to discuss gyroresonat scattering.
 Upstream of the shock particles have not enough time to interact resonantly with low-frequeney waves and the rough relation 7(6B/B)! (ef., Upstream of the shock particles have not enough time to interact resonantly with low-frequency waves and the rough relation $\tau \sim (\delta B/B)^{4}$ (cf.
 Blandford Eichler 1987) cannot be deduce from the interaction there., Blandford Eichler 1987) cannot be deduce from the interaction there.
 However. for erowing τί and fixed AW the time between scattering acts decreases what is equivalent to increasing the magnetic field fluctuations.," However, for growing $\tau^{U}$ and fixed $\Delta \vartheta$ the time between scattering acts decreases what is equivalent to increasing the magnetic field fluctuations."
 In the [following simulations we consider shocks with ~= 20. 40. SO. 160. 320. magnetic field inclinations wos1ο.45°. 607.757.90 and downstream: values of magnetic field Uuctuations 7?=O.L0-107.1.1 0.11.0.69.," In the following simulations we consider shocks with $\gamma=$ 20, 40, 80, 160, 320, magnetic field inclinations $\psi=15^\circ, 30^\circ,
45^\circ$, $60^\circ, 75^\circ, 90^\circ$ and downstream values of magnetic field fluctuations $\tau^{D}=0, 1.0\cdot 10^{-3}, 1.1\cdot 10^{-2}, 0.11, 0.69$ ."
 ‘Thus. as a first case we consider downstream conditions without magnetic field fluctuations.," Thus, as a first case we consider downstream conditions without magnetic field fluctuations."
 By simple data inspection (cf., By simple data inspection (cf.
 Fig., Fig.
 2) we look for minimum 7 where the, 2) we look for minimum $\tau^{U}$ where the
Disks evolve on time-scales that are orders of magnitudes smaller than expected from microphysical transport processes. and various suggestions have been made over the vears to explain this discrepancy.,"Disks evolve on time-scales that are orders of magnitudes smaller than expected from microphysical transport processes, and various suggestions have been made over the years to explain this discrepancy."
 Turbulent transport. in particular. has figured among the leading candidates since the inception of the a-disk paradigm. and a number of hydrodynamic and MHD turbulent transport mechanisms have been proposed in the literature.," Turbulent transport, in particular, has figured among the leading candidates since the inception of the $\alpha$ -disk paradigm, and a number of hydrodynamic and MHD turbulent transport mechanisms have been proposed in the literature."
 On the hydrodynamic side. subcritical turbulence (? and references therein). if present. is apparently too inefficient (??)..," On the hydrodynamic side, subcritical turbulence \citealt{RZ99}
 and references therein), if present, is apparently too inefficient \citep{LL05,JBSG06}."
 Convection was up to now found too inetficient and to transport angular momentum inthe wrong direction (??).. but ἃ recent reinvestigation of the problem indicates that this might be an artifact of these simulations being performed too close to the stability threshold (2)..," Convection was up to now found too inefficient and to transport angular momentum inthe wrong direction \citep{C96,SB96}, but a recent reinvestigation of the problem indicates that this might be an artifact of these simulations being performed too close to the stability threshold \citep{LO10}."
 Two-dimensional weak turbulence driven by small-scale. incoherent gravitational instabilities (density waves) is an option (?)..," Two-dimensional weak turbulence driven by small-scale, incoherent gravitational instabilities (density waves) is an option \citep{G96}."
 Alternatively. the baroclinic instability (?) may generate vorticity. and transport through the coupling with density waves. but its conditions of existence are still controversial (??).. although ? have probably identified the root of this debate by pointing out the nonlinear nature of the instability: also the resulting vortices would be subject to 3D instabilities (?)..," Alternatively, the baroclinic instability \citep{KB03} may generate vorticity, and transport through the coupling with density waves, but its conditions of existence are still controversial \citep{JG06,PSJ07}, although \cite{LP10} have probably identified the root of this debate by pointing out the nonlinear nature of the instability; also the resulting vortices would be subject to 3D instabilities \citep{LP09}."
 ? have proposed that the magnetorotational instability (MRI) is a potentially efficient source of turbulent transport in the nonlinear regime. an expectation soon borne out in numerical simulations.," \cite{BH91a} have proposed that the magnetorotational instability (MRI) is a potentially efficient source of turbulent transport in the nonlinear regime, an expectation soon borne out in numerical simulations."
 This instability provides by now the most extensively studied transport mechanism. through local unstratified (?).. stratified (2).. and global (2?) 3D disk simulations.," This instability provides by now the most extensively studied transport mechanism, through local unstratified \citep{HGB95}, stratified \citep{SHGB96}, and global \citep{H00}
 3D disk simulations."
 These initial simulations as well as the numerous ones following them have shown that MRI turbulence is an efficient way to transport angular momentum. in the presence or absence of a mean vertical or toroidal field. with an overall transport efficiency depending on the field configuration and strength.," These initial simulations as well as the numerous ones following them have shown that MRI turbulence is an efficient way to transport angular momentum, in the presence or absence of a mean vertical or toroidal field, with an overall transport efficiency depending on the field configuration and strength."
 However. the significant role played by microphysical dissipation in the resolutions accessible to date had largely been underestimated (??)..," However, the significant role played by microphysical dissipation in the resolutions accessible to date had largely been underestimated \citep{LL07,FPLH07}."
 By now. both the field strength and dissipation dependence of the simulated turbulent transport have been studied to some extent (and only in unstratified local shearing box settings for the latter one).," By now, both the field strength and dissipation dependence of the simulated turbulent transport have been studied to some extent (and only in unstratified local shearing box settings for the latter one)."
" The dependence of the Shakura-Sunyaev « parameter has been characterized very early on by ?. who showed that momentum transport o:7!7 both fora net vertical or toroidal (albeit with very different efficiencies in the two configurations), a scaling further confirmed in later simulations. as summarized in ?.."," The dependence of the Shakura-Sunyaev $\alpha$ parameter has been characterized very early on by \cite{HGB95} who showed that momentum transport $\propto \beta^{-1/2}$ both for a net vertical or toroidal (albeit with very different efficiencies in the two configurations), a scaling further confirmed in later simulations, as summarized in \cite{PCP07}."
 Until recently. the effect of physical viscosity (v) and resistivity Gp on the transport had been neglected. under the implicit assumption that these should not matter too much once inertial turbulent scales are resolved in the simulations.," Until recently, the effect of physical viscosity $\nu$ ) and resistivity $\eta$ ) on the transport had been neglected, under the implicit assumption that these should not matter too much once inertial turbulent scales are resolved in the simulations."
 However. ? have shown that. in the presence of à mean vertical field. the MRI-driven turbulent transport did exhibit a substantial dependence on the magnetic Prandtl number Pri v/n. with no clear trends with respect to either viscosity of resistivityalone.," However, \cite{LL07} have shown that, in the presence of a mean vertical field, the MRI-driven turbulent transport did exhibit a substantial dependence on the magnetic Prandtl number $Pm=\nu/\eta$ , with no clear trends with respect to either viscosity of resistivity."
. Recently. ? found similar results in shearing boxes with a mean toroidal field instead of a mean vertical one.," Recently, \cite{SH09} found similar results in shearing boxes with a mean toroidal field instead of a mean vertical one."
 When the mean magnetic flux vanishes. the transport behavior is more complex.," When the mean magnetic flux vanishes, the transport behavior is more complex."
 The initial investigation by ? concluded that the transport was converging to a finite value. but ? found that the transport efficiency was dependent on the simulation. resolution.," The initial investigation by \cite{HGB96} concluded that the transport was converging to a finite value, but \cite{GS05} found that the transport efficiency was dependent on the simulation resolution."
 More recently. the role of the magnetic Prandtl number Pri has been identified in this setting (2): turbulence exists only for magnetic Prandtl numbers larger than about 2. which requires the explicit inclusion of viscous and resistive terms in the fluid equations for numerical simulations to correctly capture the physies of the problem.," More recently, the role of the magnetic Prandtl number $Pm$ has been identified in this setting \citep{FPLH07}: turbulence exists only for magnetic Prandtl numbers larger than about 2, which requires the explicit inclusion of viscous and resistive terms in the fluid equations for numerical simulations to correctly capture the physics of the problem."
 The disappearance of turbulence at low Pri. as well as the need of large enough amplitudes in the initial conditions at Par> 2. indicate that the zero net flux magnetized shearing box is a subcritical system rather than a linearly unstable one (??)..," The disappearance of turbulence at low $Pm$, as well as the need of large enough amplitudes in the initial conditions at $Pm > 2$ , indicate that the zero net flux magnetized shearing box is a subcritical system rather than a linearly unstable one \citep{LO08b,LO08a}. ."
observations is presented in Table 1..,observations is presented in Table \ref{tabobs}.
 Spectra of the stellar templates 61 Cvg X B (IN5V and INTV) were also observed during the 2008 run for the purpose of radial velocities and rotational broadening analysis., Spectra of the stellar templates 61 Cyg A B (K5V and K7V) were also observed during the 2008 run for the purpose of radial velocities and rotational broadening analysis.
 In addition. eight templates of spectral tvpes G6-MOV. were collected during the 2003 campaign and another KSVY in 1999.," In addition, eight templates of spectral types G6-M0V were collected during the 2003 campaign and another K5V in 1999."
 The images were bias corrected and Hat-fielded. and the spectra subsequently. extracted. using conventional optimal extraction techniques in order to optimize the signal-to-noise ratio of the output (Llorne1986).," The images were bias corrected and flat-fielded, and the spectra subsequently extracted using conventional optimal extraction techniques in order to optimize the signal-to-noise ratio of the output \citep{Horne86}."
. Every target was bracketed with observations of a comparison CuAdr|CuNe are lamp and the pixel-to-wavelongth scale was derived through polynomial fits to a large number οἱ identified reference lines., Every target was bracketed with observations of a comparison CuAr+CuNe arc lamp and the pixel-to-wavelength scale was derived through polynomial fits to a large number of identified reference lines.
 The final rms scatter of the fit was always «1/30 of the spectral dispersion., The final rms scatter of the fit was always $<$ 1/30 of the spectral dispersion.
 We rectified the 27 individual spectra by subtracting à low-order spline fit to the continuum. after masking out the main emission and atmospheric absorption lines.," We rectified the 27 individual spectra by subtracting a low-order spline fit to the continuum, after masking out the main emission and atmospheric absorption lines."
 Phe spectra were subsequently rebinned into a uniform velocity scale of 36 kin 5 pix.1., The spectra were subsequently rebinned into a uniform velocity scale of 36 km $^{-1}$ $^{-1}$.
 Phe IK5V template 61 Cve A was broadened to 78 uns+ to match the width of the donor photospheric lines (sce Sect., The K5V template 61 Cyg A was broadened to 78 km $^{-1}$ to match the width of the donor photospheric lines (see Sect.
 4)., 4).
 Every spectrum of SY Cne was then cross-correlated: against the broadened. template in the spectral regions free [rom emission and telluric absorption features., Every spectrum of SY Cnc was then cross-correlated against the broadened template in the spectral regions free from emission and telluric absorption features.
 tadial velocities. were extracted. following the method. of lTonry&Davis(1979).. where parabolic fitshi were performed o the peak of the cross-correlation functions. and the uncertainties are purely statistical.," Radial velocities were extracted following the method of \cite{ton79}, where parabolic fits were performed to the peak of the cross-correlation functions, and the uncertainties are purely statistical."
 Since the orbital period is poorly constrained to 0.380. +0.001 d. we. performed a power spectrum analysis on the radial velocities in the range V.1-2 davs. and the results are. clisplavec in Fig. l..," Since the orbital period is poorly constrained to 0.380 $\pm0.001$ d we performed a power spectrum analysis on the radial velocities in the range 0.1-2 days, and the results are displayed in Fig. \ref{figperiod}. ."
 Here we have rescaled the errorbars by a factor 1.3 so that. the minimum v; is 1.0., Here we have rescaled the errorbars by a factor 1.3 so that the minimum $\chi^{2}_{\nu}$ is 1.0.
 The periodogram is dominated by strong aliasing due to the sparse sampling of our observations., The periodogram is dominated by strong aliasing due to the sparse sampling of our observations.
 In order to test the significance of the cillerent peaks above the noise level we performed. a Monte Carlo simulation., In order to test the significance of the different peaks above the noise level we performed a Monte Carlo simulation.
 Synthetic 47. periodograms were computed. from. a large population (10°) of velocities randomly picked from a white-noise distribution ancl with identical time sampling as our data., Synthetic $\chi^{2}$ periodograms were computed from a large population $^5$ ) of velocities randomly picked from a white-noise distribution and with identical time sampling as our data.
 A te significance level is defined by the 99.99. per cent of the computed X7. values and this is indicated. in lig., A $\sigma$ significance level is defined by the 99.99 per cent of the computed $\chi^{2}_\nu$ values and this is indicated in Fig.
 1 bv a horizontal dashed. line., \ref{figperiod} by a horizontal dashed line.
 Most. of the peaks between. [requencies 2.5-2.7  are under the line and hence periods in the range 0.37-0.40 days are significantIv above noise at the 99.99. per cent level., Most of the peaks between frequencies 2.5-2.7 $^{-1}$ are under the line and hence periods in the range 0.37-0.40 days are significantly above noise at the 99.99 per cent level.
 The two deeper peaks correspond to 0.3824 cay ancl 0.3837. day and. have v»2—L0 and 2.3 respectively. forB 24 degrees of ⋅⋅freedom., The two deeper peaks correspond to 0.3824 day and 0.3837 day and have $\chi^{2}_{\nu}=1.0$ and 2.3 respectively for 24 degrees of freedom.
 X deo significance level around the minimm peak will exclude the second. peak (Lamptonetal.1976). and hence we can conclude that 0.3824 dav is by far the most significant period., A $\sigma$ significance level around the minimm peak will exclude the second peak \citep{lampton76} and hence we can conclude that 0.3824 day is by far the most significant period.
 A least-squares sinewave fit to the radial velocities. using 22=(0.3824 day as input parameter. vielcls the following parameters where 75 corresponds to the Heliocentrie Julian date of the inferior conjunction of the donor star.," A least-squares sine-wave fit to the radial velocities, using $P=0.3824$ day as input parameter, yields the following parameters where $T_{0}$ corresponds to the Heliocentric Julian date of the inferior conjunction of the donor star."
 ALL quoted. errors are 68 per cent confidence., All quoted errors are 68 per cent confidence.
 “Lhe svstemic velocity 5. has been corrected fromthe radial velocity of 61 €vg A. that we take as -64.3 20.9 kms. 3 (Wilson1953).," The systemic velocity $\gamma$ has been corrected fromthe radial velocity of 61 Cyg A, that we take as -64.3 $\pm$ 0.9 km $^{-1}$ \citep{wilson53}."
.. Note that a sinewave fit fixing P=(0.3837 clay also vields A»=ὅτιE33 km 1. an indication⋠⋠⋠ that our Av»⊳ value is. robust and its. error realistic. irrespectively of the true value of the orbital period.," Note that a sinewave fit fixing $P=0.3837$ day also yields $K_2=87.3\pm3.3$ km $^{-1}$, an indication that our $K_2$ value is robust and its error realistic, irrespectively of the true value of the orbital period."
 The same is found when other peaks around the minimum are taken., The same is found when other peaks around the minimum are taken.
 mEe., Fig.
 2 displays the radial velocity. points folded on our favoured orbital period. £2=0.3823753 clay together with the best sine fitsolution., \ref{figrv} displays the radial velocity points folded on our favoured orbital period $P=0.3823753$ day together with the best sine fitsolution.
 Our A» velocity disagrees with the one reported. by Smithetal.(2005). which was obtained using the Na doublet., Our $K_2$ velocity disagrees with the one reported by \cite{smith05} which was obtained using the Na doublet.
 We note that the Na doublet can be quenched— by heating elfects. as opposed. to the metallic lines usec by us (Alartinetal. 1989)..," We note that the Na doublet can be quenched by heating effects, as opposed to the metallic lines used by us \citep{martin89}. ."
 LIÉ this were the case. the light centre of the Na lines would be displaced towards the back side of the star.," If this were the case, the light centre of the Na lines would be displaced towards the back side of the star."
 The effects of irradiation can be estimated. using the Ix-correction approach of Wade&Lorne(1988). Advefive=οαλά|q). where Adv is the increment in A» velocity clue to irradiation. re the radius of the donor star. @ the binary separation anc f the fractional displacement of the absorption lines site with respect to the center of mass of the star.," The effects of irradiation can be estimated using the K-correction approach of \cite{wadehorne88} $\Delta
K_2/K_2= f{\bf (}r_2/a{\bf )} (1+q)$, where $\Delta K_2$ is the increment in $K_2$ velocity due to irradiation, $r_2$ the radius of the donor star, $a$ the binary separation and $f$ the fractional displacement of the absorption line's site with respect to the center of mass of the star."
 In the extreme case. when the Na absorption is completely supressed from the irracliatecl hemisphere. f.=4/(32).," In the extreme case, when the Na absorption is completely supressed from the irradiated hemisphere, $f=4/(3\pi)$."
 Thus. replacing τω by I5eeleton's equation (Eeeleton1983) and adopting q=1.18 (sec next section) we find NA»=32 kms 1.," Thus, replacing $r_2/a$ by Eggleton's equation \citep{eggleton83} and adopting $q=1.18$ (see next section) we find $\Delta 
K_2=32$ km $^{-1}$."
 This cüllerence is just enough to accommodate Smith ct als value with ours. for maxiniuun quenching of the Na1 lines.," This difference is just enough to accommodate Smith et al's value with ours, for maximum quenching of the Na lines."
 Pherefore. we conclude that Smith ct als A» is likely a significant overestimate due to irradiation while our value is much less allected by these ellects.," Therefore, we conclude that Smith et al's $K_2$ is likely a significant overestimate due to irradiation while our value is much less affected by these effects."
 In order to measure the rotational broadening of the donor's absorption features we have [focused on the 10 highest resolution spectra gathered during the 2003 campaign., In order to measure the rotational broadening of the donor's absorption features we have focused on the 10 highest resolution spectra gathered during the 2003 campaign.
 We broadenedthe G6-MÓ templates [roni 5 to 100 km + in μαeps of 5 km using a Gray profile (Gray1992)and continuum limb-darkening coellicients appropriate for every μα»ectral type ancl our wavelength range.," We broadenedthe G6-M0 templates from 5 to 100 km $^{-1}$ in steps of 5 km $^{-1}$ , using a Gray profile \citep{Gray92} and continuum limb-darkening coefficients appropriate for every spectral type and our wavelength range."
 The broadened templates were multiplied by factors f« 1. to account Lor 16 fractional contribution to the total light.," The broadened templates were multiplied by factors $f<1$ , to account for the fractional contribution to the total light."
 These were, These were
process that clears out the disk to produce laree planets.,process that clears out the disk to produce large planets.
" Three parameters — M;/X.. py. and p, — provide good measures of the transition from oligarchy to chaos."," Three parameters – $\Sigma_l / \Sigma_s$ , $p_H$ , and $p_o$ – provide good measures of the transition from oligarchy to chaos."
" The Hill parameter measures when the oligarchs have enough mass {ο interact dwnanucally,", The Hill parameter measures when the oligarchs have enough mass to interact dynamically.
" The ratio X;/X, isolates (he Gime when planetesimals cannot camp the oligarchs and thus prevent large-scale dvnamical interactions.", The ratio $\Sigma_l / \Sigma_s$ isolates the time when planetesimals cannot damp the oligarchs and thus prevent large-scale dynamical interactions.
 The orbit overlap parameter distinguishes times when orbit overlap is important., The orbit overlap parameter distinguishes times when orbit overlap is important.
 To understand the transition from oligarchy {ο chaos in less idealized situations. we now consider complete planet formation simulations using the full hybrid code.," To understand the transition from oligarchy to chaos in less idealized situations, we now consider complete planet formation simulations using the full hybrid code."
 The caleulations start wilh 110 km planetesimals and allow all objects to collide. merge. and interact gravitationallv.," The calculations start with 1–10 km planetesimals and allow all objects to collide, merge, and interact gravitationally."
 When objects in the coagulation code reach mzz2x107(“y/8gem7) g. we promote them into the v-body code and follow their individual trajectories.," When objects in the coagulation code reach $m \approx 2 \times 10^{25}~(\Sigma_0 / {\rm 8~g~cm^{-2}})$ g, we promote them into the $n$ -body code and follow their individual trajectories."
 We describe caleulations in a small (large) torus in 823.3 (833.4)., We describe calculations in a small (large) torus in 3.3 3.4).
 The calculations begin with ὁ km planetesimals in a torus extending from 0.86 AU to 1.14 AU., The calculations begin with 1–3 km planetesimals in a torus extending from 0.86 AU to 1.14 AU.
 We divide this region into 32 annuli ancl seed each annulus with planetesimals in nearly circular and coplanar orbits (ey=10? and 4j= e4/2)., We divide this region into 32 annuli and seed each annulus with planetesimals in nearly circular and coplanar orbits $e_0 = 10^{-5}$ and $i_0 = e_0/2$ ).
 The planetesimals have surface densitv X=μίαAU)τι with X; = 116 g 7 at 1 AU.," The planetesimals have surface density $\Sigma = \Sigma_0 (a / {\rm 1 ~ AU})^{-3/2}$, with $\Sigma_0$ = 1–16 g $^{-2}$ at 1 AU."
 In these caleulations. we do nol consider lragmentation. which generally speeds up Che growth of the largest objects αἱ (he expense of mass loss [rom disruptions and gas drag (Wetherill&Stewart1993:IXenvon&Luu 1993).," In these calculations, we do not consider fragmentation, which generally speeds up the growth of the largest objects at the expense of mass loss from disruptions and gas drag \citep{ws93,kl98}."
. Weidenschillingetal.(1997) consider a similar suite of calculations.," \citet{wei97}
 consider a similar suite of calculations."
 Where il is possible to compare. our results agree with these calculations 2002 )..," Where it is possible to compare, our results agree with these calculations \citep[see also][]{kom02}."
 For Xj — 8 g 7. growth at 1 AU follows a standard pattern (Wetherill&Stewartetal.1997:IXenvon&Luu 1998).," For $\Sigma_0$ = 8 g $^{-2}$, growth at 1 AU follows a standard pattern \citep{ws93, wei97, kl98}."
. After a few thousand vears. mergers produce a few large objects with radii of ~ 10 km.," After a few thousand years, mergers produce a few large objects with radii of $\sim$ 10 km."
 As dynamical frietion. circularizes the orbits of these objects. runaway growth beeins.," As dynamical friction circularizes the orbits of these objects, runaway growth begins."
 It takes only 104 vr to produce several dozen 100300 km objects., It takes only $10^4$ yr to produce several dozen 100–300 km objects.
 At ~2xI0! vr. the first object is promoted into the »-body code.," At $\sim 2 \times 10^4$ yr, the first object is promoted into the $n$ -body code."
 As larger5 objects form farther out in the disk. more promotions occur.," As larger objects form farther out in the disk, more promotions occur."
" These objects continue to 5grow rapidly until they reach isolation! masses of ~1079 5ο, when stirring 55begins to reduce eravitational focusing factors."," These objects continue to grow rapidly until they reach `isolation' masses of $\sim 10^{26}$ g, when stirring begins to reduce gravitational focusing factors."
 The transition to oligarchic growth: beeins at the inner edge of the erid and rapidly propagates outwards., The transition to oligarchic growth begins at the inner edge of the grid and rapidly propagates outwards.
" At ~3xLO?vr. the number of oligarchs with masses m.2 107 e peaks al Ni,e T."," At $\sim$$3 \times 10^5$yr, the number of oligarchs with masses $m \gtrsim$ $10^{26}$ g peaks at $N_o \sim$ 7."
 Soon after oligarchic growth begins at the outer edge of the grid. oligarchs," Soon after oligarchic growth begins at the outer edge of the grid, oligarchs"
Part of this increase is undoubtedly due to the increase in noise in measuring the surface brightness at the faint outer wings of the stars.,Part of this increase is undoubtedly due to the increase in noise in measuring the surface brightness at the faint outer wings of the stars.
 However. we assume conservatively that the increase is entirely due to PSF variability and represent it by a parabolic eXpression where csi) 1s the uncertainty in the PSF relative to the intensity at radius r (pixels) from the PSF center.," However, we assume conservatively that the increase is entirely due to PSF variability and represent it 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."
 We then compute the expected variance of a pixel 7 with intensity / (ADU) and distance r from the center of from the expression where G is the effective gain and Ris 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.," We then compute the expected variance of a pixel $\sigma^2$ 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."
 In addition to the i-band images. two spectra with an exposure time of 900 s were obtained using grism #44 of ALFOSC. which covers the wavelength range 3200-9100A.," In addition to the i-band images, two spectra with an exposure time of 900 s were obtained using grism 4 of ALFOSC, which covers the wavelength range 3200-9100."
. The spectral resolution achieved using the 1733 slit wasÁ., The spectral resolution achieved using the 3 slit was.
. Due to the effect of the secondary spectrum and fringing longwards of 6000Á.. we studied only the wavelength range 4000-6000À.," Due to the effect of the secondary spectrum and fringing longwards of 6000, we studied only the wavelength range 4000-6000."
. The summed spectrum has a signal-to-noise ratio (S/N) of ~ 250. but the spectrum remains featureless.," The summed spectrum has a signal-to-noise ratio (S/N) of $\sim$ 250, but the spectrum remains featureless."
 Table | indicates the results of the model fits. and the upper panel of Fig.," Table \ref{tulokset} indicates the results of the model fits, and the upper panel of Fig."
 2. shows the one-dimensional surface brightness profile of and the model profiles., \ref{prof} shows the one-dimensional surface brightness profile of and the model profiles.
 The I-band magnitude of the nucleus translates into R ~ 14.2 using I) = 0.5(22).. so the nucleus had brightened by ~ 0.6 mag in the R-band from the deep minimum on Dec 17th.," The I-band magnitude of the nucleus translates into R $\sim$ 14.2 using (R-I) = 0.5, so the nucleus had brightened by $\sim$ 0.6 mag in the R-band from the deep minimum on Dec 17th."
 The host galaxy is faint (I = 17.5) compared to the nucleus and never exceeds the core brightness at any radius. but is visible as à small excess in the surface brightness profile of (see Fig. 2)).," The host galaxy is faint (I = 17.5) compared to the nucleus and never exceeds the core brightness at any radius, but is visible as a small excess in the surface brightness profile of (see Fig. \ref{prof}) )."
 Since PSFI was constructed from stars closer to than PSF2. it is unsurprising that PSFI provides a more accurate description of the data (smaller y7) than PSF2.," Since PSF1 was constructed from stars closer to than PSF2, it is unsurprising that PSF1 provides a more accurate description of the data (smaller $\chi^2$ ) than PSF2."
 The fit with core + galaxy reduces the y of the fit. but none of the fits provide formally satisfactory fits. most likely due to higher-order PSF variability unaccounted for in our radial PSF variability model.," The fit with core + galaxy reduces the $\chi^2$ of the fit, but none of the fits provide formally satisfactory fits, most likely due to higher-order PSF variability unaccounted for in our radial PSF variability model."
 We completed several tests to ensure that the observed excess Is real and not due to PSF variability over the field of view., We completed several tests to ensure that the observed excess is real and not due to PSF variability over the field of view.
 Firstly. we compared the sealed surface brightness profiles of stars 2. 3. 6. SI. and S2 and with the surface brightness profile of star 5 (see the lower panel of Fig.," Firstly, we compared the scaled surface brightness profiles of stars 2, 3, 6, S1, and S2 and with the surface brightness profile of star 5 (see the lower panel of Fig."
 2)., 2).
 There is a clear excess around that exceeds the rms scatter of stellar profiles by a factor of 3-6 at r>27, There is a clear excess around that exceeds the rms scatter of stellar profiles by a factor of 3-6 at $r >$.
"%, The excess is clearly above any PSF variability observed over the field of view.", The excess is clearly above any PSF variability observed over the field of view.
 Secondly. although the fits with PSF2 provide a slightly worse fit than with PSFI. the host galaxy component is still present with similar brightness and effective radius as with PSFI.," Secondly, although the fits with PSF2 provide a slightly worse fit than with PSF1, the host galaxy component is still present with similar brightness and effective radius as with PSF1."
 Also this test indicates that PSF variability does not affect the results significantly., Also this test indicates that PSF variability does not affect the results significantly.
 Finally. we fitted core + host galaxy models to stars 6. S2. and S4. which are similar in brightness to0716+714.," Finally, we fitted core + host galaxy models to stars 6, S2, and S4, which are similar in brightness to."
. Using PSFI. the fit converged towards a solution ry;—O in all three cases. 1.9. no host galaxy component was detected in targets known to be unresolved.," Using PSF1, the fit converged towards a solution $r_{\rm
 eff} \rightarrow 0$ in all three cases, i.e. no host galaxy component was detected in targets known to be unresolved."
 Based on the three test mentioned above. we conclude that the excess around is real and we have detected the host galaxy.," Based on the three test mentioned above, we conclude that the excess around is real and we have detected the host galaxy."
"the line-of-sight i.e., considering the imposing dimensions of the large region (e.g., Lagrois&Joncas 2009a,b)), the interface between compressive shocks and the surrounding ISM could be several times smaller than the size of the nebula itself.","the line-of-sight i.e., considering the imposing dimensions of the large region (e.g., \citealt{Lag2009a,Lag2009b}) ), the interface between compressive shocks and the surrounding ISM could be several times smaller than the size of the nebula itself."
" This said, the expected spectral signature of shock excitation could be strongly diluted by photoionized foreground/background material."," This said, the expected spectral signature of shock excitation could be strongly diluted by photoionized foreground/background material."
" A total of 23378 out of 30057 emission-line profiles retained for this study are directly associated to the bright, central structure clearly visible in Figure 8."," A total of 378 out of 057 emission-line profiles retained for this study are directly associated to the bright, central structure clearly visible in Figure 8."
" The strong emission of its ionized component (see peak intensities in Figures 3 to 7) and therefore the quality of the gathered signal make this structure, surrounded by the most massive stars of the Melotte 15 cluster, the ideal feature in our quest for shock excitation in central 11805."," The strong emission of its ionized component (see peak intensities in Figures 3 to 7) and therefore the quality of the gathered signal make this structure, surrounded by the most massive stars of the Melotte 15 cluster, the ideal feature in our quest for shock excitation in central 1805."
" First, a series of four weaker portions [in gas emission], specifically surrounding the central structure, were selected."," First, a series of four weaker portions [in gas emission], specifically surrounding the central structure, were selected."
 Each of these portions was spatially binned into a single 1x 11x 2249 x x cchannels)profile?.., Each of these portions was spatially binned into a single $\times$ $\times$ 249 $\times$ $\times$ channels).
" For each of the 23378 spectra selected and investigated in this subsection, the corresponding foreground/background emission was approximately recovered using a linear combination of these four 1x x 2249 profiles."," For each of the 378 spectra selected and investigated in this subsection, the corresponding foreground/background emission was approximately recovered using a linear combination of these four $\times$ $\times$ 249 profiles."
" Each of the four linear coefficients was weighted via the inverse of the distance separating the targeted pixel from the center of the corresponding weaker portion (e.g., the larger this distance is, the smaller is the corresponding linear coefficient and, therefore, the smaller is the statistical impact of this weaker zone on the computation of the foreground/background spectral signature at the position of this given pixel)."," Each of the four linear coefficients was weighted via the inverse of the distance separating the targeted pixel from the center of the corresponding weaker portion (e.g., the larger this distance is, the smaller is the corresponding linear coefficient and, therefore, the smaller is the statistical impact of this weaker zone on the computation of the foreground/background spectral signature at the position of this given pixel)."
" Hence, a unique foreground/background spectrum is constructed for each of the 23378 points (see below) although neighbor pixels have [as it should be] foreground/background material with very similar spectral signatures (i.e., roughly identical linear coefficients)."," Hence, a unique foreground/background spectrum is constructed for each of the 378 points (see below) although neighbor pixels have [as it should be] foreground/background material with very similar spectral signatures (i.e., roughly identical linear coefficients)."
" The subtraction of the foreground/background emission obviously had a certain incidence on the S/N of the resulting profiles, the peak intensity of all five lines being reduced in the subtracting process."," The subtraction of the foreground/background emission obviously had a certain incidence on the S/N of the resulting profiles, the peak intensity of all five lines being reduced in the subtracting process."
" Still following all conditions listed in 8 4.2.2, 22229 emission-line profiles out of 23378 were said to “survive” the procedure."," Still following all conditions listed in $\S$ 4.2.2, 229 emission-line profiles out of 378 were said to “survive” the procedure."
 Panels (a) and (b) of Figures, Panels (a) and (b) of Figures
Not all azimuthal components mm contribute al the same level.,Not all azimuthal components $m$ contribute at the same level.
" Figure 1. shows the amplitudes of the ;1,, as a function of m Dor two different values of the tube radius. &."," Figure \ref{fig.exact_and_born} shows the amplitudes of the $A_m$ as a function of $m$ for two different values of the tube radius, $R$ ."
 In this particular example. pj=5x10* ces. ec=11 km/s. and B=1 kG are fixed.," In this particular example, $\rho_0=5\times10^{-7}$ cgs, $c=11$ km/s, and $B=1$ kG are fixed."
 For IH—0.5 Mm. whieh is less than the wavelength (AJ?=0.36). only the m.—0.41 azimuthal components contribute.," For $R=0.5$ Mm, which is less than the wavelength $kR=0.86$ ), only the $m=0,\pm1$ azimuthal components contribute."
 For tubes with larger radii the higher order components have larger amplitudes. as can be seen in the case 22=2 Mm (hit=3.43).," For tubes with larger radii the higher order components have larger amplitudes, as can be seen in the case $R=2$ Mm $kR = 3.43$ )."
 Magnetic effects cause perturbations to both the steady. background state ancl the wavelfield., Magnetic effects cause perturbations to both the steady background state and the wavefield.
 In this section we use the Born approximation to derive an approximate solution {ο equations (11))-(15)) based on tlie assumption that magnetic effects are small 1995)., In this section we use the Born approximation to derive an approximate solution to equations \ref{cont}) \ref{temp}) ) based on the assumption that magnetic effects are small \citep[see e.g.][]{Rosenthal1995}.
. The Lorentz force is quadratic in the magnetic field., The Lorentz force is quadratic in the magnetic field.
 As à result we introduce a small parameter that is second order in (he magnetic field., As a result we introduce a small parameter that is second order in the magnetic field.
" We choose to expand all quantities in powers of the small dimensionless parameter In this expansion framework the magnetic field appears al order εντ,", We choose to expand all quantities in powers of the small dimensionless parameter In this expansion framework the magnetic field appears at order $\epsilon^{1/2}$.
 In particular. we write the steady background magnetic field as The magnetic field causes a shift. ey. in the steady component of the density inside the tube relative to the steady. component of the density outside the tube. py: Likewise. we write the steady. component of pressure as Thesechanges are related to (he magnetic field through equation (10)) aud (hie equationof state:," In particular, we write the steady background magnetic field as The magnetic field causes a shift, $\epsilon\rho_1$, in the steady component of the density inside the tube relative to the steady component of the density outside the tube, $\rho_0$ : Likewise, we write the steady component of pressure as Thesechanges are related to the magnetic field through equation \ref{eq.pressure_bal}) ) and the equationof state:"
"quadrant Q«c(«c 57. 5""«fOF. which comprises Daade's Window. is so far only partially released by 2\LASS.","quadrant $0^{\circ}<\ell<5^{\circ}$ , $-5^{\circ}<{\it b}<0^{\circ}$, which comprises Baade's Window, is so far only partially released by 2MASS."
 Fig., Fig.
" 4 shows a histogram of values for the cells in the 10"".10 extinction map.", 4 shows a histogram of values for the cells in the $^{\circ} \times 10^{\circ}$ extinction map.
 The mean extinction in the entire map ds <ly>=0.29 with a standard deviation o=0.12 from the mean., The mean extinction in the entire map is ${\it<A_K>} = 0.29$ with a standard deviation $\sigma = 0.12$ from the mean.
 63 of the cells fall within 2-0 of this mean value. ancl of them have τν1.0.," 63 of the cells fall within $\sigma$ of this mean value, and of them have ${\it A_K} < 1.0$."
 The upper panel of Fie., The upper panel of Fig.
 5 shows the histogram of internal crrors in extinction determination., 5 shows the histogram of internal errors in extinction determination.
 The mean internal error is «0;“=0.08. with a standard eviation of 0.02 around this mean.," The mean internal error is $<\sigma_i> = 0.08$, with a standard deviation of $0.02$ around this mean."
 of the cells have internal errors within 2 standard deviations from the mean value (0.04.<m; 0.12)., of the cells have internal errors within 2 standard deviations from the mean value $0.04 \leq \sigma_i \leq 0.12$ ).
 The lower panel shows the dependence of internal errors withely: we note that for Ay&1.5 the internal errors increase significantly., The lower panel shows the dependence of internal errors with; we note that for ${\it A_K} > 1.5$ the internal errors increase significantly.
 Since the present extinction map can be useful for a wide varicty of Galactic and extragalactic studies in such central directions. it will be provided in electronic form in the CDS. bv columns: (1) ancl (2) galactic longitude and latitucle of the cell centre. (3) the band. extinction and (4) he uncertainty in the determination στ.," Since the present extinction map can be useful for a wide variety of Galactic and extragalactic studies in such central directions, it will be provided in electronic form in the CDS, by columns: (1) and (2) galactic longitude and latitude of the cell centre, (3) the band extinction and (4) the uncertainty in the determination $\sigma_i$."
 Schultheis ct al. (, Schultheis et al. (
1999). provided. an extinction. map for he inner Galactic Bulge (covering [ff«8 and 1« 5°) obtained [rom heJ and DENIS CMDs. together with isochrones from Aertelli ct al. (,"1999) provided an extinction map for the inner Galactic Bulge (covering $|\ell|<8^{\circ}$ and $|{\it b}|<1.5^{\circ}$ ) obtained from the and DENIS CMDs, together with isochrones from Bertelli et al. ("
1994). hereafter called DISNIS extinction map.,"1994), hereafter called DENIS extinction map."
 The cdilferences tween the present extinction determinaion applied to 2NASS photometry and he procedure adopted: by Schultheis e al. (, The differences between the present extinction determination applied to 2MASS photometry and the procedure adopted by Schultheis et al. (
1999) are the ollowing: (1) the latter use an isochrone from Bertelli οἱ al. (,1999) are the following: (i) the latter use an isochrone from Bertelli et al. (
1994). with metallicity Z—0.02. age of LO Gyr and distance d=S Ipc. to represent the dereddened (Α.Ο ΔΕΗ of the Bulee fields: (ii) Schultheis ct al.,"1994), with metallicity =0.02, age of 10 Gyr and distance =8 Kpc, to represent the dereddened ) CMD of the Bulge fields; (ii) Schultheis et al."
" adopt a transformation romA to with an estimated error of 0.04 mag. in order to x able to use isochrone fitting: and (ii) the DENIS infrared ohotometry is limited to A,—11.0. with detection limits at J=160 and A.=13.0 (Schulthejs ct al."," adopt a transformation from to with an estimated error of 0.04 mag, in order to be able to use isochrone fitting; and (iii) the DENIS infrared photometry is limited to =11.0, with detection limits at $J = 16.0$ and $K_s = 13.0$ (Schultheis et al."
 1999)., 1999).
 Unavane et al. (, Unavane et al. (
"1998) estimaed a DENIS completeness imiting magnitude of A,—10.0 in the inner Bulec.",1998) estimated a DENIS completeness limiting magnitude of =10.0 in the inner Bulge.
 The postution of the DENTS extinction. map is also 4., The resolution of the DENIS extinction map is also $^{\prime}$.
 Lig., Fig.
" 6 shows the comparison of the extinction. values derived rom the 2ALASS photometry. lAoir ith those derived rom the DENIS photometry. .NODENIS- or the area in COoninioln between the two mapμα nos5"" and [0«1.5 )"," 6 shows the comparison of the extinction values derived from the 2MASS photometry, ${\it A_{K,2MASS}}$ , with those derived from the DENIS photometry, ${\it A_{K,DENIS}}$, for the area in common between the two maps $|\ell|<5^{\circ}$ and $|{\it b}|<1.5^{\circ}$ )."
 The extinction values derivec from he 2ALASS and DIELS photometric data presen ul CACOent agreement. especially αρ to =l.0.," The extinction values derived from the 2MASS and DENIS photometric data present an excellent agreement, especially up to =1.0."
 3evoned this imit dkpENIS valttes are higher than Ayospyss ones. bu if we consider ha the uncertainties in the extinction cleermination and ohotometric errors increase in these zones. he agreement is still significant.," Beyond this limit ${\it A_{K,DENIS}}$ values are higher than ${\it A_{K,2MASS}}$ ones, but if we consider that the uncertainties in the extinction determination and photometric errors increase in these zones, the agreement is still significant."
 The departure from identiv line in Fig., The departure from identity line in Fig.
 6 COUd be partially due to dilferences in the filter profiles and Zero-point calibrations adopted by the two surveys., 6 could be partially due to differences in the filter profiles and zero-point calibrations adopted by the two surveys.
 Schlegel et al. (, Schlegel et al. (
1998) (hereafter SEDOS) presented an all-sky reddening map basedon the 100 jm dust emission. modeling the emission by dust. grains with blackbocly radiation at a temperature 7—18.2 Wk. Temperature correctionswere,"1998) (hereafter SFD98) presented an all-sky reddening map basedon the 100 $\mu$ m dust emission, modeling the emission by dust grains with blackbody radiation at a temperature =18.2 K. Temperature correctionswere"
properties in molecular line observations despite their similar internal liminosities.,properties in molecular line observations despite their similar internal luminosities.
 LAA 04191 eet al., IRAM 04191 et al.
 1999: Belloche et al., 1999; Belloche et al.
 2002) n L152MEG Crapsi οἱ al., 2002) and L1521F (Crapsi et al.
associated 2004) show evidence for infall whereas L1014 does not (Crapsi et al., 2004) show evidence for infall whereas L1014 does not (Crapsi et al.
 2005)., 2005).
1 IRAM 04191 is with a well-collimated outflow eel al., IRAM 04191 is associated with a well-collimated outflow et al.
 1999): the other two are not. although al least LIOI4 and possibly L1521E feature weak. compact outflows (Bourke et al.," 1999); the other two are not, although at least L1014 and possibly L1521F feature weak, compact outflows (Bourke et al."
 2005: Bourke et al., 2005; Bourke et al.
 2006)., 2006).
" The discovery of VeLLOs with has put into question the picture of low-mass star lormation as a continuous process of constant mass accretion al the standard rate of —2x10"" vr| (Shu. Adams. Lizano 1987) through a single evolutionary sequence. (he well-established class svstem progressing from Class 0 to III (Myers Lada 1993. André et al."," The discovery of VeLLOs with has put into question the picture of low-mass star formation as a continuous process of constant mass accretion at the standard rate of $\sim 2\times 10^{-6}$ $yr^{-1}$ (Shu, Adams, Lizano 1987) through a single evolutionary sequence, the well-established class system progressing from Class 0 to III (Myers Lada 1993, André et al."
 1993)., 1993).
 This standard accretion rate predicts a much higher Iuminositv than observed for VeLLOs: VeLLOs must feature some combination of a very low central mass aud a very low accretion rate (e.g. Dunham et al., This standard accretion rate predicts a much higher luminosity than observed for VeLLOs; VeLLOs must feature some combination of a very low central mass and a very low accretion rate (e.g. Dunham et al.
 2006)., 2006).
 If the accretion continues at the current low rate to the very small central mass. it might not make a star.," If the accretion continues at the current low rate to the very small central mass, it might not make a star."
 ILowever. the accretion rate is not necessarily constant.," However, the accretion rate is not necessarily constant."
 For instance. FU Orionis (FU Ori) objects undergo outbursts (Bell et al.," For instance, FU Orionis (FU Ori) objects undergo outbursts (Bell et al."
 1995 and references therein)., 1995 and references therein).
 Studies (Vorobvoy Basu 2005 and references therein) lor the nature ol the FU Ori variables suggest. accretion bursts [rom the disk to the central star bv. the thermal instability of the disk., Studies (Vorobyov Basu 2005 and references therein) for the nature of the FU Ori variables suggest accretion bursts from the disk to the central star by the thermal instability of the disk.
 Therefore. (wo potential explanations for the verv low luminosities of VeLLOs are 1) proto-brown cdwarls. and 2) objects in a quiescent phase of the episoclic accretion process.," Therefore, two potential explanations for the very low luminosities of VeLLOs are 1) proto-brown dwarfs, and 2) objects in a quiescent phase of the episodic accretion process."
 The former can be discriminated [rom the latter with studies of the chemistry since they involve vastly different. thermal histories. which is crucial (o the chemical evolution.," The former can be discriminated from the latter with studies of the chemistry since they involve vastly different thermal histories, which is crucial to the chemical evolution."
 The thermal history is especially important in interactions between gas ancl ice: ice evaporation and eas [reeze-out [rom and onto grain surfaces. respectively. depend on the dust temperature (Lee et al.," The thermal history is especially important in interactions between gas and ice; ice evaporation and gas freeze-out from and onto grain surfaces, respectively, depend on the dust temperature (Lee et al."
 2004)., 2004).
 Proto-brown clwarls. with their very low masses. will never experience a hot phase. whereas the outbursts of a evele of episodic accretion. a short time period when (he majority of the mass is dumped onto the central protostar. involve significant. warming of the surrounding dust.," Proto-brown dwarfs, with their very low masses, will never experience a hot phase, whereas the outbursts of a cycle of episodic accretion, a short time period when the majority of the mass is dumped onto the central protostar, involve significant warming of the surrounding dust."
 The quiescent states between outbursts feature much colder dust temperatures., The quiescent states between outbursts feature much colder dust temperatures.
 As a result. envelopes of proto-brown dwarls will be similar (to starless cores in (heir chemical distributions. while objects in a quiescent state of episodic accretion will show different chemical distributions from starless cores or normal. embedded Class 0/1 objects.," As a result, envelopes of proto-brown dwarfs will be similar to starless cores in their chemical distributions, while objects in a quiescent state of episodic accretion will show different chemical distributions from starless cores or normal, embedded Class 0/I objects."
 IRAM 04191 may be undergoing episodie accretion since it features a strong outflow which predicts a higher acceretion rate by (wo orders of magnitude (han inferred. [rom the internal Iuminositv of the source, IRAM 04191 may be undergoing episodic accretion since it features a strong outflow which predicts a higher accretion rate by two orders of magnitude than inferred from the internal luminosity of the source
Aliminum size estimates can be obtained straight{οναον from lines whieh correspond to total coverage.,Miminum size estimates can be obtained straightforwardly from lines which correspond to total coverage.
 The proper separation at reshift z of the light. paths from the two images of a lensed QSO with emission redshift tem due to a lens at redshift z; is given by (Young et al., The proper separation at reshift $z$ of the light paths from the two images of a lensed QSO with emission redshift $z_{\rm em}$ due to a lens at redshift $z_{\rm l}$ is given by (Young et al.
 1981) where do is the observed angular separation of the images., 1981) where $\delta\phi$ is the observed angular separation of the images.
 Phe light paths converge at toon and the angular clameter distances. D. are given by where the comoving distance for a [lat universe (ο =0:Peebles 1993: Loge 2000).," The light paths converge at $z_{\rm con}=0$ and the angular diameter distances, $D$, are given by where is the comoving distance for a flat universe $\Omega_K=0$; Peebles 1993; Hogg 2000)."
 Equ., Equ.
 6 holds for 2. 27zas assumed. here., \ref{sep} holds for $z>z_{\rm l}$ as assumed here.
 Eeami ct al. (, Egami et al. (
2000) argue that the lensing ealaxy should be at ο~3 based on the Einstein ane core racii of their model.,2000) argue that the lensing galaxy should be at $z_{\rm l} \sim 3$ based on the Einstein and core radii of their model.
 The damped ssvsteni at 2~2.974 may |be compatible with this estimate., The damped system at $z\sim 2.974$ may be compatible with this estimate.
 Additionally. Petitjean et al. (," Additionally, Petitjean et al. ("
2000) suggest. that. strong ssvstenis are prime lense candidates because the roughly equal brightnesses of images A and D suggest the LOS are traversing the central regions of the lensing object.,2000) suggest that strong systems are prime lense candidates because the roughly equal brightnesses of images A and B suggest the LOS are traversing the central regions of the lensing object.
 We have thus calculated the separation for five values of 2; as shown in Tables 4. and 5..," We have thus calculated the separation for five values of $z_{\rm l}$ as shown in Tables \ref{tab:sep}
 and \ref{tab:sep2}."
 Phese values include an arbitrary value of 0.7. not corresponding to any known object. three values corresponding to observed. ssystenis and one value from the damped ssvstemi mentioned.," These values include an arbitrary value of 0.7, not corresponding to any known object, three values corresponding to observed systems and one value from the damped system mentioned."
 The observed. angular separation of image B from the combined A}€ image was taken to be 00=0.360 arcsec., The observed angular separation of image B from the combined A+C image was taken to be $\delta\phi=0.369$ arcsec.
 Phere is some ambiguity as to wha one means by the location of the AX|€ image. and we have simply chosen the location of the Uusx-weightecl centro of the A|C image.," There is some ambiguity as to what one means by the location of the A+C image, and we have simply chosen the location of the flux-weighted centroid of the A+C image."
 Size estimates scale with 3o (this stil holds. approximately. in the statistical approach of the nex section) and so the elfect is of the order of no more than a few percent.," Size estimates scale with $\delta\phi$ (this still holds, approximately, in the statistical approach of the next section) and so the effect is of the order of no more than a few percent."
 Thus the results remain qualitatively unallected., Thus the results remain qualitatively unaffected.
 We then applied à maximum likelihood method in order to estimate the most probable size of the absorbers. given the information that we have obtained through this work.," We then applied a maximum likelihood method in order to estimate the most probable size of the absorbers, given the information that we have obtained through this work."
 We assumed two simple gcomoetries for absorbers of uniform size., We assumed two simple geometries for absorbers of uniform size.
 The probability that asphericalcloud is intersected by both LOS. given that it is intersected by one is (MeGill 1990). for Yoo0. 1].and zero otherwise.," The probability that acloud is intersected by both LOS, given that it is intersected by one is (McGill 1990), for $X \in [0,1]$ ,and zero otherwise."
 Lore N(s)= S(MB. where A is the absorber radius.," Here $X(z)\equiv 
S(z)/2R$ , where $R$ is the absorber radius."
 bor, For
should be dominated by imerger-trieeered AGN. even at moderate Iuniünosities (L4~10 cre Ly),"should be dominated by merger-triggered AGN, even at moderate luminosities $L_{\rm x} \sim 10^{43}$ erg $^{-1}$ )."
 This is because increcr-driven accretion in this model is tied to the cosmological ealaxy merger rate. which increases rapidly with redshift (Cousclice ct al.," This is because merger-driven accretion in this model is tied to the cosmological galaxy merger rate, which increases rapidly with redshift (Conselice et al."
 2003: IKartaltepe et al., 2003; Kartaltepe et al.
 2007) whereas quiesceut accretion is related to the mass fiction aud eas fraction of late-type ealaxies. which evolve more slowly.," 2007) whereas quiescent accretion is related to the mass function and gas fraction of late-type galaxies, which evolve more slowly."
 The IIIIOG model predicts that at 2= the number density of quiesceutly accreting AGN will not equal that of mereer-fucled ACN until roughly bolow the kuee in the AGN huuinositv function., The HH06 model predicts that at $z=2$ the number density of quiescently accreting AGN will not equal that of merger-fueled AGN until roughly below the knee in the AGN luminosity function.
 Iu the hard N-rav baud this kuee occurs at a Iuuinositv of roughly Ly~10H erg s+ (Aird et al.," In the hard X-ray band this knee occurs at a luminosity of roughly $L_{\rm X} \sim
10^{44}$ erg $^{-1}$ (Aird et al."
 2010)., 2010).
 This mecaus the predicted X-ray Iuninositv at which au equal fraction of ACN are fueled by quiesceut aud. mierger-triggered accretion at 2= is roughly Lx~1072 cre Ἐν," This means the predicted X-ray luminosity at which an equal fraction of AGN are fueled by quiescent and merger-triggered accretion at $z=2$ is roughly $L_{\rm
 X} \sim 10^{42}$ erg $^{-1}$."
 TÉ we asstune that disk-like hosts are fucling their ACN via internal processes and have not experienced a major merecr in the recent past. then this prediction is at odds with the high disk fraction we observe at Lx~10% eve st. an order of maguitude above this uumositv.," If we assume that disk-like hosts are fueling their AGN via internal processes and have not experienced a major merger in the recent past, then this prediction is at odds with the high disk fraction we observe at $L_{\rm X} \sim
10^{43}$ erg $^{-1}$, an order of magnitude above this luminosity."
 We fiud that the Inuiimnositv at which au equal raction of ACN are losted by disk and splieroid galaxies is roughly Lx~10% eres+., We find that the luminosity at which an equal fraction of AGN are hosted by disk and spheroid galaxies is roughly $L_{\rm X} \sim 10^{43}$ erg $^{-1}$.
 This fiudiug suggests that he stochastic fuchue of SAIBUs is far more prevalent at moderate huuinosities than predicted by the Πο noclel., This finding suggests that the stochastic fueling of SMBHs is far more prevalent at moderate luminosities than predicted by the HH06 model.
 This appareut disagreenmient with the IIIIO6 ficling uodel was previously reported at lower redshifts o» Ceorgakakis et al. (, This apparent disagreement with the HH06 fueling model was previously reported at lower redshifts by Georgakakis et al. (
2009). who found that the contribution to the N-rav. Iuninositv function at τον1 roni ACN iu late-tvpe hosts exceeded the predicted Πο function for stochastically fueled ACN.,"2009), who found that the contribution to the X-ray luminosity function at $z\sim1$ from AGN in late-type hosts exceeded the predicted luminosity function for stochastically fueled AGN."
 It was also noted by Cisternas et al. (, It was also noted by Cisternas et al. (
2011). who found a large Yaction }) of huninous AGN (Ex>10H erg 1) at 2~L hosted by disk-dominated galaxies.,"2011), who found a large fraction ) of luminous AGN $L_{\rm X} > 10^{44}$ erg $^{-1}$ ) at $z\sim1$ hosted by disk-dominated galaxies."
 While the veh disk fraction we observe is similar to what has been weviously reported iu these studies. the disagreement οtween our Ποιος and the predictions of the IITIOG uodel is more acute eiven the higher redshift of our sample aud the strong redshift evolution predicted for nerecr-driven accretion.," While the high disk fraction we observe is similar to what has been previously reported in these studies, the disagreement between our findings and the predictions of the HH06 model is more acute given the higher redshift of our sample and the strong redshift evolution predicted for merger-driven accretion."
" Overall our findings generally agree with au ciereine consensus that inajor galaxy iereers likely plav a subdominant role in triggering mocderate-liuunosity AGN,", Overall our findings generally agree with an emerging consensus that major galaxy mergers likely play a subdominant role in triggering moderate-luminosity AGN.
 This has been asserted from a morphological standpoint by Cisternas et al. (, This has been asserted from a morphological standpoint by Cisternas et al. (
2011) aud Ceorgakalis et al. (,2011) and Georgakakis et al. (
2009) at 2~1. aud by Schawinski et al. (,"2009) at $z\sim1$, and by Schawinski et al. ("
2011) at zo2. based on the large disk fraction found. among ACN hosts.,"2011) at $z\sim2$, based on the large disk fraction found among AGN hosts."
 It has also been proposed by Miillaney et al. (, It has also been proposed by Mullaney et al. (
2011) based ou the average specific star formation rates (SSFR) of ACN hosts out to z~3.,2011) based on the average specific star formation rates (SSFR) of AGN hosts out to $z\sim3$.
 They fiud that a vast majority of hosts have SSERs consistent with the star foriing main sequence (Noeske et al., They find that a vast majority of hosts have SSFRs consistent with the star forming main sequence (Noeske et al.
 2007) and that less than appear to be undergoing a starburst phase., 2007) and that less than appear to be undergoing a starburst phase.
 From this they couclude that the muclear activity in these galaxies is being fucled by internal πουαπστις rather than violent mergers., From this they conclude that the nuclear activity in these galaxies is being fueled by internal mechanisms rather than violent mergers.
 A similar conclusion was also reached by Allevata et al. (, A similar conclusion was also reached by Allevato et al. (
2011) based on the projected clustering of AGN in the COSMOS field out to 123.,2011) based on the projected clustering of AGN in the COSMOS field out to $z\sim2.2$.
 There are several reasons why nonanerger related accretion inav contribute more to the onset of ACN activity at this redshift than previously expected., There are several reasons why non-merger related accretion may contribute more to the onset of AGN activity at this redshift than previously expected.
 This includes such thiugs as a shorter post-blowout quasar lhfetime. which would reduce the contribution frou moreci-trigecred ACN to the N-rav luminosity function. or a faster evolving gas fraction than that assumed bv IHIIO6.," This includes such things as a shorter post-blowout quasar lifetime, which would reduce the contribution from merger-triggered AGN to the X-ray luminosity function, or a faster evolving gas fraction than that assumed by HH06."
 It may also be due to the rise of violeut eravitational imstabilities iu disk ealaxics due to the effects of rapid cold flow accretion (Dekel. Sari. Coeoverino 2009).," It may also be due to the rise of violent gravitational instabilities in disk galaxies due to the effects of rapid cold flow accretion (Dekel, Sari, Ceverino 2009)."
 Such instabilities become iucreasiuelv cohbunon at z21 (Ehueercen et al., Such instabilities become increasingly common at $z>1$ (Elmegreen et al.
 05. Cenzel et al.," 05, Genzel et al."
 06) and are not accounted for in the IHHIOG model., 06) and are not accounted for in the HH06 model.
 Unlike the weaker disk iustabilities that are associated with secular evolution at low redshift (0.9.. bar iustabilities). these high redshift instabilities are highly cficicut at continuously fuuucliug seas aud stars to the centers of ealaxies on short timescales (a single disk rotation) aud at hieh inflow rates (~LOAD. ft: Cacciato Dekel 2011). potentially fueliug increased ACN activity in disk ealaxies without the need for ealaxv-ealaxy icrecrs (Bournaurd et al.," Unlike the weaker disk instabilities that are associated with secular evolution at low redshift (e.g., bar instabilities), these high redshift instabilities are highly efficient at continuously funneling gas and stars to the centers of galaxies on short timescales (a single disk rotation) and at high inflow rates $\sim10 M_{\odot}$ $^{-1}$; Cacciato Dekel 2011), potentially fueling increased AGN activity in disk galaxies without the need for galaxy-galaxy mergers (Bournaurd et al."
 2011. in prep).," 2011, in prep)."
 Of course. the disagreciuent between the high disk fraction we observe and the merger-doniüuated fueliug model is predicated on the assumption that disk-like hosts have not experienced a mereer iu the recent past.," Of course, the disagreement between the high disk fraction we observe and the merger-dominated fueling model is predicated on the assumption that disk-like hosts have not experienced a merger in the recent past."
 The two cau be reconciled if these disks lave instead survived or reformed following a merger event., The two can be reconciled if these disks have instead survived or reformed following a merger event.
 Munerical simulations have shown that disks cau reform after a auerecr if the iuteractiug systems are gas rich (Robertson et al., Numerical simulations have shown that disks can reform after a merger if the interacting systems are gas rich (Robertson et al.
 2006: Bundy et al., 2006; Bundy et al.
 2010). although it has been argued that such interactions are nof conducive to the fueling of SMDITs (Hopkius Ieruquist 20093.," 2010), although it has been argued that such interactions are not conducive to the fueling of SMBHs (Hopkins Hernquist 2009)."
 Alternatively. minor mereers provide a leans to trigger ACN-activity within galaxies without cutirely destroviug their pre-existing morphology.," Alternatively, minor mergers provide a means to trigger AGN-activity within galaxies without entirely destroying their pre-existing morphology."
 Seni-analvtic cosmological galaxy formation models in which all ACN activity is assumed to be triggered by imerecrs (Somerville ct al., Semi-analytic cosmological galaxy formation models in which all AGN activity is assumed to be triggered by mergers (Somerville et al.
 2008) do predict that the average lerecr event that trigecrs an AGN with Lx>10 eres + las a ass ratio of 115 as opposed to the more disruptive 1:1 or 1:2 mergers (Somerville ct al., 2008) do predict that the average merger event that triggers an AGN with $L_{\rm X} > 10^{42}$ erg $^{-1}$ has a mass ratio of 1:8 as opposed to the more disruptive 1:1 or 1:2 mergers (Somerville et al.
 2011)., 2011).
 Coupled with a tine delay between the merger aud the visibility of the ACN. the signatures of these mergers could prove difficult to detect.," Coupled with a time delay between the merger and the visibility of the AGN, the signatures of these mergers could prove difficult to detect."
 Therefore. since we caunot rule out such interactions. minor mergers would secun to be one of the remaining wavs to reconcile the merecr-dominated fueling mocel with the hieh disk fraction aud lack of disturbed morphologics that we observe.," Therefore, since we cannot rule out such interactions, minor mergers would seem to be one of the remaining ways to reconcile the merger-dominated fueling model with the high disk fraction and lack of disturbed morphologies that we observe."
 To explore if major galaxy mergers are the primary mechanisiu fucling AGN activity at i~2. we have used /WNTEC2 imaging to examine the rest-frame optical morphologies of ealaxics hosting nmodoerate-Iuninosity. X-ray selected ACN at 2=15weD.," To explore if major galaxy mergers are the primary mechanism fueling AGN activity at $z\sim2$, we have used /WFC3 imaging to examine the rest-frame optical morphologies of galaxies hosting moderate-luminosity, X-ray selected AGN at $z=1.5-2.5$."
 Eaploving visual classifications. we have determined both the predominant morphological type of these galaxies and the frequency at which they exhibit morphological disturbances indicative of recent interactions.," Employing visual classifications, we have determined both the predominant morphological type of these galaxies and the frequency at which they exhibit morphological disturbances indicative of recent interactions."
 To determine if the ACN hosts show ierecr or interaction signatures more often than simular non-active galaxies. we have also classified a saluple of mass-inatched coutrol galaxies at the same redshift.," To determine if the AGN hosts show merger or interaction signatures more often than similar non-active galaxies, we have also classified a sample of mass-matched control galaxies at the same redshift."
 First. we fined that just over half of the AGN reside in disk ealaxies (51.1 23%). while a smaller percentage," First, we find that just over half of the AGN reside in disk galaxies $51.4^{+5.8}_{-5.9}\%$ ), while a smaller percentage"
A reliable age for the local Galactic Dise places a valuable constraint on both elobular cluster ages ancl cosmological models.,A reliable age for the local Galactic Disc places a valuable constraint on both globular cluster ages and cosmological models.
 A number of independent methods of investigating this problem have been emploved in the past (eg., A number of independent methods of investigating this problem have been employed in the past (eg.
 Jimenez 998 and references therein). resulting in à broad consensus that the lower limit for the Disc age lies between 5 and 2 Gyr.," Jimenez 1998 and references therein), resulting in a broad consensus that the lower limit for the Disc age lies between 8 and 12 Gyr."
 Potentially one of the most. reliable means of estimating the Disc age is via cool white dwarl (CWD) stars., Potentially one of the most reliable means of estimating the Disc age is via cool white dwarf (CWD) stars.
 These estimates use the idea. first. proposed. by Schmidt (1050). that in a galaxy of finite age there will be a emperature bevond which the oldest. coolest white cars (WDs) have not had time to cool.," These estimates use the idea, first proposed by Schmidt (1959), that in a galaxy of finite age there will be a temperature beyond which the oldest, coolest white dwarfs (WDs) have not had time to cool."
 This predicted: eut-olf in the luminosity function. (LE) of Ws. if satisfactorily observed. can then be used in conjunction with WD cooling mocels to derive the Disc age.," This predicted cut-off in the luminosity function (LF) of WDs, if satisfactorily observed, can then be used in conjunction with WD cooling models to derive the Disc age."
 CWDs are difficult to finc. being both extremely faint and of similar colour to the numerous. [x and. Al-type cdwarfs: and are almost exclusively discovered. by means of their proper motion.," CWDs are difficult to find, being both extremely faint and of similar colour to the numerous K and M-type dwarfs; and are almost exclusively discovered by means of their proper motion."
 The cut-off. in the. WDLE was observed. (Liebert. Dahn Monet. LOSS. hereafter. LDAL) after thorough follow up observations of CWD candidates crawn from the Luyten Half Second. (LIIS). Catalogue (Luvten. 1979).," The cut-off in the WDLF was observed (Liebert, Dahn Monet 1988, hereafter LDM) after thorough follow up observations of CWD candidates drawn from the Luyten Half Second (LHS) Catalogue (Luyten 1979)."
 Although at that time a Disc. age of 9342 Gyr was derived from this sample (Winget et al., Although at that time a Disc age of $9.3\pm2$ Gyr was derived from this sample (Winget et al.
 987). further observations ancl improvements in moclel atmospheres (Bergeron. Ruiz Leggett 1997. hereafter BL) and theoretical LP's (Wood 1992. 1995) has prompted a recent redetermination of the Disc age for the same sample (Leggett. Ruiz Bergeron 1998. hereafter LRB). vielding a value of 8Sd1.5 Gyr.," 1987), further observations and improvements in model atmospheres (Bergeron, Ruiz Leggett 1997, hereafter BRL) and theoretical LFs (Wood 1992, 1995) has prompted a recent redetermination of the Disc age for the same sample (Leggett, Ruiz Bergeron 1998, hereafter LRB), yielding a value of $8\pm1.5$ Gyr."
 While the existence of the cut-olf in he LDM WDLE has not been challenged. by subsequent observational work. the details of its precise position and shape have.," While the existence of the cut-off in the LDM WDLF has not been challenged by subsequent observational work, the details of its precise position and shape have."
 A sample of CWDs found using conimon proper motion binaries (CPMDs). again cullecl from the Luvten surveys. suggest that there are ~5 times more very faint CWDs than found by ΕΟΝ (Oswalt et al.," A sample of CWDs found using common proper motion binaries (CPMBs), again culled from the Luyten surveys, suggest that there are $\sim5$ times more very faint CWDs than found by LDM (Oswalt et al."
 1996. hereafter ΟΝΕ).," 1996, hereafter OSWH)."
 A Disc age of 9.541 Gyr was found using this sample. and the factor of 5 increase in the faintest WDs has been confirmed by an independent search for CWDs in the south (Ruiz and Takamiva 1995).," A Disc age of $9.5\pm1$ Gyr was found using this sample, and the factor of $\sim5$ increase in the faintest WDs has been confirmed by an independent search for CWDs in the south (Ruiz and Takamiya 1995)."
 Until now. the proper motion catalogues usec to extract samples of C\WDs have been produced: by blink? comparison.," Until now, the proper motion catalogues used to extract samples of CWDs have been produced by `blink' comparison."
 While these surveys have clearly been successful in picking up individual stars of low luminosity and high, While these surveys have clearly been successful in picking up individual stars of low luminosity and high
1l min in QQ Vul (Osborne.Cropper&Christiani:1987).. and in AL Tr 6.5 7 min and 13.5. 14 min (Schwarz 1998).,"11 min in QQ Vul \cite{b24}, and in AI Tri 6.5– 7 min and 13.5 – 14 min \cite{b36}."
". We speculate that the QPOs in V1309 Ori are due to ""blobby. accretion’. already observed in X-ray data (Walter. Wolk Aclams 1995: cle Martino et al."," We speculate that the QPOs in V1309 Ori are due to 'blobby accretion', already observed in X-ray data (Walter, Wolk Adams 1995; de Martino et al."
 1998)., 1998).
 We compare these timescales with those derived: by Ixing (1989).. and King (1995). (see also Chanmugam 1995).," We compare these timescales with those derived by King \shortcite{b17}, and King \shortcite{b18} (see also Chanmugam 1995)."
" The irradiation of the accretion Dow may ionise the subsonic acecretion [low below the inner £4, point and modulate eas How through this point on the timescale of the dynamical time scale in the Roche potential near £4.", The irradiation of the accretion flow may ionise the subsonic accretion flow below the inner $L_{1}$ point and modulate gas flow through this point on the timescale of the dynamical time scale in the Roche potential near $L_{1}$ .
" The equation presented by Ixing (1989) where LL, is the scale height and c; is local sound speed near Ly. predicts the timescales for these oscillations."," The equation presented by King \shortcite{b17}
 where ${\rm H_{*}}$ is the scale height and ${\rm c_{*}}$ is local sound speed near ${\rm L_{1}}$, predicts the timescales for these oscillations."
 The orbital period of 479 minutes gives us a prediction of 26 minutes for V1309 Ori., The orbital period of 479 minutes gives us a prediction of 26 minutes for V1309 Ori.
 The observed. timescales of QPOs in V1309 Ori. LO and. 15 minutes. are approximately half of the predicted.," The observed timescales of QPOs in V1309 Ori, 10 and 15 minutes, are approximately half of the predicted."
 The QPOs are seen strongest in CDV. and are negligible in the longer wavelengths. which may due to reason that only a fraction of How. will undergo. these oscillations near Ly. as pointed out by Wing (1995).," The QPOs are seen strongest in $UBV$, and are negligible in the longer wavelengths, which may due to reason that only a fraction of flow will undergo these oscillations near ${\rm L_{1}}$, as pointed out by King \shortcite{b18}."
 several eclipsing polars have been observed with high signal to noise ancl high time resolution., Several eclipsing polars have been observed with high signal to noise and high time resolution.
 These include HU. αν (Llarrop-Allinetal.1999). and. UZ For (Perrymanctal. 2001)., These include HU Aqr \cite{b15} and UZ For \cite{b26}.
. Both of these systenis show a sharp eclipse ingress lasting several seconds., Both of these systems show a sharp eclipse ingress lasting several seconds.
 This sharp drop in intensity is associated with the eclipse of the bright accretion region on the white dwarf., This sharp drop in intensity is associated with the eclipse of the bright accretion region on the white dwarf.
 In the case of UZ For there are. two sharp intensity drops. indicating that there are two accretion regions visible during the eclipse.," In the case of UZ For there are two sharp intensity drops, indicating that there are two accretion regions visible during the eclipse."
 One of the most striking features about the eclipse profiles of V1309 Ori is the obvious lack of a sharp ingress or egress which indicate the (cisjappearance of the white dwarf and/or hot spots in the surface of the white chwarl behind the secondary., One of the most striking features about the eclipse profiles of V1309 Ori is the obvious lack of a sharp ingress or egress which indicate the (dis)appearance of the white dwarf and/or hot spots in the surface of the white dwarf behind the secondary.
 The relative faintness of the accretion region(s). compared to the bright accretion stream in VI309 Ori. may be the reason for this.," The relative faintness of the accretion region(s), compared to the bright accretion stream in V1309 Ori, may be the reason for this."
 After the eclipse of the white dwarf. the accretion stream is still visible for a length of time.," After the eclipse of the white dwarf, the accretion stream is still visible for a length of time."
 Observations of other polars such as HU Jer (Glenn et al., Observations of other polars such as HU Aqr (Glenn et al.
 1994. Dridge et al.," 1994, Bridge et al."
 2002) show that this length of time can vary [rom one cycle to the next., 2002) show that this length of time can vary from one cycle to the next.
 Phe fact that we observe a variable eclipse ingress of the accretion stream in V1309 Ori is therefore not un-typical of polars., The fact that we observe a variable eclipse ingress of the accretion stream in V1309 Ori is therefore not un-typical of polars.
 However. what does make V1309. Ori unique amongst polars is the fact that the eclipse egress is highly variable: all other polars show a rapid rise at the same phase coming out from eclipse.," However, what does make V1309 Ori unique amongst polars is the fact that the eclipse egress is highly variable: all other polars show a rapid rise at the same phase coming out from eclipse."
 Phe fact that V1309 Ori does not implies either that we can observe the stream above the orbital plane before the white dwarf is visible. or that the accretion stream travels far enough around the white dwacl so that it is visible before the white cwarl itself.," The fact that V1309 Ori does not implies either that we can observe the stream above the orbital plane before the white dwarf is visible, or that the accretion stream travels far enough around the white dwarf so that it is visible before the white dwarf itself."
 To investigate if this further. we shown in Ligure ll the view of the system. at two different phases (b= 0.96 and 1.04).," To investigate if this further, we shown in Figure \ref{system} the view of the system at two different phases $\Phi$ = 0.96 and 1.04)."
" We use the following svsten parameters in determining these: ;/=TS"". q=0.67. Alay=O.7 (Staude.Schwope&Schwarz2001) and a white dwarf - secondary star separation of 104+ em (determined using the above parameters. ancl standard. Roche lobe e&cometrv)."," We use the following system parameters in determining these: $i=78^{\circ}$, $q$ =0.67, $M_{wd}=0.7$ \cite{b42} and a white dwarf - secondary star separation of $\times10^{11}$ cm (determined using the above parameters and standard Roche lobe geometry)."
 We also show a single magnetic field. line originating from the negative circularly polarised accretion region (3=35 and lace onto the observer at 0.2. 87.3)).," We also show a single magnetic field line originating from the negative circularly polarised accretion region $\beta=35^{\circ}$ and face onto the observer at $\Phi$ =0.2, \ref{modelling}) )."
 In Figure 11.. the accretion streams leading to both poles are visible at &=0.96.," In Figure \ref{system}, the accretion streams leading to both poles are visible at $\rm \Phi=0.96$."
 At —0.04 only the stream leading to the negative pole is visible., At $\rm \Phi=0.04$ only the stream leading to the negative pole is visible.
 “Phe white dwarf. appears before the stream [leading to the positive pole is. visible., The white dwarf appears before the stream leading to the positive pole is visible.
 Even if the negative pole is not visible at 0.04. the emission from the stream leading to that pole is.," Even if the negative pole is not visible at 0.04, the emission from the stream leading to that pole is."
" We take £2),/a—0.2: thiss implies. 2,=.2.9.107Lu cm.", We take $R_{\mu}/a$ =0.2: this implies $R_{\mu}=2.9\times10^{10}$ cm.
 «prosThis is consistent.. with. our findings in ES.3.., This is consistent with our findings in \ref{mass}.
 This shows that for this accretion stream geometry the accretion stream. is. visible after the white dwarf has been eclipsed ancl also before the white cwarl comes out of eclipse., This shows that for this accretion stream geometry the accretion stream is visible after the white dwarf has been eclipsed and also before the white dwarf comes out of eclipse.
 I£ the stream emission was highly variable then this could explain the variable egress profile., If the stream emission was highly variable then this could explain the variable egress profile.
 CGarnavich et al., Garnavich et al.
 (1994) and Shafter et al., \shortcite{b12} and Shafter et al.
 (1995). noted that the secondary star in VI309. Ori is oversized. for its spectral type (MO - MI) and mass (0.4:0.6 AZ. )., \shortcite{b35} noted that the secondary star in V1309 Ori is oversized for its spectral type (M0 - M1) and mass (0.4–0.6 $M_{\odot}$ ).
 Indeed. recent binary evolution models (eg Smith Dhillon 1998 and Baralle Ixolb 2000) which assume an unevolved donor star and typical mass transfer rates. all predict either a much earlier or later spectral type for V1309 Ori than observed (cf Figure |. of Baralle Ixolb 2000).," Indeed, recent binary evolution models (eg Smith Dhillon 1998 and Baraffe Kolb 2000) which assume an unevolved donor star and typical mass transfer rates, all predict either a much earlier or later spectral type for V1309 Ori than observed (cf Figure 1 of Baraffe Kolb 2000)."
 Llowever. by assuming an evolved donor it is possible to match the observed spectral (ype to the predicted. value.," However, by assuming an evolved donor it is possible to match the observed spectral type to the predicted value."
 For instance. [rom Figure 3 of Baralle Ixolb (2000)... for an initial secondary star mass of Mo—1.2M... a central llvdrogen abundance at the start of mass transfer. of X005. and a mass transfer rate of M=15.1. 101) g 1) we fined we that a spectral type of ALL is predicted for an orbital period of S hrs.," For instance, from Figure 3 of Baraffe Kolb \shortcite{b1}, for an initial secondary star mass of $M_{2}$, a central Hydrogen abundance at the start of mass transfer of $X_{c}$ =0.05, and a mass transfer rate of $\dot{M} =1.5
\times 10^{-9} M_{\odot} (\sim1\times10^{17}$ ) g $^{-1}$ ) we find we that a spectral type of M1 is predicted for an orbital period of 8 hrs."
 This is within the estimated range of values required to satisfy. conditions for observed oversized secondary to fll its Roche lobe (the main uncertainties are the distance and the mass of the white dwarl. eg Llarrop-Allin ct al.," This is within the estimated range of values required to satisfy conditions for observed oversized secondary to fill its Roche lobe (the main uncertainties are the distance and the mass of the white dwarf, eg Harrop-Allin et al."
 1997. de Martino et al.," 1997, de Martino et al."
 1998)., 1998).
 As already noted by Ixing. Osborne Schenker (2002).. V1309 Ori is indeed. à good. possible candidate for a binarysystem which has gone through a supersoft: source. phase and contains a nuclear evolved donor star.," As already noted by King, Osborne Schenker \shortcite{b19}, V1309 Ori is indeed a good possible candidate for a binarysystem which has gone through a supersoft source phase and contains a nuclear evolved donor star."
 Szkocdy Silber (1996) ancl Schmidt Stockman (2001) have. noticed that V1309 Ori has extraordinary strong excitation lines of NVAI240 and OV AI370) which may support such an interpretation.," Szkody Silber \shortcite{b43} and Schmidt Stockman \shortcite{b33} have noticed that V1309 Ori has extraordinary strong excitation lines of ${
\rm N~V~\lambda 1240}$ and ${ \rm O~V~\lambda1370}$ ) which may support such an interpretation."
cni.,cm.
 The frequencies at which clectrous with energy Las cnt their peak power for cach cutoff mechanisin are then Of course. iu a given object. the lowest value of Eyes will be the operative value.," The frequencies at which electrons with energy $E_{\rm max}$ emit their peak power for each cutoff mechanism are then Of course, in a given object, the lowest value of $E_{\rm max}$ will be the operative value."
 Thus if one can determine the mechanism causing the spectral cutoff. its value constrains considerably more pliysical parameters than the simple observation of a radio power-law spectra.," Thus if one can determine the mechanism causing the spectral cutoff, its value constrains considerably more physical parameters than the simple observation of a radio power-law spectrum."
 Populations- of+ relativisticol- ious- aud clectrous can produci observable coutinmun radiation through four mechanisius. one hadronic aud three leptonic (reviewed iu Revuolds 20082).," Populations of relativistic ions and electrons can produce observable continuum radiation through four mechanisms, one hadronic and three leptonic (reviewed in Reynolds 2008a)."
 The hadronic mechanisi is the inelastic κοπομοβ of cosuiccray protons on thermal uncle. producing vious.," The hadronic mechanism is the inelastic scattering of cosmic-ray protons on thermal nuclei, producing pions."
 The charged pious decay to electrous aud yositrous. making a (probably) uceligible contribution o the relativistic lepton pool.but the zx decay o eununa rays of comparable cnerey E-(uin)=m2/2~70 MeV. The spectrum of emitted photons should be that of the ious that produce them.," The charged pions decay to electrons and positrons, making a (probably) negligible contribution to the relativistic lepton pool,but the $\pi^0$ 's decay to gamma rays of comparable energy $E_\gamma ({\rm min}) = m_\pi c^2/2 \sim 70$ MeV. The spectrum of emitted photons should be that of the ions that produce them."
 The three eptonie processes are svichrotron emission. described above. as well ax nouthermal brouisstralilung. with the same spectrum as that of the nonthermal electrons. (IC) upscatterine of photons from on observed any significant ambicut radiation.," The three leptonic processes are synchrotron emission, described above, as well as nonthermal bremsstrahlung, with the same spectrum as that of the nonthermal electrons, and inverse-Compton (IC) upscattering of photons from any significant ambient radiation."
" Tn practice. this is ikely to be primarily the cosmic microwave backerouud (CMB). though iu. some cases. IC. from"" UV-opticaTc-IR photons may be competitive."," In practice, this is likely to be primarily the cosmic microwave background (CMB), though in some cases, IC from UV-optical-IR photons may be competitive."
" The spectrum wil jio the same as that of the svuchrotron ciission from whatever population of particles is responsible. that ix. considerable harder iu the keV TeV range than hat of the other processes,"," The spectrum will be the same as that of the synchrotron emission from whatever population of particles is responsible, that is, considerably harder in the keV – TeV range than that of the other processes."
 Whilen. the synchrotron Xxocess is clearly operating. it is not clear which of he other processes unight be responsible for cussion roni any particular object.," While the synchrotron process is clearly operating, it is not clear which of the other processes might be responsible for emission from any particular object."
 The best evidence for ion acceleration is the spectral feature resulting from tlie photon2. οσον from a+ created pion nearly. at either rest. about 70 MeV. Detailed caleulatious (c.g.Dariugetal.1999) show that this feature may not be liehly distinct in a real object.," The best evidence for ion acceleration is the spectral feature resulting from the minimum photon energy from a created pion nearly at rest, about 70 MeV. Detailed calculations \citep[e.g.,][]{baring99} show that this feature may not be highly distinct in a real object."
 Several shell (ie. not containing a pulsar) SNR& have been detected in TeV sanuua rays. using air-Coereukoy detectors such as the IHigb-Euergy Stereoscopie System iu Namibia.," Several shell (i.e., not containing a pulsar) SNRs have been detected in TeV gamma rays, using air-\v{C}eerenkov detectors such as the High-Energy Stereoscopic System in Namibia."
 These include €0317.3-0.5. Vela Jr. ROW 86. and SN 1006sa).," These include G347.3-0.5, Vela Jr., RCW 86, and SN 1006."
 The spectra are steep. with photon indices PT~2 (FLx(hv)DP for G317.3-0.5.the spectrum. is observed to steepen above 1 TeV. Elaborate models for jese four cases have been constructed (e.g..Berezhko&VOIk2006).," The spectra are steep, with photon indices $\Gamma \sim 2$ $F_\gamma
\propto (h\nu)^-\Gamma$ ); for G347.3-0.5, the spectrum is observed to steepen above 1 TeV. Elaborate models for these four cases have been constructed \citep[e.g.,][]{berezhko06}."
. The TeV emission in these models can jo due either to IC from the CAMB. or to z? decay.," The TeV emission in these models can be due either to IC from the CMB, or to $\pi^0$ decay."
 Both classes of ΠοΟΙ have difficulty., Both classes of model have difficulty.
 The former require low chline factors of magnetic feld. aud imply ineficicut shock acceleration. while the latter sufferfrou severe imuits ou thermal gas fou X-rayobservations. implying insufiicicnt targets for the relativistic protons (Ellison(al.2010).," The former require low filling factors of magnetic field, and imply inefficient shock acceleration, while the latter sufferfrom severe limits on thermal gas from X-ray observations, implying insufficient targets for the relativistic protons \citep{ellison10}."
.. For complex objects such as C317.3- and Vela Jh. simple ouc-zone models nay be inadequate. though there is as vet no clear path to consistent models.," For complex objects such as G347.3-0.5 and Vela Jr., simple one-zone models may be inadequate, though there is as yet no clear path to consistent models."
 An iniportant result first suggested many years ago. but only recently put ou a firmer observational foundation. 18 the merease in magnetic-feld strength in SNRs over 2 ο factor of the compression ratio expected iu üghlv ionized gases.," An important result first suggested many years ago, but only recently put on a firmer observational foundation, is the increase in magnetic-field strength in SNRs over a simple factor of the compression ratio expected in highly ionized gases."
 The possibility of uch stronger Shock amplification of maguctic fields was first proposed or heliospherie shocks by Chevalier (1977). while the uodels of Revuolds&Chevalier(1981)— demanded Sthstantial amplification.," The possibility of much stronger shock amplification of magnetic fields was first proposed for heliospheric shocks by Chevalier (1977), while the models of \cite{reynolds81} demanded substantial amplification."
 Early gamma-ray upper Inti roni Cas A bounded the electron population from 2bove (from the inferred absence of brenisstralilung). vottnding the maguetic-field strength from below based radio svuchrotroutae] fluxes.," Early gamma-ray upper limits from Cas A bounded the electron population from above (from the inferred absence of bremsstrahlung), bounding the magnetic-field strength from below based on observed radio synchrotron fluxes."
"⋅ Cowsikware! &rate heSarkarSalesaudinverse-Compton(L980) used this argument to deduce a dium uagneticetic fieldFe streusthdvo of to.about 7L1 πλαν1 The""e slightele COLllCAN(.yo or]curvatiununD observedOTIO iuH radiom1 spectra:vetas of""CONDO. SNRs was explained by Revnolds&Ellison(1992). as due o cficicut shock acceleration modifying the shock structure, but values of maguotic field were also implied of 100 μέ and more — far higher than a few times the vpical ⋅interstellar∖⋅↴∖⊀⋅ maeuctic ⊀↜∖⋅⋡⋅∖field of ⊳∙≻3 54 ""es Ilish-resolution.[m] N-yav nuagesoO from show hat the “thin ritus"" present in most historical shell SNRs ave in fact very thin so thin that au unusual depletion (beyond. simple post-shock expansion) of relativisticàwe clectrousvolatixsetie orlest↴ magnetic⋅ field ⊀≼⊾∖⋅⊳∎∖is required ivonininm to explain the sudden disappearauce of svuchrotrou cuissivity downstream (Bambaetal,2003:Vink& 2003).."," \cite{cowsik80} used this argument to deduce a minimum magnetic field strength of about 1 mG. The slight concave curvature observed in radio spectra of SNRs was explained by \cite{reynolds92}
 as due to efficient shock acceleration modifying the shock structure, but values of magnetic field were also implied of 100 $\mu$ G and more – far higher than a few times the typical interstellar magnetic field of 3 – 5 $\mu$ G. High-resolution X-ray images from show that the “thin rims” present in most historical shell SNRs are in fact very thin – so thin that an unusual depletion (beyond simple post-shock expansion) of either relativistic electrons or magnetic field is required to explain the sudden disappearance of synchrotron emissivity downstream \citep{bamba03,vink03}. ."
 If synchrotron losses are depleting the Clectrons. we iufer (Parizotetal.2006) where the shock speed ay=«4/1000 kins 1 and the fibuneut width is e.," If synchrotron losses are depleting the electrons, we infer \citep{parizot06}
 where the shock speed $u_8 \equiv u/1000$ km $^{-1}$ and the filament width is $w$ ."
 However. if the amplified magnetic," However, if the amplified magnetic"
is also found as temperature dependence of Ale IIA.& emission in stars (both dwarls and ejants) of minimal activitv. a likely indicator of the overall chromospheric energv. density (Buchholz.Ulnschneider.&Cuntz1993.seetheirFig.15)..,"is also found as temperature dependence of Mg II, emission in stars (both dwarfs and giants) of minimal activity, a likely indicator of the overall chromospheric energy density \cite*[][see their Fig.~15]{buc98}."
" If the mechanical energy. [αν Fy, utilized [ου generating stellar mass loss is assumed as FNxLi. the surlace-integratecd mechanical energy. [lux Lay can now be expressed as Next we consider the characteristic chromospheric radius. Rey."," If the mechanical energy flux $F_{\rm M}$ utilized for generating stellar mass loss is assumed as $F_{\rm M} \propto T_{\rm eff}^{7.5}$, the surface-integrated mechanical energy flux $L_{\rm M}$ can now be expressed as Next we consider the characteristic chromospheric radius $R_{\rm Chr}$."
 For cool giants aud supereiants. no well-defined boundary between (he chromosphere and (he wind exists.," For cool giants and supergiants, no well-defined boundary between the chromosphere and the wind exists."
 Hence. we use (he sonic point of the average velocity fiekl as relerence.," Hence, we use the sonic point of the average velocity field as reference."
" For the well-stuclied Ix supergiant ¢ Aur (with logg,~ 0.8). Rey, is found to be close to 222, (Dadeοἱal.1996).. and for general giants and supergiants. (Rep,—1,)/R. is assumed (o varv as g,|. which elves With the above temperature dependence of the mechanical energy flux. (eq. ["," For the well-studied K supergiant $\zeta$ Aur (with $\log{g_*} \simeq 0.8$ ), $R_{\rm Chr}$ is found to be close to $2 R_*$ \citep{baa96}, and for general giants and supergiants, $(R_{\rm Chr}-R_*)/R_*$ is assumed to vary as $g_*^{-1}$, which gives With the above temperature dependence of the mechanical energy flux (eq. ["
"2]) and chromospheric radius Rey, (eq. [",2]) and chromospheric radius $R_{\rm Chr}$ (eq. [
3]). we finally obtain as mass loss rate JM. see eq. (,"3]), we finally obtain as mass loss rate $\dot{M}$, see eq. ("
"1). with f2,. M,. and L, as stellar raclius. mass. ancl luminosity given in solar units. aud gy, and g. as stellar ancl solar surface gravity. respectively.","1), with $R_*$ , $M_*$, and $L_*$ as stellar radius, mass, and luminosity given in solar units, and $g_*$ and $g_\odot$ as stellar and solar surface gravity, respectively."
 This is. apart from the two new factors. indeed the old Reimers law.," This is, apart from the two new factors, indeed the old Reimers law."
" To satisfy the well-constrained RGB mass-loss of globular cluster stars (see 3). the fitting parameter 7 will be set to S(E1)xLOM, fl"," To satisfy the well-constrained RGB mass-loss of globular cluster stars (see 3), the fitting parameter $\eta$ will be set to $8 (\pm1) \times 10^{-14} M_{\odot}$ $^{-1}$."
 For the newly developed mass loss formula various tests and applications have been devised., For the newly developed mass loss formula various tests and applications have been devised.
" In particular. we want to oblain insight into the importance of the new [actors eiven bv the stellar effective temperature Zur and eravily g,."," In particular, we want to obtain insight into the importance of the new factors given by the stellar effective temperature $T_{\rm eff}$ and gravity $g_*$."
 In fact. for ordinary giants the (vo new [actors do not make much difference. which explains (he long-lasting success of the Reimers relation.," In fact, for ordinary giants the two new factors do not make much difference, which explains the long-lasting success of the Reimers relation."
 In particular. the Z5. factor is. despite its high power. restricted in iis impact by the small band of relevant. effective temperatures (3000 (ο 4500 Ix). and the g-sensilive factor remains of the order of one for all but the smallest gravities.," In particular, the $T_{\rm eff}^{3.5}$ factor is, despite its high power, restricted in its impact by the small band of relevant effective temperatures (3000 to 4500 K), and the $g$ -sensitive factor remains of the order of one for all but the smallest gravities."
 In fact. as previously discussed. the Z;r exponentin eq. (," In fact, as previously discussed, the $T_{\rm eff}$ exponentin eq. ("
4) is somewhat uncertain.,4) is somewhat uncertain.
 However. due to the," However, due to the"
the wavelength.,the wavelength.
 If the grain size changes with cometocentric distance. the light scattered by the dust will be affected in a different manner at different wavelengths.," If the grain size changes with cometocentric distance, the light scattered by the dust will be affected in a different manner at different wavelengths."
 The images obtained with SUSI 2 on March 24. 2007 did not allow us to detect any cometary activity.," The images obtained with SUSI 2 on March 24, 2007 did not allow us to detect any cometary activity."
 Because no photometric standard star could be observed during the same night (nonphotometric night) we have used the predicted Echeclus average magnitude for calibrating the profile with absolute magnitudes., Because no photometric standard star could be observed during the same night (nonphotometric night) we have used the predicted Echeclus average magnitude for calibrating the profile with absolute magnitudes.
 We base this prediction on the phase curve published by ?.. and the heliocentric and geocentric distances. as well as the phase angle during the observations.," We base this prediction on the phase curve published by \cite{rousselot:2005a}, and the heliocentric and geocentric distances, as well as the phase angle during the observations."
 Figure 7. presents the surface brightness profile of Echeclus obtained when all the R-band images are co-added (after centering)., Figure \ref{f:sbp} presents the surface brightness profile of Echeclus obtained when all the R-band images are co-added (after centering).
 We compare this profile to the one of a star apprearingσι in the same field of view and adjusted in intensity., We compare this profile to the one of a star apprearing in the same field of view and adjusted in intensity.
 No differences can be detected up to the sky background. i.e. up to Rz27/arcsec.," No differences can be detected up to the sky background, i.e. up to $\simeq$ $^2$."
 We have used the method mentioned by ? to derive an order of magnitude of the maximum Afp parameter that can be derived from these data.," We have used the method mentioned by \cite{jewitt:1984}
 to derive an order of magnitude of the maximum $Af\rho$ parameter that can be derived from these data."
" We have used the R,,,, magnitude where the surface brightness profile reaches the sky background. i.e. about 27/aresec."," We have used the $R_{max}$ magnitude where the surface brightness profile reaches the sky background, i.e. about $^2$."
" From this magnitude we have derived a lower limit for the coma magnitude. if it exists. by using the formula: A,==?5Log(207)Ryu."," From this magnitude we have derived a lower limit for the coma magnitude, if it exists, by using the formula: $R_{coma}\simeq-2.5Log_{10}(2\pi r^2)+R_{max}$."
" 1n this formula. + represents the diameter corresponding to the A,,,, magnitude. ie. 5 arcsec."," In this formula, $r$ represents the diameter corresponding to the $R_{max}$ magnitude, i.e. 5 arcsec."
 These values lead to an upper limit for A£p equal to about 75 em., These values lead to an upper limit for $Af\rho$ equal to about 75 cm.
 This upper limit can be compared to the one measured one year before (10.000 cm for the R-band. see above).," This upper limit can be compared to the one measured one year before (10,000 cm for the R-band, see above)."
 The ratio is about 130., The ratio is about 130.
 By using similar parameters and formulae as the one mentionned in Sect., By using similar parameters and formulae as the one mentionned in Sect.
 3 we derive an upper dust production rate Qiu;=0.0 kes!., 3 we derive an upper dust production rate $Q_{max}\simeq 0.6$ $^{-1}$.
 In the range covered by our spectra different emission bands corresponding to different radicals could be observed., In the range covered by our spectra different emission bands corresponding to different radicals could be observed.
 Among these emission bands the more intense are CN (3880 A)) and C: , Among these emission bands the more intense are CN (3880 ) and $_2$ 
that. moclels he propagation of photons from the neutron star surface towards a distant observer.,that models the propagation of photons from the neutron star surface towards a distant observer.
 To generate ight) curves we use the Oblate Schwarzschile (OS) approximation οἱ Morsinkelal.(2007) to model reativistic light-bending. Doppler shifts ancl eravitational recshift.," To generate light curves we use the Oblate Schwarzschild (OS) approximation of \citet{MLCB} to model relativistic light-bending, Doppler shifts and gravitational redshift."
 Phe OS model. which takes into account rotation-incuced oblateness. is more appropriate for very rapidly rotating neutron stars than the more usual Schwarzschilel Doppler approximation (Poutanen&Cicrliiski20 3).," The OS model, which takes into account rotation-induced oblateness, is more appropriate for very rapidly rotating neutron stars than the more usual Schwarzschild + Doppler approximation \citep{PGie}."
. Assuming that burst oscillation [requenevy is à good measure of stellarspin.. 4U. 1636-536 rotates al zz580 112.," Assuming that burst oscillation frequency is a good measure of stellar, 4U 1636-536 rotates at $\approx 580$ Hz."
 The associated rotational cleformation is a [ow percent. depending on the assumed mass ancl nuclear equation of state. leading to a small but noticeable inlluence on the light curve (Morsinkctal.2007).," The associated rotational deformation is a few percent, depending on the assumed mass and nuclear equation of state, leading to a small but noticeable influence on the light curve \citep{MLCB}."
. We neglect both special relativistic time delay and the additional time delay experienced. by initially αναον wopagatingoὃν photons. since these delays are much smaller (~101 s. Poutanen&Beloborodoy(2006))) than the time bins we consider (~0.1 8).," We neglect both special relativistic time delay and the additional time delay experienced by initially inwardly propagating photons, since these delays are much smaller $\sim 10^{-4}$ s, \citet{PBel}) ) than the time bins we consider $\sim 0.1$ s)."
 We specify stellar mass AZ and equatorial radius Ro and then compute the deformed spherical surface using the OS mocel., We specify stellar mass $M$ and equatorial radius $R_\mathrm{eq}$ and then compute the deformed spherical surface using the OS model.
 To start the burst we specify an initial small burning area and then track the propagation of the burning front across the star., To start the burst we specify an initial small burning area and then track the propagation of the burning front across the star.
 The stellar surface is divided into a grid of xuches with area ~ 0.1 kim. and we consider a patch to be ienited as soon as the burning front reaches the centre point of the patch.," The stellar surface is divided into a grid of patches with area $\sim$ 0.1 $^2$, and we consider a patch to be ignited as soon as the burning front reaches the centre point of the patch."
 Once a patch has started burning. we need o specify how its emission. varies with tine.," Once a patch has started burning, we need to specify how its emission varies with time."
 As discussed in Section L.. there are various numoerically-generated single »oint emission models. but no simple analvtie models.," As discussed in Section \ref{intro}, there are various numerically-generated single point emission models, but no simple analytic models."
" In his stuely we follow Bhattacharvva&Strohmaver(2006a.b) and assume that the temperature. of the burning front ollows the following profile after ignition: where /,, is the time at which the temperature reaches HS maximum d,=dug0.99(173du)."," In this study we follow \citet{BSa, BSb} and assume that the temperature of the burning front follows the following profile after ignition: where $t_m$ is the time at which the temperature reaches its maximum $T_m = T_0 + 0.99(T_1-T_0)$."
" Vhe time scale fiy ses the time scale at whichthe temperature increases. while /,, sets the time scale on which it cecavs."," The time scale $t_\mathrm{lr}$ sets the time scale at whichthe temperature increases, while $t_\mathrm{ld}$ sets the time scale on which it decays."
 Unburnt patches are assumed to have a temperature Zi until ignition., Unburnt patches are assumed to have a temperature $T_0$ until ignition.
 The xwameters in this model depend. primarily on. the composition of the burning maerial. which will vary with on accretion rate (see Section 1)).," The parameters in this model depend primarily on the composition of the burning material, which will vary with on accretion rate (see Section \ref{intro}) )."
 Bursts which are helium-rich. or example. would be expected to have shorter timescales han those which contain a higher fraction of hvdrogen.," Bursts which are helium-rich, for example, would be expected to have shorter timescales than those which contain a higher fraction of hydrogen."
 Our »arameter space must tjerefore. be wide enough to ake into account the expected level of variation., Our parameter space must therefore be wide enough to take into account the expected level of variation.
 One ollow-on question is whether his exponential temperature model remains valid across all burning regimes., One follow-on question is whether this exponential temperature model remains valid across all burning regimes.
 lt has oen used successfully in. detailed: spectral modelling of yUPStS at dilferent accretion rates by Bhattacharyya&Strohmaver (2006a.b).. but ultimately one would. like to see this confirmed by detailed nuclear physics caleulations.," It has been used successfully in detailed spectral modelling of bursts at different accretion rates by \citet{BSa,BSb}, but ultimately one would like to see this confirmed by detailed nuclear physics calculations."
 We then assume black body. emission at the specified temperature from cach patch., We then assume black body emission at the specified temperature from each patch.
 In Figure 7 we show the tvpical single pateh light curve., In Figure \ref{singlepoint} we show the typical single patch light curve.
 The rise portion is similar in shape to the bursts shown in. Woosleyctal.(2004)., The rise portion is similar in shape to the bursts shown in \citet{W}.
. The decay portion does dilfer from that seen in some of the cases studied by Woosleyetal. (2004).. but because we are focusing on the rise we never reach the points late in the," The decay portion does differ from that seen in some of the cases studied by \citet{W}, , but because we are focusing on the rise we never reach the points late in the"
requency distributions of the relative 5;dv colours of the 69 quasars in our sample. and the relative ον colours of the LBQS quasars in Figure S... a Ixuiper test. shower he two cistributions to be different at the confidence evel.,"frequency distributions of the relative $b_J - K$ colours of the 69 quasars in our sample, and the relative $b_J - K$ colours of the LBQS quasars in Figure \ref{relbjMkCDFs}, a Kuiper test showed the two distributions to be different at the confidence level."
" Phe statistics of the two relative 6,AN distributions are consistent with our quasar sample having a broader anc rededer relative byA distribution than the LBQS quasars.", The statistics of the two relative $b_J - K$ distributions are consistent with our quasar sample having a broader and redder relative $b_J - K$ distribution than the LBQS quasars.
 Dased on the dillerences between Figures 4. and 7.. and our statistical analyses here and in Section 3.2... we determine hat whether we used observed or relative quasar colours. he results of comparinge our quasar sample to a sample of LBQS quasars is the same. that our sample has a broader. redder 6.)AN distribution than the LBOS quasars.," Based on the differences between Figures \ref{bjMKhistogram} and \ref{relbjMkHist}, and our statistical analyses here and in Section \ref{comparison}, we determined that whether we used observed or relative quasar colours, the results of comparing our quasar sample to a sample of LBQS quasars is the same, that our sample has a broader, redder $b_J - K$ distribution than the LBQS quasars."
 The acditional information gleaned from our comparison of the relative 5; ἰν colours is that the small excess of blue quasars in our sample. when compared to the LDOS. appears to be caused. by. Ix-correction. effects. and is not. actually. an intrinsic property of the 5;ἐν distribution.," The additional information gleaned from our comparison of the relative $b_J - K$ colours is that the small excess of blue quasars in our sample, when compared to the LBQS, appears to be caused by K-correction effects, and is not actually an intrinsic property of the $b_J - K$ distribution."
 Because the excess of blue. quasars in our quasar sample. when compared to the LBQS quasar sample. is probably caused by WKeeorrection ellects. the small excess of blue quasars is much less significant than the excess of τος quasars when Comparing our quasar sample to the LBOS sample.," Because the excess of blue quasars in our quasar sample, when compared to the LBQS quasar sample, is probably caused by K-correction effects, the small excess of blue quasars is much less significant than the excess of red quasars when comparing our quasar sample to the LBQS sample."
" To understand. the effects of cur sample sclection on the b, and IX magnitude parameter space. we plottec 5; and"," To understand the effects of our sample selection on the $b_J$ and K magnitude parameter space, we plotted $b_J$ and"
 software. oue having 30 slitets and he other 27 slilets on M31 RGB star candidates. with the targets distribued more-or-less uuiforuilv over he LRIS feld «of view.," software, one having 30 slitlets and the other 27 slitlets on M31 RGB star candidates, with the targets distributed more-or-less uniformly over the LRIS field of view."
 Multi-slit spectroscopic observations were carried. usdic Week/LRIS chiving a wo-nigh run on Septeuiber 28.29. 1998 (UT).," Multi-slit spectroscopic observations were carried using Keck/LRIS during a two-night run on September 28–29, 1998 (UT)."
 Each spectiuu covers the spectral rauee AATHHO ssh0A conainiue t16 near-nfrarec triplet: AAS198. 8512. audAA.," Each spectrum covers the spectral range $\lambda\lambda7550\>$ $\>8850\,\rm\AA$ containing the near-infrared triplet: $\lambda\lambda8498$, 8542, and."
. The instruneutal spectral/velocity resolutions AA//GS km (FWIIAD: see ROO2 or detais., The instrumental spectral/velocity resolution is /68 km $^{-1}$ (FWHM); see RG02 for details.
 The total exposure times for the two masks is 1.7 aud 2.0 hr., The total exposure times for the two masks is 1.7 and 2.0 hr.
 Individual exposures are typicallv 30 1 lone. althougLsome exposures were stopped short due to telescope. iustmeut. and weather preleis.," Individual exposures are typically 30 min long, although some exposures were stopped short due to telescope, instrument, and weather problems."
 Overscan correction. 2D bias structure sitraction. and cosmic rav removal are accomplished using standard tasks.," Overscan correction, 2D bias structure subtraction, and cosmic ray removal are accomplished using standard tasks."
 Cosmic ravs are remove’ ou the basis of object sharpness aud peal pixel brightuess., Cosmic rays are removed on the basis of object sharpness and peak pixel brightness.
 Thev are nasked from cach image aloug with :Q summrouncdiue l-pixel buffer to remove the low-level wines of each eveut., They are masked from each image along with a surrounding 1-pixel buffer to remove the low-level wings of each event.
 A flat-ficld correction or each cata frame is performed using a spectral dome flat that is wellanatched to the data frame in terms of LRIS flexure effects., A flat-field correction for each data frame is performed using a spectral dome flat that is well-matched to the data frame in terms of LRIS flexure effects.
 Data reduction issues/complicatious for LRIS specPa atὉ cdiseussed in sole dezül iu RC2: that study did not however use the reduction software that is used here., Data reduction issues/complications for LRIS spectra are discussed in some detail in RG02; that study did not however use the \citet{phi03} reduction software that is used here.
 Wavelenetli calibration. skv subraction and extraction are all accomplished using the LRIS data reduction pipeline developed by Phillipsctal.(2000).," Wavelength calibration, sky subtraction and extraction are all accomplished using the LRIS data reduction pipeline developed by \citet{phi03}."
.. The software uses an optical model for the various ements of the LRIS spectrograph (colnuator. erating. canuera. ete.)," The software uses an optical model for the various elements of the LRIS spectrograph (collimator, grating, camera, etc.)"
 to derive a mapping from the slitinask to ic CCD detector as a fiction of waveleneth., to derive a mapping from the slitmask to the CCD detector as a function of wavelength.
 This optical model is based on spectrograpl design daawines. id has been empiricallv refined usinο calibration spectra (arc lamp through a ericd-of-holes mask) taken ‘lose to the observing run to account or nusaligninent of auv of LRIS’ clemenuts.," This optical model is based on spectrograph design drawings, and has been empirically refined using calibration spectra (arc lamp through a grid-of-holes mask) taken close to the observing run to account for misalignment of any of LRIS' elements."
 Even so the model is oulv ood to about 0.3 pixels (0.06 in the spatial direction and ~0.2A 1 the dispersion direction)., Even so the model is only good to about 0.3 pixels $0.06''$ in the spatial direction and $\sim$ in the dispersion direction).
 Iu addition. one iust allow for small time-dependent aliguiment/focus errors. id variations from exposure to exposure because of instruneut flexure.," In addition, one must allow for small time-dependent alignment/focus errors, and variations from exposure to exposure because of instrument flexure."
 These errors are removed by corrections: a zeropoint correction iu waveleueth measured from a bright uieht-kv. emission line. aud a plate-scale and offset which are solved for using the measured locii of slitlet edges.," These errors are removed by low-order corrections: a zeropoint correction in wavelength measured from a bright night-sky emission line, and a plate-scale and offset which are solved for using the measured locii of slitlet edges."
 Iu order to determine the radial velocity of cach object. its final coadded προςτα is cross-correlated against a tempate spectimi (average spectrin of three control sample stars from ROO2. where the coadditiou is done after süfting each to zero velocity).," In order to determine the radial velocity of each object, its final coadded spectrum is cross-correlated against a template spectrum (average spectrum of three control sample stars from RG02, where the coaddition is done after shifting each to zero velocity)."
 The cross-correlation fiction (CCE) is computed from 1000 O |1000 Iau 1. covering a plausible range of radial velocities for stars associated with ΑΠ).," The cross-correlation function (CCF) is computed from $-1000$ to +1000 km $^{-1}$, covering a plausible range of radial velocities for stars associated with M31."
 The CCF echuique vieks an unanibieuous peak and a reliable radial velocity for 11 of the 57 objects coiuprising the uain sample o| potential M31 targets., The CCF technique yields an unambiguous peak and a reliable radial velocity for 41 of the 57 objects comprising the main sample of potential M31 targets.
 A complete description of this CCF procedure is given d1i RGO2., A complete description of this CCF procedure is given in RG02.
 The radia velocity determined from the location of the CCF peak. ei. is corrected to the heliocentric yale using he task iu IRAE.," The radial velocity determined from the location of the CCF peak, $v_{\rm
obs}$, is corrected to the heliocentric frame using the task in IRAF."
 We estimate the ris error in radial velocity from t16 degree of sienificance of the CCF peak: σι=clIP(1το).b (Toury&Davis1979)., We estimate the rms error in radial velocity from the degree of significance of the CCF peak: $\sigma_v=\sigma_v^{\rm TD}(1+r_{\rm TD})^{-1}$ \citep{ton79}.
". The value of σTD"" Is enipirically ound to be 7T damus+ for ow ustimuental set-up (RGO2).", The value of $\sigma_v^{\rm TD}$ is empirically found to be 77 km $^{-1}$ for our instrumental set-up (RG02).
 The mean lo error in velocity for lis Cl fold salple is 21 xls1, The mean $1\sigma$ error in velocity for this G1 field sample is 21 km $^{-1}$ .
out” due to the nulling condition.,out” due to the nulling condition.
 If only a limited number of redshift bins is available. (20)) holds only approximately. leading to a residual 1n (21)).," If only a limited number of redshift bins is available, \ref{eq:GGIsplit}) ) holds only approximately, leading to a residual in \ref{eq:tomonull}) )."
" Since the nulling condition i$ the only condition. that the weight 7? must satisfy in order to ""null"". there is much freedom in choosing the form of it."," Since the nulling condition is the only condition that the weight $T^{(ij)}$ must satisfy in order to “null”, there is much freedom in choosing the form of it."
 We would like to further specify its form such that it preserves as much Fisher information in Y? as possible., We would like to further specify its form such that it preserves as much Fisher information in $Y^{(ij)}$ as possible.
 The method we have adopted for the nulling weight construction will be detailed in 4..," The method we have adopted for the nulling weight construction will be detailed in $\,$ \ref{sec:weight}."
 Note that for each (7.{) combination. one can in principle apply more than one nulling weight to the original bispectrum. and obtain more nulled measures.," Note that for each $(i,j)$ combination, one can in principle apply more than one nulling weight to the original bispectrum, and obtain more nulled measures."
 If one retains the condition of maximizing the Fisher information and demands that all the weight functions built for one (./) combination are orthogonal to each other. one arrives at higher-order modes that have the second-most. third-most. ete..," If one retains the condition of maximizing the Fisher information and demands that all the weight functions built for one $(i,j)$ combination are orthogonal to each other, one arrives at higher-order modes that have the second-most, third-most, etc.,"
 information content (higher-order weights. see JSO8).," information content (higher-order weights, see JS08)."
 The total number of such linearly independent nulled measures for a certain (7./) equals the possible values ofk>i1.," The total number of such linearly independent nulled measures for a certain $(i,j)$ equals the possible values of $k \ge i+1$."
 In this schematic study we will only use the optimum. re. the first-order nulling weights.," In this schematic study we will only use the optimum, i.e. the first-order nulling weights."
 We will assess the information loss due to this limitation in 5.4..," We will assess the information loss due to this limitation in $\,$ \ref{sec:compression}. ."
 We set up a fictitious survey with a survey size of A = 4000 deg? which is similar to the survey size of DES., We set up a fictitious survey with a survey size of $A$ = 4000 $^2$ which is similar to the survey size of DES.
 This can be easily scaled to any survey size using the proportionality of statistical errors to A7!7., This can be easily scaled to any survey size using the proportionality of statistical errors to $A^{-1/2}$.
" We assume a galaxy intrinsic ellipticity dispersion σε=|"")0.35.", We assume a galaxy intrinsic ellipticity dispersion $\sigma_{\epsilon}=\sigma(\epsilon_{\rm I}^{\rm ran})=0.35$.
 As galaxy redshift probability distribution we adopt the frequently used parameterization (Smailetal.1994).. and use zy=0.64. a=2. =1.5.," As galaxy redshift probability distribution we adopt the frequently used parameterization \citep{smail94}, and use $z_0 = 0.64$, $\alpha = 2$, $\beta = 1.5$."
 The distribution is cut at tna=3 and normalized to 1., The distribution is cut at $z_{\textrm{max}} = 3$ and normalized to 1.
 The corresponding median redshift of this fictitious survey is zi=0.9. which is compatible to a survey like EUCLID.," The corresponding median redshift of this fictitious survey is $z_{\rm m}=0.9$, which is compatible to a survey like EUCLID."
 We adopt an average galaxy number density 5.=40aremin. which ts again EUCLID-like., We adopt an average galaxy number density $\bar{n}_{\textrm{g}}=40~ {\textrm{arcmin}}^{-2}$ which is again EUCLID-like.
" Disjunct redshift bins without photo-z error are assumed. which means that the galaxy redshift probability distribution in redshift bin / takes the form p(2)ep(x) if and only if the redshift that corresponds to comoving distance y, is within the boundaries of redshift bin ;."," Disjunct redshift bins without photo-z error are assumed, which means that the galaxy redshift probability distribution in redshift bin $i$ takes the form $p^{(i)}_{\textrm{s}}(z) \propto p_{\textrm{s}}(z)$ if and only if the redshift that corresponds to comoving distance $\chi_{\rm s}$ is within the boundaries of redshift bin $i$."
 A number of 10 redshift bins are used by default., A number of 10 redshift bins are used by default.
 The boundaries of the redshift bins are set such that each bin contains the same number of galaxies., The boundaries of the redshift bins are set such that each bin contains the same number of galaxies.
 We adopt 20 angular frequency bins spaced logarithmically between Gyin=50 and ἕμ=3000. and denote the characteristic angular frequency of a bin as £.," We adopt 20 angular frequency bins spaced logarithmically between $\ell_{\textrm{min}}=50$ and $\ell_{\textrm{max}}=3000$, and denote the characteristic angular frequency of a bin as $\bar{\ell}$."
 Within this range the noise properties of the cosmic shear field are still not too far in the non-Gaussian regime. allowing à more realistic theoretical estimation. of the bispectrum and its covariance.," Within this range the noise properties of the cosmic shear field are still not too far in the non-Gaussian regime, allowing a more realistic theoretical estimation of the bispectrum and its covariance."
 Whether this number of angular frequency bins can reconstruct the angular frequency dependence of the bispectrum is tested. and 20 bins are found to be sufficient for our requirements on precision.," Whether this number of angular frequency bins can reconstruct the angular frequency dependence of the bispectrum is tested, and 20 bins are found to be sufficient for our requirements on precision."
 This is also expected since thebispectrum ts rather featureless as a function of angular frequency., This is also expected since thebispectrum is rather featureless as a function of angular frequency.
 We show the modeling of Boog and its covariance in this section., We show the modeling of $B_{\textrm{GGG}}$ and its covariance in this section.
 We will only consider the tomographic bispectrum at redshift bins satisfying z;< and z;<z. which already ensures an elimination of By; and Bey systematics in our case.," We will only consider the tomographic bispectrum at redshift bins satisfying $z_i<z_j$ and $z_i<z_k$ , which already ensures an elimination of $B_{\textrm{III}}$ and $B_{\textrm{GII}}$ systematics in our case."
" Applying Limber’s equation. it can be shown that the tomographic convergence bispectrum can be written as à projection of the three-dimensional matter bispectrum By(ki.koΚαὶy) (see e.g. TIOF): To compute B,. we employ the fitting formula by Scoceimarro&Couchman(2001)... which is based on hyper-extended perturbation theory (Scoccimarro&Frieman 1999)."," Applying Limber's equation, it can be shown that the tomographic convergence bispectrum can be written as a projection of the three-dimensional matter bispectrum $B_\delta\!\left({k_1},{k_2},{k_3}; \chi\right)$ (see e.g. TJ04): To compute $B_{\delta}$, we employ the fitting formula by \citet{sco01}, which is based on hyper-extended perturbation theory \citep{sco99}. ."
 A comparison of this formula with the halo model results can be found in Takada&Jain(2003a.b).," A comparison of this formula with the halo model results can be found in \citet{TJ03b,TJ03c}."
". Estimating the bispectrum covariance is often done within a flat-sky spherical harmonie. formalism (Hu2000.Hu00Ohereafter)... which suffers - at least formally - from drawbacks since its basis functions Yj, are only defined for discrete angular frequency values and full sky coverage."," Estimating the bispectrum covariance is often done within a flat-sky spherical harmonic formalism \citep[Hu00 hereafter]{hu00}, which suffers - at least formally - from drawbacks since its basis functions $Y_{lm}$ are only defined for discrete angular frequency values and full sky coverage."
 In Joachimietal.(2009) another approach exclusively based on the two- Fourier formalism was constructed. which we use here.," In \citet{joachimi09b} another approach exclusively based on the two-dimensional Fourier formalism was constructed, which we use here."
 The bispectrum covariance is à six-point correlation function which can be expanded into its connected parts as outlined in e.g. Bernardeauetal.(2002a)., The bispectrum covariance is a six-point correlation function which can be expanded into its connected parts as outlined in e.g. \citet{ber02a}.
. As argued in TJOA. for angular scales £€3000. the term which is a triple of two-point functions still dominates.," As argued in TJ04, for angular scales $\ell \le 3000$, the term which is a triple of two-point functions still dominates."
" Keepingonly this term. the bispectrum covariance readsin which AZ; is the bin width of the angular frequency bin with typical value £;. and P""(£) is theobserved powerspectrum which contains the intrinsic ellipticity noise (e.g.Kaiser 2008): where i; is the galaxy number density in redshift bin {."," Keepingonly this term, the bispectrum covariance readsin which $\Delta \bar{\ell}_i$ is the bin width of the angular frequency bin with typical value $\bar{\ell}_i$ , and $\bar{P}^{(ij)}(\bar{\ell})$ is theobserved powerspectrum which contains the intrinsic ellipticity noise \citep[e.g.][]{kaiser92,hu99,JS08a}: : where $\bar{n}_i$ is the galaxy number density in redshift bin $i$ ."
 The term A(£j.£5.(3) 18 defined as," The term $\Lambda \br{\ell_1,\ell_2,\ell_3}$ is defined as \eqa{"
 Anomalous Microwave Emission ((AME) is the name given to excess microwave emission. observed. at. frequencies in the range ~1060 CGllz. that is stronely correlated with fr infrared. (ELI) dust. emission2010).," Anomalous Microwave Emission (AME) is the name given to excess microwave emission, observed at frequencies in the range $\sim10-60$ GHz, that is strongly correlated with far infrared (FIR) dust emission."
. This dust-correlatecd emission is known to be a significant source of contamination [or Cosmic Microwave jJackground (CMD). data. that must be separated: accurately from the CAIB signal (Bennett2008:Goldetal.2009:Dickinson 2009b).," This dust-correlated emission is known to be a significant source of contamination for Cosmic Microwave Background (CMB) data, that must be separated accurately from the CMB signal \citep{Bennett03,Bonaldi07,Miville-Deschenes08,Gold09,Dickinson09b}."
. Although there is still some debate about the physical mechanism that is responsible for the emission. the most. favoured explanation is in terms of small. rapidly spinning cust grains (Draine&Lazarian1998a.b).," Although there is still some debate about the physical mechanism that is responsible for the emission, the most favoured explanation is in terms of small, rapidly spinning dust grains \citep{Draine98a,Draine98b}."
. Assumine this is the case. microwave observations of spinning dust represent à new wav of studying the properties of interstellar dust. grains and its environment within the interstellar medium," Assuming this is the case, microwave observations of spinning dust represent a new way of studying the properties of interstellar dust grains and its environment within the interstellar medium"
. Assumine this is the case. microwave observations of spinning dust represent à new wav of studying the properties of interstellar dust. grains and its environment within the interstellar medium.," Assuming this is the case, microwave observations of spinning dust represent a new way of studying the properties of interstellar dust grains and its environment within the interstellar medium"
The A-values and Auger widths for the Ar isonuclear sequence previously computed by Diémontetal.(2000) with and the MCDE code have given us the opportunity {ο benchmark the accuracy of the present data sets.,The $A$ -values and Auger widths for the Ar isonuclear sequence previously computed by \citet{bie00} with and the MCDF code have given us the opportunity to benchmark the accuracy of the present data sets.
 This has been useful to sort oul the large discrepancies reported by Palmerietal.(2008a) for Fe and Palmeretal.(2006). [or 5 with the MCDE data of Chenetal.(1997)., This has been useful to sort out the large discrepancies reported by \citet{pal03a} for Fe and \citet{pal06} for S with the MCDF data of \citet{che97}.
. Similar discrepancies are found with (their data for Ar which. in the light of the good agreement of HERI with Diémmont et al..," Similar discrepancies are found with their data for Ar which, in the light of the good agreement of HFR1 with Biémmont et al.,"
 are cerlainiv not due to the more formal relativistic representation of the MCDE method., are certainly not due to the more formal relativistic representation of the MCDF method.
 Sienilieant cliscords of FRI with the Auger widths reported by Faenovetal.(1994) [or the Ale. Si. aud 5 isonuclear sequences and with those by Behar&Netzer(2002) [or levels in the D and C isoelectronic sequences are difficult to explain.," Significant discords of HFR1 with the Auger widths reported by \citet{fae94}
 for the Mg, Si, and S isonuclear sequences and with those by \citet{beh02} for levels in the B and C isoelectronic sequences are difficult to explain."
 Accuracy ratings of the present A-values and Auger widths greater than LOM ! are conticently assigned (to and20'4... respectively. except lor special cases: e.g. transitions involving Ix-vacancy levels in the carbon isoelectronic sequence where strong admixture causes severe departures.," Accuracy ratings of the present $A$ -values and Auger widths greater than $10^{13}$ $^{-1}$ are confidently assigned to and, respectively, except for special cases; e.g. transitions involving K-vacancy levels in the carbon isoelectronic sequence where strong admixture causes severe departures."
 The present radiative and Auger widths will be used in the computations of the IX-shell photoionization cross sections of these medium-Z ions which are required in NSTAR for the modelling of interesting spectral features., The present radiative and Auger widths will be used in the computations of the K-shell photoionization cross sections of these $Z$ ions which are required in XSTAR for the modelling of interesting spectral features.
 Future work will involve extension of the present approach to the Ni isonuclear sequence., Future work will involve extension of the present approach to the Ni isonuclear sequence.
 ALIAD acknowledges partial support from FONACIT. Venezuela. under contract No. 20011000912.," MAB acknowledges partial support from FONACIT, Venezuela, under contract No. S1-20011000912."
 This work was funded in part by the NASA Astronomy and Physics Research and Analysis Program., This work was funded in part by the NASA Astronomy and Physics Research and Analysis Program.
 Partial support for CAL was provided by a travel grant. [rom IVIC and the FNRS of Belgium., Partial support for CM was provided by a travel grant from IVIC and the FNRS of Belgium.
ADI.200L).. similar to the Large Magellanic Cloud.,"$_{\odot}$, similar to the Large Magellanic Cloud."
 It is oue of the nearest. examples of a starburst. and was the host of a Type I supernova ancl recent nova2010).," It is one of the nearest examples of a starburst, and was the host of a Type I supernova and recent nova."
. It therefore has been studied im detail. especially at UV and NIR waveleneths that probe the voune stellar populations.," It therefore has been studied in detail, especially at UV and NIR wavelengths that probe the young stellar populations."
 For example. were able to make use of &£u-UV imagine from the Ultraviolet Imaging Telescope aud ground-based Ebaud to decompose the galaxy iuto a centrally concentrated intense star forming coniponent and an extended soothi disk.," For example, were able to make use of far-UV imaging from the Ultraviolet Imaging Telescope and ground-based I-band to decompose the galaxy into a centrally concentrated intense star forming component and an extended smooth disk."
 They sugeest the two component structure is the result of a recent merger., They suggest the two component structure is the result of a recent merger.
 The buelt resolved stellar populations lave also been characterized., The bright resolved stellar populations have also been characterized.
 A dominant red giant brauch was detected. and a large metallicity spread was 1ieasured by using WEPC2 aud NICAIOS photometry.," A dominant red giant branch was detected, and a large metallicity spread was measured by using WFPC2 and NICMOS photometry."
 used UST/WEPC2 and ΠοΤΙΣ Πασάς to sugeest that the stellar initial mass functiou appears to be steeper than that of at high masses (£720 ΔΕ.) even though the ealaxw’s metallicity is similar to that of the Simall Magellanic Cloud.," used HST/WFPC2 and HST/STIS imaging to suggest that the stellar initial mass function appears to be steeper than that of at high masses $>$ 20 $_{\odot}$ ), even though the galaxy's metallicity is similar to that of the Small Magellanic Cloud."
 However. it is unclear how this slope is affected by fluctuations in the star formation rate.," However, it is unclear how this slope is affected by fluctuations in the star formation rate."
 We Πάτμο examine the stellar populations of NGC 1211 using deeper data than previous studies., We further examine the stellar populations of NGC 4214 using deeper data than previous studies.
 The jw data include two DUST WFPC2 fields of the outer disk aud one WECS3 field of the immer regions., The new data include two HST WFPC2 fields of the outer disk and one WFC3 field of the inner regions.
 Through detailed. stellar evolution iuodel fitting. we measure he star formation history (SEII) of these regions ik ook for differences between quiesceut and intensely star ornüue regions.," Through detailed stellar evolution model fitting, we measure the star formation history (SFH) of these regions and look for differences between quiescent and intensely star forming regions."
 Section 2 describes our data set am xiotonmetry., Section 2 describes our data set and photometry.
 Section 3 prescuts the results of our mode chting procedure., Section 3 presents the results of our model fitting procedure.
 Section [| discusses the results of the ucasureineuts m the context of galaxy morphology iik euvironment. aud Section 5 sumunarizes our conclusions.," Section 4 discusses the results of the measurements in the context of galaxy morphology and environment, and Section 5 summarizes our conclusions."
 Weasstune (0)Mjoy= 3.03 Mpc Or conversions of angular measurements to plivsica distances and assume an inclination of 38° for surface density measurements., We assume $(m-M)_0 =$ 3.03 Mpc for conversions of angular measurements to physical distances and assume an inclination of $^{\circ}$ for surface density measurements.
 We adopt five-vear WMAP cosiuologv. for ] conversions between time and redshift., We adopt a five-year WMAP cosmology for all conversions between time and redshift.
 As part of the ANGST. program after the fülure of ACS. from 2007-Dec-0L1 to 2007-Dec-23. we performed deep WFPC? observations a field in the NGC 1211 disk located at R.À. (2000)—182.817375. (12:15:23.10). decl. (," As part of the ANGST program after the failure of ACS, from 2007-Dec-04 to 2007-Dec-23, we performed deep WFPC2 observations a field in the NGC 4214 disk located at R.A. (2000)=183.847375 (12:15:23.4), decl. ("
2000)=36.362333 (136:21: with a rotation angle PA_VV3=119.9 degrees (GO-10915).,2000)=36.362333 (+36:21:44) with a rotation angle V3=119.9 degrees (GO-10915).
.L1) To. iniprove our radial coverage to be closer to the goal of the original ANGST program. from 2009-Feb-23 to 2009-Feb-26. we performed shallower supplemental observations for a field located at R.A. (2000)=183.878103 (12:15:30.7). decl. (," To improve our radial coverage to be closer to the goal of the original ANGST program, from 2009-Feb-23 to 2009-Feb-26, we performed shallower supplemental observations for a field located at R.A. (2000)=183.878103 (12:15:30.7), decl. ("
2000)236.356808 (1236:21:21.5). with a rotation anele PA_-VV3=51.0 (GO-11986).,2000)=36.356808 (+36:21:24.5) with a rotation angle V3=51.0 (GO-11986).
 At the same time. WEC3/IR and WEC3/UVIS imaging of the central area of the ealaxy. proposed by the Scicutific Oversight Committee (SOC). were released.," At the same time, WFC3/IR and WFC3/UVIS imaging of the central area of the galaxy, proposed by the Scientific Oversight Committee (SOC), were released."
 Both fields were observed from) 2009-Dec-22 to 2009-Dec-23 located. at R.A. (2000) =183.913333 (12:15:39.2). decl. (," Both fields were observed from 2009-Dec-22 to 2009-Dec-23 located at R.A. (2000) =183.913333 (12:15:39.2), decl. ("
2000) = 36.226911. (126:19:297.0). with a rotation anele VV3=120.0 (CGO-11360. PE: O'Connell).,"2000) = 36.326944 (+36:19:37.0) with a rotation angle V3=120.0 (GO-11360, PI: O'Connell)."
 Figure 1 shows outlines of the fields’ locations.," Figure \ref{field_loc}
 shows outlines of the fields' locations."
 Our field locations were chosen to maximize the uuuber of disk stars aud avoid crowding., Our field locations were chosen to maximize the number of disk stars and avoid crowding.
 Iu the deep field. we obtained 15 full-orbit exposures with the ΝΕΤΟΖ through the F606W (wide V) filter. aud 29 full-orbit exposures through the ESIWV (7 equivalent) filter.," In the deep field, we obtained 15 full-orbit exposures with the WFPC2 through the F606W (wide $V$ ) filter, and 29 full-orbit exposures through the F814W $I$ equivalent) filter."
 These data totaled 39000 s and 75100 s of exposure time in F606W and FSIIWNV. respectively.," These data totaled 39000 s and 75400 s of exposure time in F606W and F814W, respectively."
 In the other WFDPC? field. we obtained 2 orbits through FG6OGW. totaling [800 s. aud. { orbits through ESLIW. totaling 9600 s. The WEC3/TR data contained 1198 s aud 2308 s of exposure in FLIOW aud. F160W. respectively.," In the other WFPC2 field, we obtained 2 orbits through F606W, totaling 4800 s, and 4 orbits through F814W, totaling 9600 s. The WFC3/IR data contained 1198 s and 2398 s of exposure in F110W and F160W, respectively."
" The WEC3/UVIS data contained 1683 s. 1510 s. aud 1339 « in F336W. EDW. and FaliW, respectively."," The WFC3/UVIS data contained 1683 s, 1540 s, and 1339 s in F336W, F438W, and F814W, respectively."
 All WEPC2 taecs were calibrated in the IST pipeline ith CALWDP?2 using OPUS version 2006_66a for the 2007 data and 2008_55¢ for the 2009 data., All WFPC2 images were calibrated in the HST pipeline with CALWP2 using OPUS version 6a for the 2007 data and 5c for the 2009 data.
 All WEC'3 images were calibrated in the IST pipeline with CALWE3 version 2.0., All WFC3 images were calibrated in the HST pipeline with CALWF3 version 2.0.
 All of the data were processed through the ACS Nearby Calaxy Survey Treasury (ANGST) data analysis pipeline2009).. updated to include WEC3 UVIS and IR data2011).," All of the data were processed through the ACS Nearby Galaxy Survey Treasury (ANGST) data analysis pipeline, updated to include WFC3 UVIS and IR data."
. As a brief sumunary. the photometry was measured sinultaueouslwv for all of the objects in the uncombiued inages using the software packages IISTPIIOT (for WFPC2) and DOLPIIOT 1.2 including the ΝΕΟ module.," As a brief summary, the photometry was measured simultaneously for all of the objects in the uncombined images using the software packages HSTPHOT (for WFPC2) and DOLPHOT 1.2 including the WFC3 module."
 These packages are optimized for 1icasuriug photometry of stars on DIST tages using the well-characterizecd aud stable point spread function (PSF) calculated with The software fits the PSF to all of the stars in cach individual frame to fud PSF magnitudes., These packages are optimized for measuring photometry of stars on HST images using the well-characterized and stable point spread function (PSF) calculated with The software fits the PSF to all of the stars in each individual frame to find PSF magnitudes.
 It then determines aud applies the aperture correction for cach iuase uxing the most isolated stars. corrects for the charec transfer efficiency of the WEPC?detector®.. conibines the results frou the individual exposures. aud converts the measured count rates to the VECAmae svsteni.," It then determines and applies the aperture correction for each image using the most isolated stars, corrects for the charge transfer efficiency of the WFPC2, combines the results from the individual exposures, and converts the measured count rates to the VEGAmag system."
 Our photometric error aud completeness were then assessed. by running at least LO° artificial star tests for each field. in which a star of known color and magnitude was placed iuto the images aud the photometry rerun to determine if the star was recovered aud. if so. the difference between the iiput aud output magnitudo.," Our photometric error and completeness were then assessed by running at least $^6$ artificial star tests for each field, in which a star of known color and magnitude was placed into the images and the photometry rerun to determine if the star was recovered and, if so, the difference between the input and output magnitude."
 The photometry output was then filtered to ouly allow objects classifies as stars with signal-to-noise > lain both filters., The photometry output was then filtered to only allow objects classified as stars with signal-to-noise $>$ 4 in both filters.
" The list was further culled using sharpucss μαper,m1BSLUVspar)<<0.27 for WEPC2. [E336]aarp|BIBSως0.27 for UVIS. and [PILave,|F1601Tery|(0.35 for IR) anc crowding (F606ord|PSLi<0.7 for WEPC2. Puord|PIBoe0.7 for UVIS. anc PUNod|EIT601T.4<0.18 for IR)."," The list was further culled using sharpness $|F606W_{sharp} +
F814W_{sharp}| < 0.27$ for WFPC2, $|F336W_{sharp} + F438W_{sharp}| <
0.27$ for UVIS, and $|F110W_{sharp} + F160W_{sharp}| < 0.35$ for IR) and crowding $F606W_{crowd} + F814W_{crowd} < 0.7$ for WFPC2, $F336W_{crowd} + F438W_{crowd} < 0.7$ for UVIS, and $F110W_{crowd} +
F160W_{crowd} < 0.48$ for IR)."
 The sharpuess xuwanmeter returus zero if a star is porfectlv fit. negative if it is broader. aud positive if it is sharper than a tvpica star2000).," The sharpness parameter returns zero if a star is perfectly fit, negative if it is broader, and positive if it is sharper than a typical star."
. The crowding parameter gives the difference between the magnitude of a star measure: vefore and after subtracting the neieliboring stars in the nuaec., The crowding parameter gives the difference between the magnitude of a star measured before and after subtracting the neighboring stars in the image.
 When this value is large. it sugeests that the stars photometry was siguificautly affected by crowding. auc," When this value is large, it suggests that the star's photometry was significantly affected by crowding, and"
power spectrum.,power spectrum.
 When (he covariance matrix is used. we have six parameters.," When the covariance matrix is used, we have six parameters."
 We use (he Monte-Carlo Markov. Chain (MCMC) method to explore the parameter space., We use the Monte-Carlo Markov Chain (MCMC) method to explore the parameter space.
 Our MCMC code is based on the publicly available package COSMOAIC (Lewis&Bridle2002)., Our MCMC code is based on the publicly available package COSMOMC \citep{cosmomc}.
. The paper is organized as follows., The paper is organized as follows.
" In section Il. we give all (he formulae and show the constraint on the cosmic curvature is much better by using the parameters A. /, and (heir covariance malrix (an (hat by using the parameter 7? only."," In section II, we give all the formulae and show the constraint on the cosmic curvature is much better by using the parameters $R$, $l_a$ and their covariance matrix than that by using the parameter $R$ only."
" We also discuss the effect of the radiation component Q, on /,.", We also discuss the effect of the radiation component $\Omega_r$ on $l_a$.
 In section III. we eive our results.," In section III, we give our results."
" We discuss (he analvtical marginalization over //j in appendix A. For the SN Ia data.we caleulate where the extinction-corrected distance modulus (2)=5logy{dp(2)/Alpe}]+25. σι is the total uncertainty in the SN Ia data. and the luminosity distance is here and the dimensionless Hubble parameter is where Ὁ=8zCp/(3Iz). p,=oT}emb* oy is the Stelan-Doltzmann constant. the CMD temperature 7,,,)= 2.726). and Qp pis the DE density."," We discuss the analytical marginalization over $H_0$ in appendix A. For the SN Ia data,we calculate where the extinction-corrected distance modulus $\mu(z)=5\log_{10}[d_L(z)/{\rm Mpc}]+25$, $\sigma_i$ is the total uncertainty in the SN Ia data, and the luminosity distance is here and the dimensionless Hubble parameter is where $\Omega=8\pi G\rho/(3H^2_0)$, $\rho_r=\sigma_bT_{cmb}^4$, $\sigma_b$ is the Stefan-Boltzmann constant, the CMB temperature $T_{cmb}=2.726$ K, and $\Omega_{DE}$ is the DE density."
" Note (hat the clistance normalization is arbitrary in the SN Ia data. the Hubble constant //, determined [rom the SN data is also an arbitrary number. not the observed IInbble constant."," Note that the distance normalization is arbitrary in the SN Ia data, the Hubble constant $H_0$ determined from the SN data is also an arbitrary number, not the observed Hubble constant."
 Therefore we need to marginalize over (his nuisance parameter f/f)., Therefore we need to marginalize over this nuisance parameter $H_0$ .
" The parameter Jf, is mareinalized over with flat prior. the"," The parameter $H_0$ is marginalized over with flat prior, the"
planetesimals shape the protoplanetary disc.,planetesimals shape the protoplanetary disc.
" Kokubo Ida (1998)) coined this growth regime ""oligarchie growth"". “in the sense that not only one but several protoplanets dominate the planetesimal system’."," Kokubo Ida \cite{kokubo1}) ) coined this growth regime “oligarchic growth”, `in the sense that not only one but several protoplanets dominate the planetesimal system'."
 Kokubo Ida (2000)) investigated. through 3D N-body simulations. the growth from planetesimals to. protoplanets including the effect of the nebular gas drag.," Kokubo Ida \cite{kokubo2}) ) investigated, through 3D N–body simulations, the growth from planetesimals to protoplanets including the effect of the nebular gas drag."
 They confirmed the existence of an initial runaway phase in protoplanetary growth and a second stage of oligarchic growth of protoplanets. in the same sense as in Kokubo Ida (1998)).," They confirmed the existence of an initial runaway phase in protoplanetary growth and a second stage of oligarchic growth of protoplanets, in the same sense as in Kokubo Ida \cite{kokubo1}) )."
 Also. when analysing the evolution of the RMS eccentricity and inclination of the planetesimal system during the oligarchic growth stage they found that their results agree with those predicted by the semi- theory of Ida Makino (1993)) for their second stage.," Also, when analysing the evolution of the RMS eccentricity and inclination of the planetesimal system during the oligarchic growth stage they found that their results agree with those predicted by the semi--analytical theory of Ida Makino \cite{ida}) ) for their second stage."
 In the present version of the code we have introduced a time-dependent planetesimal aecretion. rate., In the present version of the code we have introduced a time–dependent planetesimal accretion rate.
 In this kind of calculation most authors (Pollack et al. 1996:;, In this kind of calculation most authors (Pollack et al. \cite{pollack};
 Hubickyy et al. 2005::, Hubickyj et al. \cite{hubickyj};
 Alibert et al. 2005)), Alibert et al. \cite{alibert}) )
 usually prescribe that obtained by Greenzweig Lissauer (1992)) which assumes a rapid growth regime for the core., usually prescribe that obtained by Greenzweig Lissauer \cite{greenzweig}) ) which assumes a rapid growth regime for the core.
 Instead. we adopt that corresponding to the oligarchie growth of Ida Makino (1993)): a slower aceretion rate that still has not been explored with a self-consistent code for giant planet formation.," Instead, we adopt that corresponding to the oligarchic growth of Ida Makino \cite{ida}) ); a slower accretion rate that still has not been explored with a self–consistent code for giant planet formation."
 The condition for the dominance of the oligarchic growth over the (previous) runaway growth of a protoplanet. M/ii=50— 100. was derived semi-analytically by Ida Makino (1993)).," The condition for the dominance of the oligarchic growth over the (previous) runaway growth of a protoplanet, $M/m \simeq 50-100$ , was derived semi–analytically by Ida Makino \cite{ida}) )."
 In all the cases of interest for this study. the cross-over mass Is very low — 1.8. some orders of magnitude below the Earth mass (Thommes et al. 2003)).," In all the cases of interest for this study, the cross–over mass is very low — i.e. some orders of magnitude below the Earth mass (Thommes et al. \cite{thommes}) )."
 Due to the initial runaway regime. the cross-over mass is reached m a negligible time and thus. the formation time of a giant planet's core 15 almost entirely regulated by the oligarchic growth.," Due to the initial runaway regime, the cross–over mass is reached in a negligible time and thus, the formation time of a giant planet's core is almost entirely regulated by the oligarchic growth."
 For this reason. we prescribe the oligarchic growth for the core since the very beginning of our simulations.," For this reason, we prescribe the oligarchic growth for the core since the very beginning of our simulations."
" In the dispersion-dominated regime. a solid embryo growth rate 1s well described by the particle-in-a-box approximation (Safranov. 1969)). where M, is the mass of the solid protoplanet (in. our case. the core of the giant planet). / is the planetesimals? dise scale height. Rey is the effective capture radius and F is a factor mtroduced to compensate for the underestimation of the accretion rate by a two-body algorithm when considering the velocity dispersion of a population of planetesimals modelled by a single eccentricity and inclination equal to the RMS values."," In the dispersion–dominated regime, a solid embryo growth rate is well described by the particle–in–a–box approximation (Safranov, \cite{safranov}) ), where $M_\mathrm{c}$ is the mass of the solid protoplanet (in our case, the core of the giant planet), $h$ is the planetesimals' disc scale height, $R_\mathrm{eff}$ is the effective capture radius and $F$ is a factor introduced to compensate for the underestimation of the accretion rate by a two–body algorithm when considering the velocity dispersion of a population of planetesimals modelled by a single eccentricity and inclination equal to the RMS values."
 F is estimated to be ~3 (Greenzweig Lissauer 1992))., $F$ is estimated to be $\sim 3$ (Greenzweig Lissauer \cite{greenzweig}) ).
" Due to gravitational focusing. the effective capture radius of a protoplanet. Ray. is larger than its geometrical radius with A. the geometrical radius of the solid embryo. v4. the escape velocity from its surface and οι the relative velocity between the protoplanet and planetesimals. where. hereafter. ¢=(62)p αἱ=(i2374) is planetesimals RMS eccentricity (inclination) with respect to the dise (e. / << 1) and O, is the Keplerian angular velocity."," Due to gravitational focusing, the effective capture radius of a protoplanet, $R_\mathrm{eff}$, is larger than its geometrical radius with $R_\mathrm{c}$ the geometrical radius of the solid embryo, $v_\mathrm{esc}$ the escape velocity from its surface and $v_\mathrm{rel}$ the relative velocity between the protoplanet and planetesimals, where, hereafter, $e=\mean{e_m^2}^{1/2}$ $i= \mean{i_m^2}^{1/2}$) is planetesimals RMS eccentricity (inclination) with respect to the disc $e$, $i$ $<<$ 1) and $\Omega_\mathrm{k}$ is the Keplerian angular velocity."
 We apply the approximations 7=e/2 and ἡ=ai., We apply the approximations $i \simeq e/2$ and $h \simeq ai$.
 Following Thommes et al. (2003)).," Following Thommes et al. \cite{thommes}) ),"
" we adopt for e the equilibrium expression that is deduced for the case when gravitational perturbations due to the protoplanet are balanced by dissipation due to gas drag. where M is the protoplanet mass (here we assume M to be the total mass of the proto-giant planet. meaning the core mass plus the envelope mass). oj is the planetesimal bulk density. p is the gas volume density of the protoplanetary disc. Cp is the drag coefficient (dimensionless and of the order of 1). M, is the mass of the central star and £ is the width of the “heated region"" in units of the Hill radius (considering that potentially other embryos are growing as well. 6 is of the order of 10)."," we adopt for $e$ the equilibrium expression that is deduced for the case when gravitational perturbations due to the protoplanet are balanced by dissipation due to gas drag, where $M$ is the protoplanet mass (here we assume $M$ to be the total mass of the proto–giant planet, meaning the core mass plus the envelope mass), $\rho_m$ is the planetesimal bulk density, $\rho$ is the gas volume density of the protoplanetary disc, $C_\mathrm{D}$ is the drag coefficient (dimensionless and of the order of 1), $M_{\star}$ is the mass of the central star and $\beta$ is the width of the “heated region” in units of the Hill radius (considering that potentially other embryos are growing as well, $\beta$ is of the order of 10)."
 For calculations that do not assume equilibrium values for the eccentricity and the inclination. we refer the reader to Chambers (2006)).," For calculations that do not assume equilibrium values for the eccentricity and the inclination, we refer the reader to Chambers \cite{chambers}) )."
 His results show that using equilibrium values for e and / is an acceptable approximation when considering embryos of a2107?M.. consistent with the initial core mass of all our simulations (see Sect. 3)).," His results show that using equilibrium values for $e$ and $i$ is an acceptable approximation when considering embryos of $m \ga 10^{-2} \; \mathrm{M_{\oplus}}$, consistent with the initial core mass of all our simulations (see Sect. \ref{sec:results}) )."
 However. when the solid embryo is massive enough to gravitationally bind gas from the surrounding nebula. the presence of this envelope should be considered when calculating the effective capture radius. Ra.," However, when the solid embryo is massive enough to gravitationally bind gas from the surrounding nebula, the presence of this envelope should be considered when calculating the effective capture radius, $R_\mathrm{eff}$."
 The gaseous envelope modifies the trajectory of incoming planetesimals. as they are affected by the gas drag that enlarges the capture radius of the protoplanet.," The gaseous envelope modifies the trajectory of incoming planetesimals, as they are affected by the gas drag that enlarges the capture radius of the protoplanet."
" In addition to the gravitational focusing. a ""viscous focusing"" due to gas drag should be considered."," In addition to the gravitational focusing, a “viscous focusing” due to gas drag should be considered."
 The effective radius Ryy. 1n. the form stated in Eq. (115) ," The effective radius $R_\mathrm{eff}$, in the form stated in Eq. \ref{eq:reff}) )"
dominates when the mass of bound gas tis negligible., dominates when the mass of bound gas is negligible.
 But when the embryo acquires enough gas to form a thir atmosphere. its effective radius becomes larger and it separates from that caleulated previously.," But when the embryo acquires enough gas to form a thin atmosphere, its effective radius becomes larger and it separates from that calculated previously."
 To caleulate the effective radius of the protoplanet in the presence of gas around the solid embryo we take into account. the action of gravity and gas drag on. the incoming planetesimals., To calculate the effective radius of the protoplanet in the presence of gas around the solid embryo we take into account the action of gravity and gas drag on the incoming planetesimals.
" Consider one planetesimal entering the protoplanet atmosphere with a velocity ve,(Eq. (12))).", Consider one planetesimal entering the protoplanet atmosphere with a velocity $v_\mathrm{rel}$(Eq. \ref{eq:vrel}) )).
 Its equation of motion results from the action of gravity together with the action of gas drag., Its equation of motion results from the action of gravity together with the action of gas drag.
 In Eq. (14)).," In Eq. \ref{eq:grav}) ),"
" ¢ ts the radial coordinate from the centre of the protoplanet and M, is the mass contained within r.", $r$ is the radial coordinate from the centre of the protoplanet and $M_r$ is the mass contained within $r$ .
" For the gas drag force acting on a spherical body of radius 7, travelling with velocity v. we adopt the Stokes law"," For the gas drag force acting on a spherical body of radius $r_m$ travelling with velocity $v$ , we adopt the Stokes law"
The 30 Doradus nebula in the Large Alagellanie Cloucl (LAIC) is the closest example of a giant extragalactic rregion and the largest in the local group of galaxies.,The 30 Doradus nebula in the Large Magellanic Cloud (LMC) is the closest example of a giant extragalactic region and the largest in the local group of galaxies.
" Lt is regarded as undergoing intense enough star formation to be referred to as a ""mini-starburst, by Leitherer(1998). ancl as such is an important nearby laboratory of both massive star formation and starburst. phenomena.", It is regarded as undergoing intense enough star formation to be referred to as a `mini-starburst' by \citet{leitherer98} and as such is an important nearby laboratory of both massive star formation and starburst phenomena.
 The highly: dynamic nebulosity (e.g. Meaburn1981:Meaburn. 1987)) is powered by a super star cluster of LOO massive stars.," The highly dynamic nebulosity (e.g. \citealt{meaburn81,meaburn87}) ) is powered by a super star cluster of $\sim 100$ massive stars."
 Remarkable UST imagery of the environment. of the central cluster ol massive stars has recently: been presented by Walbornetal. (2002)., Remarkable HST imagery of the environment of the central cluster of massive stars has recently been presented by \citet{walborn.et.al02}.
.. Phe combined. winds. UW radiation and supernova explosions [roni so many massive stars ab a similar evolutionary epoch enables the generation. of the nested elant (20 300 pe diameter) shells that comprise the giant rregion (Aleaburn1980:Meaburn1990:Leitherer 1998)).," The combined winds, UV radiation and supernova explosions from so many massive stars at a similar evolutionary epoch enables the generation of the nested giant (20 – 300 pc diameter) shells that comprise the giant region \citealt{meaburn80,meaburn90,leitherer98}) )."
 On the largest scales. surrounding 30 Dor are supergiant. (600 1400 pe diameter) interstellar shells such as LMC3.," On the largest scales, surrounding 30 Dor are supergiant (600 – 1400 pc diameter) interstellar shells such as LMC3."
" The term ""shell will be used in this paper rather than the commonly. used term “bubble” since it is preferable to use a term that is civnamically neutral and constrained. to no specilie geometry. (c.g. spherical).", The term 'shell' will be used in this paper rather than the commonly used term `bubble' since it is preferable to use a term that is dynamically neutral and constrained to no specific geometry (e.g. spherical).
 Phe term ‘bubble’. often erroneously presupposes a roughly spherical. pressuredriven. energv-conserving shell.," The term `bubble', often erroneously presupposes a roughly spherical, pressure--driven, energy-conserving shell."
 Ες is certainlv not the case For the supergiant shells which are unlikely to be either spherical or energv-conserving., This is certainly not the case for the supergiant shells which are unlikely to be either spherical or energy-conserving.
" The division between ‘giant’ and ""supergiant when applied to the LMC shells will be for the diameter ranges above and recently confirmed. by the LE observations of Iximetal.(1999).", The division between 'giant' and 'supergiant' when applied to the LMC shells will be for the diameter ranges above and recently confirmed by the H observations of \citet{kim.et.al99}.
.. Dilferent. though related. mechanisms must be involved in the formation of LAIC shells in these distinctly. separate diameter ranges.," Different, though related, mechanisms must be involved in the formation of LMC shells in these distinctly separate diameter ranges."
 The most important dillerence is that supergiant shells have diameters in excess of the neutral gas scale-height. of the LMC., The most important difference is that supergiant shells have diameters in excess of the neutral gas scale-height of the LMC.
 The overlapping giant shells comprising the halo of 30 Doradus have been shown to be expanding at around. 50 ((c.@. Aleaburn1984:Chu&Ixennicutt. 1994)). whereas a multitude of 15pe diameter regions exhibit outllows of =200kms. .," The overlapping giant shells comprising the halo of 30 Doradus have been shown to be expanding at around 50 (e.g. \citealt{meaburn84,chu&kennicutt94}) ) whereas a multitude of $15~{\rm pc}$ diameter regions exhibit outflows of $\ga 200~{\rm
km~s^{-1}}$ ."
 Lhe latter were interpreted as voung supernova remnants in the perimeters of giant shells (Meaburn1988)., The latter were interpreted as young supernova remnants in the perimeters of giant shells \citep{meaburn88}.
. lis a cartoon that illustrates the hierarchy. of scale sizes present in a giant rreeion like 30 Doradus., 1 is a cartoon that illustrates the hierarchy of scale sizes present in a giant region like 30 Doradus.
 The brightest. dominant velocity. components of 30 Doradus are complex. but seem to be comprised. of three distinct velocity regimes corresponding to LL sheets along the sightline.," The brightest, dominant velocity components of 30 Doradus are complex but seem to be comprised of three distinct velocity regimes corresponding to H sheets along the sightline."
 These are at. 250kms 270kms! and 300kms (AleGeeotal.1978:Chu&Wennicutt1994:timetal. 1999)..," These are at $250~{\rm km~s^{-1}}$, $270~{\rm km~s^{-1}}$ and $300~{\rm km~s^{-1}}$ \citep{mcgee.et.al78,chu&kennicutt94,kim.et.al99}. ."
 In this paper. the systemic velocity. Voy. is taken to be the average heliocentric velocity (Viger) of these components. 270kms.l. in agreement with previous observations (see for example Peckctal.1997:1991:Garayetal.1993:Clayton 1987)).," In this paper, the systemic velocity, $V_{\rm sys}$ is taken to be the average heliocentric velocity $V_{\rm HEL}$ ) of these components, $270~{\rm km~s^{-1}}$, in agreement with previous observations (see for example \citealt{peck.et.al97short,meaburn91,garay.et.al93,clayton87}) )."
 In this work. the aim is to investigate the faint. highest speed. phenomena in the halo of 30 Doradus in order to completethe kinematical characterisation of this important eu rregion.," In this work, the aim is to investigate the faint highest speed phenomena in the halo of 30 Doradus in order to completethe kinematical characterisation of this important giant region."
 New echelle observations of the line profiles. of, New echelle observations of the line profiles of
In order {ο evaluate equation AS we assume that both the point sources and the noise have zero-mean Gaussian distribution functions with the variances 97 and 07 and are not correlated. spatially.,In order to evaluate equation \ref{eq5} we assume that both the point sources and the noise have zero-mean Gaussian distribution functions with the variances $\sigma_x^2$ and $\sigma_n^2$ and are not correlated spatially.
 We also assume that there is only one point source in the windowW., We also assume that there is only one point source in the window.
 Under these assumptions the integration in À4 can be performed to vield: Ilere After substituting equation AG in equation AS we obtain the linal expression lor the point source probability: The original downhill simplex algorithm described in Nelder&Mead:(1965) ancl (1971) minimizes a [unetion of No variables by using the values of the function at several verlices and (irving to move away from the highest vertex.," Under these assumptions the integration in \ref{eq4} can be performed to yield: Here After substituting equation \ref{eq6} in equation \ref{eq5} we obtain the final expression for the point source probability: The original downhill simplex algorithm described in \citet{Nelder}
 and \citet{O'Neill} minimizes a function of $N$ variables by using the values of the function at several vertices and trying to move away from the highest vertex."
 In our paper the function being minimized is the goodness-ol-lit measure |? delined in equation 4.., In our paper the function being minimized is the goodness-of-fit measure $\chi^2$ defined in equation \ref{chi2}.
 There are [our basic wavs lo move a vertex: reflection. expansion. contraction and shrinkage.," There are four basic ways to move a vertex: reflection, expansion, contraction and shrinkage."
 We adopted the simplex algorithin with a number of improvements., We adopted the simplex algorithm with a number of improvements.
 The changes areilustrated in Figure 13.., The changes areillustrated in Figure \ref{Simplex}.
 First. we modified reflection.," First, we modified reflection."
 If the change in \? is smaller than a user-specilied threshold. it is an indication that the reflection is done almost parallel to the iso-\7 lines.," If the change in $\chi^2$ is smaller than a user-specified threshold, it is an indication that the reflection is done almost parallel to the $\chi^2$ lines."
 In (his case an attempt is made to replace the reflection with a move in a perpendicular direction., In this case an attempt is made to replace the reflection with a move in a perpendicular direction.
 The number of perpendicular directions are equal to ΔΝ—1)., The number of perpendicular directions are equal to $2(N-1)$.
 The move is performed. if it results in a A? lower than the one achieved by the reflection.," The move is performed, if it results in a $\chi^2$ lower than the one achieved by the reflection."
 Another modification is that contraction and shrinkage have been replaced with the line minimization of 4? along the unsuccessful reflection direction., Another modification is that contraction and shrinkage have been replaced with the line minimization of $\chi^2$ along the unsuccessful reflection direction.
 Le. if the reflection results in point with higher than the original V7. a point with the lowest 4? is found on the line of the unsuccessful reflection.," I.e. if the reflection results in point with higher than the original $\chi^2$ , a point with the lowest $\chi^2$ is found on the line of the unsuccessful reflection."
outskirts.,outskirts.
 On the largest scale the pressure as well as the image appears to be quite smooth., On the largest scale the pressure as well as the image appears to be quite smooth.
 A weak lensing mass reconstruction of Clowe et al. (, A weak lensing mass reconstruction of Clowe et al. (
2003) shows that the cluster exhibits three dark matter peaks. with only the weakest of them corresponding to a X-ray peak. yet all of them are preceded by a shocked zone. seen in the entropy map.,"2003) shows that the cluster exhibits three dark matter peaks, with only the weakest of them corresponding to an X-ray peak, yet all of them are preceded by a shocked zone, seen in the entropy map."
 The main pressure peak is approximately located at the position of the center of the mass distribution from the weak lensing reconstruction., The main pressure peak is approximately located at the position of the center of the mass distribution from the weak lensing reconstruction.
 This center is adopted for the volume calculations and reported in Table .., This center is adopted for the volume calculations and reported in Table \ref{t:ol}.
 The entropy dip of the bullet is offset from the potential minimum and there 15 no entropy dip associated with the potential minimum of the main cluster., The entropy dip of the bullet is offset from the potential minimum and there is no entropy dip associated with the potential minimum of the main cluster.
 There are. however. entropy fluctuations in the pressure core. possibly associated with debris of the eitropy core of the main cluster.," There are, however, entropy fluctuations in the pressure core, possibly associated with debris of the entropy core of the main cluster."
 The spectroscopic analysis is reported in Table 13. and Fig.13.., The spectroscopic analysis is reported in Table \ref{t:cl14:t} and \ref{f:cl14}.
 It reveals temperature fluctuations by a factor of 1.5., It reveals temperature fluctuations by a factor of 1.5.
 The temperature of the bullet is only slightly lower than the bulk of the cluster., The temperature of the bullet is only slightly lower than the bulk of the cluster.
 However it exhibits a distinctly low entropy. which also allows us to trace the tail of the bullet.," However it exhibits a distinctly low entropy, which also allows us to trace the tail of the bullet."
 The zone assigned to bullet can be seen as negative deviation in entropy profile in Fig.13.., The zone assigned to bullet can be seen as negative deviation in entropy profile in \ref{f:cl14}.
 The bullet pressure peak is confirmed: it amounts to and is located behind the zone of lowest entropy in the bullet., The bullet pressure peak is confirmed; it amounts to and is located behind the zone of lowest entropy in the bullet.
 By combining together all the high-entropy zones. associated with the shock heating. we have achieved significance in the temperature variation. from 10.," By combining together all the high-entropy zones, associated with the shock heating, we have achieved significance in the temperature variation, from 10."
 to 1422 keV. This corresponds to à Mach number of 1.4+0.2., to $14\pm2$ keV. This corresponds to a Mach number of $1.4\pm0.2$.
 This estimate is lower. compared to the shock parameters deduced from the image showing the Mach cone.," This estimate is lower, compared to the shock parameters deduced from the image showing the Mach cone."
 A higher Mach number would be obtained from the entropy enhancement: 2.6+0.2., A higher Mach number would be obtained from the entropy enhancement: $2.6\pm0.2$.
 It is plausible that the extraction region captures both shock and postshock gas., It is plausible that the extraction region captures both shock and postshock gas.
 The later has lower pressure. but records its state in the entropy.," The later has lower pressure, but records its state in the entropy."
 As was noted above. the observed shock is located in front of the outward moving dark matter potential.," As was noted above, the observed shock is located in front of the outward moving dark matter potential."
 Since the potentials carry no longer any gas. they do not cause this shock. but just travel at the same speed.," Since the potentials carry no longer any gas, they do not cause this shock, but just travel at the same speed."
 This implies that we observe the initial forward shock propagating through the cluster., This implies that we observe the initial forward shock propagating through the cluster.
 The entropy ratio also shows that the eastern part of the cluster has lower entropy. as due to the stripping of the bulk of the bullet cluster.," The entropy ratio also shows that the eastern part of the cluster has lower entropy, as due to the stripping of the bulk of the bullet cluster."
 An analysis of the two-dimensional structure 1n the REFLEX clusters. as seen in the images and spectral hardness ratio maps. reveals statistically significant substructure. probably originating from different stages of cluster merger.," An analysis of the two-dimensional structure in the REFLEX clusters, as seen in the images and spectral hardness ratio maps, reveals statistically significant substructure, probably originating from different stages of cluster merger."
 We are able to see the substructure even at very late merger stages. where for example the X-ray image appears to be quite symmetric.," We are able to see the substructure even at very late merger stages, where for example the X-ray image appears to be quite symmetric."
 We identify the entropy to be most sensitive to both late stage mergers with the associated slow buoyancy action of relaxation of the cluster and to strong shocks. which change the entropy.," We identify the entropy to be most sensitive to both late stage mergers with the associated slow buoyancy action of relaxation of the cluster and to strong shocks, which change the entropy."
 Two mergers with large Mach numbers are found., Two mergers with large Mach numbers are found.
 A statistical analysis of the substructure in the pressure and entropy maps. reveals significant fluctuations around the mean profile.," A statistical analysis of the substructure in the pressure and entropy maps, reveals significant fluctuations around the mean profile."
 Typically. pressure fluctuations are found on the level. while the entropy fluctuations are at the level.," Typically, pressure fluctuations are found on the level, while the entropy fluctuations are at the level."
 Apparently. smoother appearance of the pressure maps should be attribtted to the larger dynamical range of the map. covering typically two orders of magnitude.," Apparently, smoother appearance of the pressure maps should be attributed to the larger dynamical range of the map, covering typically two orders of magnitude."
 A comparison of our sample with a similar analysis of hydro-dynamical simulations. by Finoguenov et al. (, A comparison of our sample with a similar analysis of hydro-dynamical simulations by Finoguenov et al. (
2005) reveals a similar distribution of clusters vs the level of the substructure in both entropy and pressure.,2005) reveals a similar distribution of clusters vs the level of the substructure in both entropy and pressure.
 Anunber of clusters exhibit a presence of low entropy gas in the outskirts. deviating by at least an order of magnitude from the prescription of gravitational heating.," A number of clusters exhibit a presence of low entropy gas in the outskirts, deviating by at least an order of magnitude from the prescription of gravitational heating."
 Surprisingly enough. these regions have gas pressures similar to that of," Surprisingly enough, these regions have gas pressures similar to that of"
of perturbation wave-vector perpendicular to the magnetic Ποια.,of perturbation wave-vector perpendicular to the magnetic field.
" Simplilving the linearized equations (8))-(13)) by repeated use of (he unperturbed background equations (15))-(13)). we obtain OD, The perturbed αν velocity due to the perturbed magnetic lied changes the ambipolar diffusion heating rate."," Simplifying the linearized equations \ref{drift}) \ref{ebstate}) ) by repeated use of the unperturbed background equations \ref{backgden}) \ref{st}) ), we obtain where The perturbed drift velocity due to the perturbed magnetic field changes the ambipolar diffusion heating rate."
 In fact. the net cooling Function € in equation (21)) must contain also (he ambipolear heating due to perturbed magnetic field.," In fact, the net cooling function $\Omega$ in equation \ref{lineng}) ) must contain also the ambipolar heating due to perturbed magnetic field."
 Iowever. because we have expressed (he ambipolar heating in equation (6)) bv parameters. in numerical calculations for drawing ligures. we can involve this effect by the amount of parameters.," However, because we have expressed the ambipolar heating in equation \ref{heatad}) ) by parameters, in numerical calculations for drawing figures, we can involve this effect by the amount of parameters."
 Thus. for simplicitv. we neglect the explicit calculations of the changes of ambipolar heating due to perturbed crilt velociv.," Thus, for simplicity, we neglect the explicit calculations of the changes of ambipolar heating due to perturbed drift velocity."
 since (he equilibrium is Gime dependent. (he normal modes of the svstem are time dependent too. (hus. we must apply some approximation techniques namely WID approximation (e.g.. Dora aud Daruah 2008) to gain valuable insight into the nature of problem: this analysis mav be a challenging task in a subsequent research.," Since the equilibrium is time dependent, the normal modes of the system are time dependent too, thus, we must apply some approximation techniques namely WKB approximation (e.g., Bora and Baruah 2008) to gain valuable insight into the nature of problem; this analysis may be a challenging task in a subsequent research."
 Here. we consider the isobaric TI. which is more realistic phenomena in (he interstellar gases.," Here, we consider the isobaric TI, which is more realistic phenomena in the interstellar gases."
 Gathering the equations (23)) and (21)) with equation (19)). in isobaric case (py= 0). leads to an exponential growth as follows," Gathering the equations \ref{linsta}) ) and \ref{lineng}) ) with equation \ref{linden}) ), in isobaric case $p_1=0$ ), leads to an exponential growth as follows"
After optimizing the Πί for the autocorrelation measurements. we then looked at how conditioning the cross-ccorrelation covariance matrices affects the overall reconstruction ofο,"After optimizing the fits for the autocorrelation measurements, we then looked at how conditioning the correlation covariance matrices affects the overall reconstruction of."
"ρ), Since the uncertainty in iis dominated by the uncertainty inz).. this conditioning should have the greatest impact onw,,(7. the reconstruction."," Since the uncertainty in is dominated by the uncertainty in, this conditioning should have the greatest impact on the reconstruction."
" We generate 10 pick-4 measurements by averaging the correlation measurements from four randomly selected fields out of the 24. which we then use to calculate©,(<)."," We generate $10^4$ pick-4 measurements by averaging the correlation measurements from four randomly selected fields out of the 24, which we then use to calculate."
. For calculating the risk. we know the true redshift distribution in each field perfectly from the simulation. so we do not need to rely on synthetic techniques as in refsec:optwpp..," For calculating the risk, we know the true redshift distribution in each field perfectly from the simulation, so we do not need to rely on synthetic techniques as in \\ref{sec:optwpp}."
" Since the fits for both aand wwere wt,best) with a few percent ridge regression conditioning refsec:optwpp.. refsec:optwprp)). for simplicity we adopt f=3.5% üs the optimal conditioning in both cases."," Since the fits for both and were best with a few percent ridge regression conditioning \\ref{sec:optwpp}, \\ref{sec:optwprp}) ), for simplicity we adopt $f$ as the optimal conditioning in both cases."
" For each pick-+ measurement. we determine the autocorrelation parameters of the photometric sample by fitting the ffrom the selected 1,,,,(7)4 fields using the optimally conditioned covariance matrix calculated from the 24 fields."," For each pick-4 measurement, we determine the autocorrelation parameters of the photometric sample by fitting the from the selected 4 fields using the optimally conditioned covariance matrix calculated from the 24 fields."
" All three parameters(A,).. and are left free and fit simultaneously."," All three parameters, and ) are left free and fit simultaneously."
" To measure555. the C,,)evolution of the correlation function parameters of the spectroscopic sample. we calculated un 10 z-bins 0)covering the range 0.11selected-1.4. where the size and location of each z-bin was such that there were approximately the same number of objects in each one."," To measure the evolution of the correlation function parameters of the spectroscopic sample, we calculated in 10 $z$ -bins covering the range $0.11<z<1.4$, where the size and location of each z-bin was selected such that there were approximately the same number of objects in each one."
 In each z-bin we calculate the covariance matrix from the 24 fields and fit each pick-4 measurement using the optimal conditioning to determine ro4) 5ol5).," In each $z$ -bin we calculate the covariance matrix from the 24 fields and fit each pick-4 measurement using the optimal conditioning to determine $r_{0,ss}(z)$ and $\gamma_{ss}(z)$."
" In one redshift bin (Q.11<andz« the values of Foxy and >,. Obtained with these methods 0.268).were significantly different from the values determined when assuming no covariance."," In one redshift bin $0.11<z<0.268$ ), the values of $r_{0,ss}$ and $\gamma_{ss}$ obtained with these methods were significantly different from the values determined when assuming no covariance."
 We investigated the likelihood contours in detail and found they were not well behaved: not only were the median parameter values different from. the result with no covariance. the standard deviation of the 10 pick-4 measurements proved to be an underestimate of the uncertainty in. that bin. which had significant effects when performing an error-weighted linear fit to sGoandAZ) .," We investigated the likelihood contours in detail and found they were not well behaved; not only were the median parameter values different from the result with no covariance, the standard deviation of the $10^4$ pick-4 measurements proved to be an underestimate of the uncertainty in that bin, which had significant effects when performing an error-weighted linear fit to $(z)$ and $(z)$."
"Weattemptedavarietyo fme thodsforestimatingthetngs""theJ", We attempted a variety of methods for estimating the errors in that bin with poor results.
 MNSand Sotand SHORSNto (hesethe aoserthesh« rather than 0.1<ry10f !Mpc. gave more well behavedMpe. values (more consistent with the values in other redshift bins or those obtained when ignoring covariance) and improved the reconstruction.," However, we found that fitting over the shorter range $0.25<r_p<10\ h^{-1}$ Mpc, rather than $0.1<r_p<10\ h^{-1}$ Mpc, gave more well behaved values (more consistent with the values in other redshift bins or those obtained when ignoring covariance) and improved the reconstruction."
 For consistency we fit over this range for all bins where z«0.8., For consistency we fit over this range for all bins where $z<0.8$.
" As in MNIO we continue to fit over the range 1.0<r,«10h Mpe for z>0.8. as in the Millennium simulations (though less so in real datasets) ddivergesWU) significantly from a power law at 0.1«ry<A7! Mpc."," As in MN10 we continue to fit over the range $1.0<r_p<10\ h^{-1}$ Mpc for $z>0.8$, as in the Millennium simulations (though less so in real datasets) diverges significantly from a power law at $0.1<r_p<1\ h^{-1}$ Mpc."
 While the conditioning of the fits for the autocorrelation parameters was kept the same for each measurement. we varied the conditioning of the cross-ccorrelation fits to see how it affects the reconstruction.," While the conditioning of the fits for the autocorrelation parameters was kept the same for each measurement, we varied the conditioning of the correlation fits to see how it affects the reconstruction."
 We bin the spectroscopic sample over the range 0.19<z«1.39 with a bin size of Az=0.04 and measure ον(0) in each bin., We bin the spectroscopic sample over the range $0.19<z<1.39$ with a bin size of $\Delta z=0.04$ and measure $w_{sp}(\theta)$ in each bin.
 At each level of conditioning we: In step 5. we calculate the risk over a slightly limited redshift range to eliminate bins where noise dominates the measurements. which diluted our ability to assess the impact of ridge regression.," At each level of conditioning we: In step 5, we calculate the risk over a slightly limited redshift range to eliminate bins where noise dominates the measurements, which diluted our ability to assess the impact of ridge regression."
 Figure 5. shows both mean risks as a function of the conditioning. compared to the risk using methods identical to MNIO.," Figure \ref{fig:phirisk} shows both mean risks as a function of the conditioning, compared to the risk using methods identical to MN10."
 We optimized for the mean risk over the redshift range rather than the maximum risk as the latter was dominated by random outliers (due to the smaller number of objects in the redshift bins used. errors in aare much larger. and hencew40. random excursions extend further. than for the autocorrelations).," We optimized for the mean risk over the redshift range rather than the maximum risk as the latter was dominated by random outliers (due to the smaller number of objects in the redshift bins used, errors in are much larger, and hence random excursions extend further, than for the autocorrelations)."
 Both techniques indicate that the minimum risk is obtained at around a few percent conditioning., Both techniques indicate that the minimum risk is obtained at around a few percent conditioning.
 There is a substantial improvement in both measures. but particularly in the risk comparing the redshift distribution for the four chosen fields to the overall (e.g. universal) mean.," There is a substantial improvement in both measures, but particularly in the risk comparing the redshift distribution for the four chosen fields to the overall (e.g. universal) mean."
 Figure 6 shows the reconstruction for Sn if i. oHDy) Bl as variance bias. compares reconstruction using methods identical to MNIO.," Figure \ref{fig:phiplot} shows the reconstruction for conditioning (i.e. the same for all three fits) as well as the variance and bias, and compares to the reconstruction using methods identical to MN10."
 The decrease in the variance is significant in each redshift bin while the bias is relatively unchanged in all but a few z-bins., The decrease in the variance is significant in each redshift bin while the bias is relatively unchanged in all but a few $z$ -bins.
 By incorporating full covariance information and ridge regression methods. the square root of the fractional risk is <40% smaller than that resulting from our prior methods.," By incorporating full covariance information and ridge regression methods, the square root of the fractional risk is $<40\%$ smaller than that resulting from our prior methods."
 In this paper we have improved on the cross-ccorrelation techniques presented. in. ? by incorporating full covariance information., In this paper we have improved on the correlation techniques presented in \citet{2010ApJ...721..456M} by incorporating full covariance information.
 In addition. we have demonstrated the improvements that result from incorporating. ridge regression in. fiting for correlation. function. parameters.," In addition, we have demonstrated the improvements that result from incorporating ridge regression in fitting for correlation function parameters."
 Conditioning using ridge regression allowed us to obtain a more stable inversion of the covariance matrix by reducing the impact of noise in the off diagonal elements. resulting in better estimates of the correlation function parameter values:," Conditioning using ridge regression allowed us to obtain a more stable inversion of the covariance matrix by reducing the impact of noise in the off diagonal elements, resulting in better estimates of the correlation function parameter values;"
The discovery of bright N-ray sources in external galaxies has generated significant interest. due to the possibility that. they may be intermediate-mass black holes suggested by the masses required (ο not violate the Edclington limit ∙∙≯for the inferredE high16 Iumimosities1Ales (ColbertByor&-Mushotzky:yleyal. 2001).,"The discovery of bright X-ray sources in external galaxies has generated significant interest due to the possibility that they may be `intermediate-mass' black holes suggested by the masses required to not violate the Eddington limit for the inferred high luminosities \citep{colbert99,makishima00,kaaret01}."
 However. there are alternative interpretations of these so-called uliraluminous X-ray sources (ULXs) in which the X-ravs are beamecl or the Eddington limit is violated and the existence of intermecdiate-mass black holes are not required.," However, there are alternative interpretations of these so-called `ultraluminous X-ray sources' (ULXs) in which the X-rays are beamed or the Eddington limit is violated and the existence of intermediate-mass black holes are not required."
 The huninosity [unction of X-ray. point sources in external galaxies does not appear to have a break around the Eddington limit [or. ‘normal’. stellar mass black holes. i.e.. with M<20M.. suggesting that objects above and below this limit and members," The luminosity function of X-ray point sources in external galaxies does not appear to have a break around the Eddington limit for `normal' stellar mass black holes, i.e. with $M < 20 M_{\odot}$, suggesting that objects above and below this limit and members"
for the first generation of stars. showed that. if eiubedded in halos with hieh DM densities. stars can be fuelled oulv by the enerey frou: DM annihilation.,"for the first generation of stars, showed that, if embedded in halos with high DM densities, stars can be fuelled only by the energy from DM annihilation."
 Mauy authors confirmed aud developed these results., Many authors confirmed and developed these results.
 DAL capture aud annihilation was studied for white dwarfs and neutron stars bv ? aud 7S, DM capture and annihilation was studied for white dwarfs and neutron stars by \citet{art-BertoneFairbairn2008} and \citet{art-Kouvaris2008}.
" The same was done for first stars by 7. 7.. ?.. and. ?.. aud their imuplicatious for cosmic relonization aud the pair production supernova rate were studied bv ? and ὧν, respectively,"," The same was done for first stars by \citet{art-Iocco2008}, \citet{art-Taosoetal2008}, \citet{art-FreeseSpolyarAguirre2008}, and \citet{art-YoonIoccoAkiyama2008}, and their implications for cosmic reionization and the pair production supernova rate were studied by \citet{art-SchleicherBK2008} and \citet{art-Iocco2009}, respectively."
 Tn the case of first stars formed at the ceuter of DM. iminuilialos. DM αλαπο may become the primary source of energy during the adiabatic coutractiou regime of the DM halo. counteracting the gravitational collapse even before DM capture becomes efficient (22777).," In the case of first stars formed at the center of DM minihalos, DM annihilation may become the primary source of energy during the adiabatic contraction regime of the DM halo, counteracting the gravitational collapse even before DM capture becomes efficient \citep{let-Spolyaretal2008, art-FreeseGS2008, art-FreeseBSG08, art-Natarajan2009, art-RipamontiIoccoetal2009}."
 When the star ruus out of its original DM. which takes about a 1ailliou years (?).. this stalling phase euds aud the collapse coutiuues util capture replenishes the fuel for DAL aunihilation (?)..," When the star runs out of its original DM, which takes about a million years \citep{art-FreeseReview2009}, this stalling phase ends and the collapse continues until capture replenishes the fuel for DM annihilation \citep{art-Ioccoetal2008}."
 Regarding low-nass stars. 7. carried out an extensive y.udy on how the stars located at the ealactic ceuter may )o affected by the dense DM densities which are expected o exist there.," Regarding low-mass stars, \citet{art-Scottetal2009} carried out an extensive study on how the stars located at the galactic center may be affected by the dense DM densities which are expected to exist there."
 They assmued that these stars were dready formed in a scenario without DAL, They assumed that these stars were already formed in a scenario without DM.
 Therefore. icr evolution on a DM halo was considered from the ZAMS.," Therefore, their evolution on a DM halo was considered from the ZAMS."
 In addition to that. we also took iuto account re capture of DAL particles from the collapse of the xotostar. as ?. did for first stars.," In addition to that, we also took into account the capture of DM particles from the collapse of the protostar, as \citet{art-Ioccoetal2008} did for first stars."
 The influence of DM oei that early stage max have dramatic consequences in ie forming star. as will be shown.," The influence of DM in that early stage may have dramatic consequences in the forming star, as will be shown."
 The paper is orgauised as follows: iu section 2. there is a short description of the stellar evolution code used here., The paper is organised as follows: in section \ref{sec-thecode} there is a short description of the stellar evolution code used here.
 The process of accretion of annihilating DAL particles In stars is briefly reviewed in section 3.., The process of accretion of annihilating DM particles in stars is briefly reviewed in section \ref{sec-ADM}.
 The differeut evolution paths found for low-mass stars evolving witlin DAL halos are discussed in section L., The different evolution paths found for low-mass stars evolving within DM halos are discussed in section \ref{sec-Evol}.
 Iu section 5 we discuss the implication of our results to observational stellar astroplivsics., In section \ref{sec-Stellar} we discuss the implication of our results to observational stellar astrophysics.
 Finally. iu section 6.. we preseut a byief πα of the results aud discuss the implication of such results for DM research.," Finally, in section \ref{sec-conclusion}, we present a brief summary of the results and discuss the implication of such results for DM research."
 The models computed in this work were made using the stellar evolution code CESAAL (?).., The models computed in this work were made using the stellar evolution code CESAM \citep{art-Morel1997}.
" The CESAM code is a cousistent set of programs and routines which performs calculations of 1D quasi-static stellar evolution inchiding diffusion. rotation and mass loss,"," The CESAM code is a consistent set of programs and routines which performs calculations of 1D quasi-static stellar evolution including diffusion, rotation and mass loss."
 CESAM conrputes structure equilibriuni models bv collocation method based on piecewise polvuonmials approximations projected on their B-spline basis. the evolution of the chemical composition is solved by stifüv stable schemes of orders up to four. the solution of the diffusion equation cluplovs the Petrov-Galerkin scheme (7).," CESAM computes structure equilibrium models by collocation method based on piecewise polynomials approximations projected on their B-spline basis, the evolution of the chemical composition is solved by stiffly stable schemes of orders up to four, the solution of the diffusion equation employs the Petrov-Galerkin scheme \citep{art-Morel1997}."
. The code deteriuues the evolution of a star by the inteeration of the set of conservation structure equations. coupled with a set of nuclear reactions describing the nucleosvuthliesis of chemical elements aud the production of energy. using an adaptive space-time mesh. Le. the code chooses an optimal step im time and space by computing the rate of variation of the equilibrium quautitics.," The code determines the evolution of a star by the integration of the set of conservation structure equations, coupled with a set of nuclear reactions describing the nucleosynthesis of chemical elements and the production of energy, using an adaptive space-time mesh, i.e., the code chooses an optimal step in time and space by computing the rate of variation of the equilibrium quantities."
 CESAM allows calculations of stellar models witlh various physical assumptions. plivsical data. external boundary conditions. nuinercal methods aud iuuierical accuracy: in this work the accuracy is set to 107.," CESAM allows calculations of stellar models with various physical assumptions, physical data, external boundary conditions, numerical methods and numerical accuracy; in this work the accuracy is set to $10^{-5}$."
 This code has the ability to compute the evolution of stars from the pre-main sequence up until the begiunius of the Ie burning cycle for various stellar masses. and for a range of metallicities (0.0001 «0.01. in the preseut configuration).," This code has the ability to compute the evolution of stars from the pre-main sequence up until the beginning of the $^4$ He burning cycle for various stellar masses, and for a range of metallicities $<\;$ $\;<$ 0.04, in the present configuration)."
 The initial chemical abundanuces are the solar ones (2). and the initial metallicity. uuless stated otherwise. is assunied to be Z=0.011. simular to that of the Sun.," The initial chemical abundances are the solar ones \citep{art-AsplundGrevesseSauval2005} and the initial metallicity, unless stated otherwise, is assumed to be Z=0.014, similar to that of the Sun."
 The mücroscopic phvsies. includiug equation of state. opacities. nucleosvuthesis. nücroscopic diffusion aud chenucal abundances. is verv refined in this stellar evolution code. since this part of the code was tested in the case of the Sun against helioseisnüc data.," The microscopic physics, including equation of state, opacities, nucleosynthesis, microscopic diffusion and chemical abundances, is very refined in this stellar evolution code, since this part of the code was tested in the case of the Sun against helioseismic data."
" The CESAM evolution code is well established iu the solar and stellar plysics comunity. being used cither to computeo solar models (?7) or models of other stars (amone others: PMSóScuti star V316 Ori (?).. JC hopei Star p LDrideni (?7).. Q.s-s\0.. PAIS aud. MS. stars ο), Veea-like stars aPsA. νου, Pic. zEri aud rCet (?).. a CM star Procvon A (2)... solar-like INE. stars aud PMS stars (2).. aCen A&DD stars (2). and MSstars with couvective overshooting (?2)))."," The CESAM evolution code is well established in the solar and stellar physics comunity, being used either to compute solar models \citep{art-CouvidatTuChi2003, art-Berthomieu1993} or models of other stars (among others: $\delta$ Scuti star V346 Ori \citep{art-Bernabeietal2009}, $\beta$ Chepei star $\nu$ Eridani \citep {art-Suarezetal2009}, $_{\odot}$ PMS and MS stars \citep{art-Marques2008}, Vega-like stars $\alpha$ PsA, $\beta$ Leo, $\beta$ Pic, $\varepsilon$ Eri and $\tau$ Cet \citep{art-DiFolco2004}, $\alpha$ CMi star Procyon A \citep{art-Kervella2004}, solar-like $_{\odot}$ stars and PMS stars \citep{art-PiauTurkChieze2002}, $\alpha$ Cen B stars \citep{art-Thevenin2002} and MSstars with convective overshooting \citep{art-Audardetal1995}) )."
 We modified this code to take into account the impact of DAL particles iu the evolution of low-mass stars., We modified this code to take into account the impact of DM particles in the evolution of low-mass stars.
 Iu particular. we have already used this code to study the evolution of the Sun within a halo of dark matter (7773: we predicted the results of two eroups of observables: solar neutrinos and helioseisiunologw data Guchicdine the sound speed profile).," In particular, we have already used this code to study the evolution of the Sun within a halo of dark matter \citep{art-LopesBS2002,art-LopesSH2002,let-LopesSilk2002}; we predicted the results of two groups of observables: solar neutrinos and helioseismology data (including the sound speed profile)."
 WIAIPS travel through stars. where they expericuce scattering off the uuclei that they encouuter in the stellar cores.," WIMPs travel through stars, where they experience scattering off the nuclei that they encounter in the stellar cores."
 Although most WIMPs travel right through the star without sufferiug any type of interaction. some of them will scatter off nuclei losing cucrey.," Although most WIMPs travel right through the star without suffering any type of interaction, some of them will scatter off nuclei losing energy."
 If they lose enough enerev. they wont be able to escape the eravitational field of the star amvinore.," If they lose enough energy, they won't be able to escape the gravitational field of the star anymore."
 The umber of WIMPs trapped inside the star is measured by the capture rate C4., The number of WIMPs trapped inside the star is measured by the capture rate $C_\chi$.
" The captive rate is proportional to the WIMD scattering cross section off unclei o, aud the dark inatter density of the halo py.", The capture rate is proportional to the WIMP scattering cross section off nuclei $\sigma_\chi$ and the dark matter density of the halo $\rho_\chi$.
" It is inversely proportional to the WIMP inass sj, and to the WIMP dispersion velocity oy.", It is inversely proportional to the WIMP mass $m_\chi$ and to the WIMP dispersion velocity $\bar{v}_\chi$.
 Capture rates of WIAIPs were first calculated bv? in the case of the Sum. bv ?. for ecueric massive bodies aud for the earth im particular. aud bv ?/— for Iain sequence stars;," Capture rates of WIMPs were first calculated by \citet{art-PressSpergel1985} in the case of the Sun, by \citet{art-Gould1987} for generic massive bodies and for the earth in particular, and by \citet{art-BouquetSalati1989} for main sequence stars."
" We computed the capture rate using equation 2.31 of ?:: where esuV26M, Rods the escape volocitv. M, and AR, ave the mass aud radius of the star. ss, 15d the proton mass. aude; and A; are. respectively. the lass fraction and the number of uncleous of clement ἐν"," We computed the capture rate using equation 2.31 of \citet{art-Gould1987}: where $v_{esc}=\sqrt{2GM_{\star}/R_{\star}}$ is the escape velocity, $M_{\star}$ and $R_{\star}$ are the mass and radius of the star, $m_p$ is the proton mass, and $x_i$ and $A_{i}$ are, respectively, the mass fraction and the number of nucleons of element $i$."
" We included the coutributious from 12 nuclei IL. Ne. 12€, HN, 150, 211 7e. ""Lii Πο, BOL PN, PO aud ""De. the abunudauces of which are followed by our code."," We included the contributions from 13 nuclei: H, $^4$ He, $^{12}$ C, $^{14}$ N, $^{16}$ O, $^{2}$ H, $^{3}$ He, $^{7}$ Li, $^{7}$ Be, $^{13}$ C, $^{15}$ N, $^{17}$ O and $^{9}$ Be, the abundances of which are followed by our code."
 For all of them. we computed their spiu- (SI) interaction with WIMPs (colerent scattering). which scales as the fourth power of the nucleus mass umuber. leading to a scattering cross section a;=Afoy sy.," For all of them, we computed their spin-independent (SI) interaction with WIMPs (coherent scattering), which scales as the fourth power of the nucleus mass number, leading to a scattering cross section $\sigma_{i}=A_i^4\sigma_{\chi,SI}$ ."
 For hydrogen. we also took iuto," For hydrogen, we also took into"
values indicates that either the SED fit is bad (due to incorrect line classification). or the line feature seen inthe FP scan is awlaard.,"values indicates that either the SED fit is bad (due to incorrect line classification), or the line feature seen in the FP scan is awkward."
 This information is used to produce another 7 contributor. calculated from the discrepancy aud the intrinsic errors of the two flux levels.," This information is used to produce another $\chi^2$ contributor, calculated from the discrepancy and the intrinsic errors of the two flux levels."
 Iuchuliug the coutiumun level derived above. a new revised line fit is performed iu the same way as m step 1.," Including the continuum level derived above, a new revised line fit is performed in the same way as in step 1."
 Tu this second iteration. however. bad line classifications sxecomme Visible by large offsets between the coutimmun evel eiven bv the SED fit and the corrected pre-filter baseline. leading to a lower S/N ratio for the corresponding Lue fit.," In this second iteration, however, bad line classifications become visible by large offsets between the continuum level given by the SED fit and the corrected pre-filter baseline, leading to a lower S/N ratio for the corresponding line fit."
 This also happens when the line eature in the FP data is in reality a fake detection due to a statistical fuctuation of the FP signals., This also happens when the line feature in the FP data is in reality a fake detection due to a statistical fluctuation of the FP signals.
 The new S/N value is later used as critical indicator or the reality of the Lue feature seen in the FP sean and whether the object will be accepted as an enmission liue ealaxy., The new S/N value is later used as critical indicator for the reality of the line feature seen in the FP scan and whether the object will be accepted as an emission line galaxy.
 The coutimmun fit found this way serves as a baseliue to determine the fluxes of the other (secondary) cussion lines., The continuum fit found this way serves as a baseline to determine the fluxes of the other (secondary) emission lines.
 For cach of the eight possible line identifications. the fluxes for the secondary lines were derived from the difference between the flux. measured in the cveto filters’ (the filters where other promunent lines would be expected) aud the continua level predicted from the template fit.," For each of the eight possible line identifications, the fluxes for the secondary lines were derived from the difference between the flux measured in the `veto filters' (the filters where other prominent lines would be expected) and the continuum level predicted from the template fit."
 For the calculation of the line streugth. the excess flux was normalized by the trausimission of the filter at the wavelength eiven by the redshift derived from the FP observation.," For the calculation of the line strength, the excess flux was normalized by the transmission of the filter at the wavelength given by the redshift derived from the FP observation."
 In cases where the veto filter tuchided more thaw oue euissionu line which generally Lappeus for the 0501190 doublet. the line flux ratio was fixed to the canonical values.," In cases where the veto filter included more than one emission line which generally happens for the 501/496 doublet, the line flux ratio was fixed to the canonical values."
 11 shows such fits for several objects iu tle left pancls., 1 shows such fits for several objects in the left panels.
 Tn all spectra. except for the ΟΠ galaxy. a uuiuber of ‘secoudary dues can be seen.," In all spectra, except for the ] galaxy, a number of `secondary' lines can be seen."
 Thc thus determined line fiux ratios are compared witli a catalog of observed line ratios for about 500 nearby ealaxics based ou data frou the literature. including a yange of Sevfert galaxies to compact dwarfs. most notably from French (1980). MeCall et al. (," The thus determined line flux ratios are compared with a catalog of observed line ratios for about 500 nearby galaxies based on data from the literature, including a range of Seyfert galaxies to compact dwarfs, most notably from French (1980), McCall et al. ("
L985). Popeseu Hopp (2000). Veilleux Osterbrock (1987). Vogel et al. (,"1985), Popescu Hopp (2000), Veilleux Osterbrock (1987), Vogel et al. ("
1993).,1993).
 À X7D ds: estimatedH inH the following+ wav: Iu the uultidinensiounal space iade up by all Lue ratio combination possible for the respective classification of the line observed iu the EPI window. the mean distance of the considered galaxy to the three closest tabulated ratios was normalized by the error ellipse of the observed ratios.," A $\chi^2$ is estimated in the following way: In the multidimensional space made up by all line ratio combination possible for the respective classification of the line observed in the FPI window, the mean distance of the considered galaxy to the three closest tabulated ratios was normalized by the error ellipse of the observed ratios."
 This method prevents a bias towards the most colon line ratios. but allows also galaxies with extreme line," This method prevents a bias towards the most common line ratios, but allows also galaxies with extreme line"
llere we derive an expression for a non-linear filler based on the notion of point source probability.,Here we derive an expression for a non-linear filter based on the notion of point source probability.
 We assume that image s observed by (he detector consists of (he point source contribution x convolved with the PRE and additive noise n: Ilere His a (ranslationally invariant matrix constructed from the PRE: j).," We assume that image ${\mathbf
s}$ observed by the detector consists of the point source contribution ${\mathbf x}$ convolved with the PRF and additive noise ${\mathbf n}$: Here ${\mathbf H}$ is a translationally invariant matrix constructed from the PRF: $ {\mathbf H_{ij}} = {\mathbf H_{i-j}} = PRF(i-j)$ ."
 The general problem is to estimate the probability of (he point source being al a particular pixel? given a measurement vector s in a certain windowW surrounding this pixel., The general problem is to estimate the probability of the point source being at a particular pixel given a measurement vector ${\mathbf s}$ in a certain window surrounding this pixel.
 The size of the windowΗ 15 determined by the size of the PRE., The size of the window is determined by the size of the PRF.
" We assume that the point source and background noise are characterized bv the distribution functions f.(s) aud f,(S). corvespondingly."," We assume that the point source and background noise are characterized by the distribution functions $f_x({\mathbf s})$ and $f_n({\mathbf s})$, correspondingly."
 We consider (wo hypotheses for the pixel#, We consider two hypotheses for the pixel.
 The first hypothesis hy is (hat there is a point source al the pixel. and the second à (null hypothesis) is that there is not a point source at the pixel.," The first hypothesis $h_1$ is that there is a point source at the pixel, and the second $h_2$ (null hypothesis) is that there is not a point source at the pixel."
" The probability of the A!” hypothesis conditioned on the measurement s is givenby the Bavesian theorem: where P(h,) is the a priori probability of the A hypothesis. f(s) is the probability density of observing the set of pixel values s."," The probability of the $k^{th}$ hypothesis conditioned on the measurement ${\mathbf s}$ is givenby the Bayesian theorem: where $P(h_k)$ is the a priori probability of the $k^{th}$ hypothesis, $f({\mathbf s})$ is the probability density of observing the set of pixel values ${\mathbf s}$ ."
" Assuming; completeness of the hypothesis set. i.e. P(h4)+Pe)=1. we obtain lor f(s) The probability density of measurement s under the null hypothesis is given simply bv (he noise distribution function. f,(s)."," Assuming completeness of the hypothesis set, i.e. $P(h_1) + P(h_2) = 1$, we obtain for $f({\mathbf s})$ The probability density of measurement ${\mathbf s}$ under the null hypothesis is given simply by the noise distribution function $f_n({\mathbf s})$."
 The probability. density of measurement s under (he point source hypothesis is (he result of integration over all possible point source contributions X: Combining evervthing we obtainfor the quantity in question:," The probability density of measurement ${\mathbf
s}$ under the point source hypothesis is the result of integration over all possible point source contributions ${\mathbf x}$ : Combining everything we obtainfor the quantity in question:"
,————.
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Norman (2000). in the simplest version of the relativistic precession model the data ou the lower frequency horizoutal ranch type buieltucss oscillatious combined with he kilohe:rtz oscillations requires that he stelar spin frequency exceed 2500. ITz in all sources. and be especially arge {up to  900 ITz) iu the Z sources.,"Norman (2000), in the simplest version of the relativistic precession model the data on the lower frequency horizontal branch type brightness oscillations combined with the kilohertz oscillations requires that the stellar spin frequency exceed $\sim$ 500 Hz in all sources, and be especially large (up to $\sim$ 900 Hz) in the Z sources."
 Simple estimates of corrections due to fluid viscosity (Psaltis Norman 2000) sugeest that hese corrections would nerease fi1’ther he reqired spin frequency., Simple estimates of corrections due to fluid viscosity (Psaltis Norman 2000) suggest that these corrections would increase further the required spin frequency.
 The predictions Or Moin therefore differ ercatly in the wo picures. and if a clear spin frequency is detected in the persistent. acciDon-powered cluission then this will have strong carne onu t16 modeling.," The predictions for $\nu_{\rm spin}$ therefore differ greatly in the two pictures, and if a clear spin frequency is detected in the persistent, accretion-powered emission then this will have strong bearing on the modeling."
 In addilon. 1 ait lapous that the spin frequeney is not equal to the]must oscillation frequency or half of it. this wil overturn the apparently conipelliug areunmieuts hat burst brightuess oscillations cui only be explained via a rotational modilatio1 mechanisin.," In addition, if it happens that the spin frequency is not equal to the burst oscillation frequency or half of it, this will overturn the apparently compelling arguments that burst brightness oscillations can only be explained via a rotational modulation mechanism."
 The two inodels aso differ in them predictions of tje frequencies. of other. weak. brightuCSS οscillaions.," The two models also differ in their predictions of the frequencies of other, weak brightness oscillations."
 In the relativistic precession uodel. there is an upper sidebaud oft10 OLvital ryequency expected (Psaltis Norman 2010).," In the relativistic precession model, there is an upper sideband of the orbital frequency expected (Psaltis Norman 2000)."
 The iuplitude of this oscillation necd uot be as great as the amplitide a the ower sidebaud: for exanuple. if the ταiderlvine perturvations of the accrction disk cut off sharply at yequencies CXCCCEine t16 observed orbital frequency tren the τα)per sideband could )o weaker.," The amplitude of this oscillation need not be as great as the amplitude at the lower sideband; for example, if the underlying perturbations of the accretion disk cut off sharply at frequencies exceeding the observed orbital frequency then the upper sideband could be weaker."
 Nonetheless. the oscillaion is expected to be there a some level.," Nonetheless, the oscillation is expected to be there at some level."
 In the cat frequency model there is an «wscillation expectec at twice the observed lower oeakk frequency., In the beat frequency model there is an oscillation expected at twice the observed lower peak frequency.
 This 1s expected Or several reasons., This is expected for several reasons.
 For exauple. the waveform eenerated by the beat is not expectc to be perfectly sinusoidal. and hence overtones will be generated.," For example, the waveform generated by the beat is not expected to be perfectly sinusoidal, and hence overtones will be generated."
 In acdition. if the radiation pattern rotating with the star. which iclps produce the beat. is itself no pertectly sinusoidal. then i will ecnerate weak overtones of the lower peak. frequeracy.," In addition, if the radiation pattern rotating with the star, which helps produce the beat, is itself not perfectly sinusoidal, then it will generate weak overtones of the lower peak frequency."
lot of damping because the panets were of course unbound eravitationally at the beeiunine.,lot of damping because the planets were of course unbound gravitationally at the beginning.
 The damping is here provided by tre disk. aud is visible between 5900 and 6000 vears in Figure 11 the imaxiuuni distance between the plaucts decreases until it becomes snaller than the Till radius.," The damping is here provided by the disk, and is visible between $5900$ and $6000$ years in Figure \ref{fig:UN}: the maximum distance between the planets decreases until it becomes smaller than the Hill radius."
 The ugh deusity of this nebula makes such a satellite capture possile., The high density of this nebula makes such a satellite capture possible.
 the dist:uice between Uranus aud Neoue ds slightly oscillating around 2 AU afer 600) voars., the distance between Uranus and Neptune is slightly oscillating around $2$ AU after $6000$ years.
 They share 1ο same orbit. with au eccenutriicitv πιαλος thiui LOL and Urauus lenaine about 30° before10: he right y-axis.," They share the same orbit, with an eccentricity smaller than $0.04$ and Uranus leading about $30^\circ$ before: the right $y$ -axis."
 The two plajets arc| not at a Lagrange »oiut of the other oue., The two planets are not at a Lagrange point of the other one.
 This unusual coufietvation mav be due to the mean motion resonaice with Saturn. or more probably o the strone dissipation bv the disk. that moclifies the local dyaandces.," This unusual configuration may be due to the mean motion resonance with Saturn, or more probably to the strong dissipation by the disk, that modifies the local dynamics."
 Ouce the disk will have evaporated. tjose two plaucts should have a tadpole motion around one of their LfLs Lagrange poiifs;," Once the disk will have evaporated, these two planets should have a tadpole motion around one of their $L_4/L_5$ Lagrange points."
 Such a ccufguration have already. been stucied. for instance by Thouuuesson (2008).," Such a configuration have already been studied, for instance by \citet{Thommes2005,Beauge-etal-2007,Cresswell-Nelson-2008}."
 GeneraIv. if ds asstuned that a eiu plaiot captures a nieratiue terrestrial planet. tha luav fieu erow.," Generally, it is assumed that a giant planet captures a migrating terrestrial planet, that may then grow."
 Iu coutrast. here the effect of the ejaut lanuet is to put two 15 Earth mass plaucts together ou the same orbit.," In contrast, here the effect of the giant planet is to put two 15 Earth mass planets together on the same orbit."
 This is a new wav of formingOo Trojan planets., This is a new way of forming Trojan planets.
 the distauceB between the two ice+ ejauts (shown as the | svinbols in Figure 1)) setles to about 2 times their seini-major axis., the distance between the two ice giants (shown as the $+$ symbols in Figure \ref{fig:2_lowres}) ) settles to about 2 times their semi-major axis.
 This indicates that the two planets are iu opposition. ibrating around their £3 poiit.," This indicates that the two planets are in opposition, librating around their $L_3$ point."
 Once again. this configuration is not possible i1 the absence of gas dissipation or of the other plajets; because the L5 »oiut is uustable.," Once again, this configuration is not possible in the absence of gas dissipation or of the other planets, because the $L_3$ point is unstable."
 Ift16 other plancts aud the disk disappear. these two dlanets should have coorbital rorseshoe orbits.," If the other planets and the disk disappear, these two planets should have coorbital horseshoe orbits."
 For comparison. fle Sane experiment as du section 2. has been perorued in the The equation of state is locally isothermal.," For comparison, the same experiment as in section \ref{sec:isoth} has been performed in the The equation of state is locally isothermal."
" The planes are initially placed ou cmalar orbits at their resent :: a;=5.2 AU. ay=9.6 AU, ap=19.2 AU. ay=30 AU."," The planets are initially placed on circular orbits at their present : $a_J=5.2$ AU, $a_S=9.6$ AU, $a_U=19.2$ AU, $a_N=30$ AU."
 The resolutiou is LO7. the 2D erid exteuds radially from 2.15 to 52 AU. while the 1D ο covers the range [0.312:101] AU.," The resolution is $10^{-2}$, the 2D grid extends radially from $2.15$ to $52$ AU, while the 1D grid covers the range $[0.312\,;\,104]$ AU."
 The planets do uotfee the disk poteutial during 100.Jupiter orbits (170years). and are then released.," The planets do notfeel the disk potential during 400Jupiter orbits (4740years), and are then released."
absolute magnitude -10 and -20 (for which extreme value there is no evidence).,absolute magnitude -10 and -20 (for which extreme value there is no evidence).
 Then there should still be a change in dispersion by a factor of ten: while apparently (here is none al all., Then there should still be a change in dispersion by a factor of ten; while apparently there is none at all.
 searching for more sophisticated predictions. one quickly encounters one major problem with concentrating on the Local Volume: it is much smaller than the regions dealt. with in studies of structure formation.," Searching for more sophisticated predictions, one quickly encounters one major problem with concentrating on the Local Volume: it is much smaller than the regions dealt with in studies of structure formation."
 In Jenkinsetal.(1998).. Figure 10. lor instance (which shows predicted velocity dispersions for various length scales). a volume of 10 Alpe radius is bevond (he small-scale edge of (he curves drawn.," In \citet{JF98}, Figure 10, for instance (which shows predicted velocity dispersions for various length scales), a volume of 10 Mpc radius is beyond the small-scale edge of the curves drawn."
 One can extrapolate them. but the result is quite uncertain: they might. predict a velocity dispersion of anvwhere [rom 50 to 500 km /.," One can extrapolate them, but the result is quite uncertain; they might predict a velocity dispersion of anywhere from 50 to 500 km $^{-1}$."
 Ol more use for the question at hand. Inagaki.Itoh&Saslaw(1992) compared the analvlic expression [or velocity dispersion in Saslawetal.(1990) with a series of η-ροςν caleulations.," Of more use for the question at hand, \citet{IIS92} compared the analytic expression for velocity dispersion in \citet{SCI2} with a series of n-body calculations."
 They found a rather small difference in rms velocity dispersion wilh mass: 169 compared to 179 (in scaled units) for sets of galaxies dillering in mass by a [factor of 100., They found a rather small difference in rms velocity dispersion with mass: 169 compared to 179 (in scaled units) for sets of galaxies differing in mass by a factor of 100.
 llowever. thev note that in their work the tendency [or smaller galaxies to be affected more by gravitational interactions is offset by the tendency for larger galaxies to be found in rich clusters.," However, they note that in their work the tendency for smaller galaxies to be affected more by gravitational interactions is offset by the tendency for larger galaxies to be found in rich clusters."
 Within the Local Volume there are no rich clusters., Within the Local Volume there are no rich clusters.
 For such an effect to operate here. one would have to have groups made up of similarlv-sized galaxies: cdwarl groups ancl ejant eroups.," For such an effect to operate here, one would have to have groups made up of similarly-sized galaxies: dwarf groups and giant groups."
 This is certainly not (rue. each group in (he Volume having galaxies with a vide distribution of size and brightness.," This is certainly not true, each group in the Volume having galaxies with a wide distribution of size and brightness."
 This brings us to the question ol field ancl group galaxies., This brings us to the question of field and group galaxies.
 Inagaki. predict that field galaxies should have much smaller peculiar velocities (han cluster galaxies., \citet{IIS92} predict that field galaxies should have much smaller peculiar velocities than cluster galaxies.
 This is a equite plausible prediction: field galaxies should have little or no interactions which take (hem away Irom the Hubble flow. while group (and cluster) galaxies should be at least in (he process of virializing.," This is a quite plausible prediction: field galaxies should have little or no interactions which take them away from the Hubble flow, while group (and cluster) galaxies should be at least in the process of virializing."
 The situation with the present sample is shown in Figures 19 and 20..," The situation with the present sample is shown in Figures \ref{group97}
 and \ref{group35}."
 There is no tendency for brighter galaxies to occur in groups. aud no obvious tendency for eroup galaxies to have a Larger velocity dispersion.," There is no tendency for brighter galaxies to occur in groups, and no obvious tendency for group galaxies to have a larger velocity dispersion."
 To compare (he situations quantitatively. we again use (he dispersions about the models and employ (cautiously) the F-ratio test.," To compare the situations quantitatively, we again use the dispersions about the models and employ (cautiously) the F-ratio test."
 For the 98-galaxy sample. isotropic model. the group dispersion is 120kins J| against a field dispersion of 113 kms |.," For the 98-galaxy sample, isotropic model, the group dispersion is 120km $^{-1}$ against a field dispersion of 113 km $^{-1}$."
 The F-ratio test gives a significance to this: the difference is probably not significant., The F-ratio test gives a significance to this; the difference is probably not significant.
 For the 98-ealaxy anisotropic model the eroup dispersion is 93 kin st against 122 km ! for the field. eiving a significancebut with thefield galaxies having a higher dispersion.," For the 98-galaxy anisotropic model the group dispersion is 93 km $^{-1}$ against 122 km $^{-1}$ for the field, giving a significance—but with the galaxies having a higher dispersion."
 For the 35-galaxy isotropic model the field and group dispersions are 68 and 98 km |. respectively. [or a significance of 8756: the anisotropic model has 47 and 86 km ! [or 82%...," For the 35-galaxy isotropic model the field and group dispersions are 68 and 98 km $^{-1}$, respectively, for a significance of ; the anisotropic model has 47 and 86 km $^{-1}$ for ."
The fluid state in ucielboring volumes ds formally interpreted as the initial data of a Riemann problem. the solution for which is frequently sought by iieaus of an exact or approximate Ricmanu solver aud vields the iutercell fluxes. F7.,"The fluid state in neighboring volumes is formally interpreted as the initial data of a Riemann problem, the solution for which is frequently sought by means of an exact or approximate Riemann solver and yields the intercell fluxes $\hat{\vct{F}}^j$."
 Iu this study. we have utilized the IILLD approximate Rican solver (?).. which computes Godunov fuxes at zone interfaces by resolving the fast. Alfvénn. and contact waves.," In this study, we have utilized the HLLD approximate Riemann solver \citep{Mignone:2009p2269}, which computes Godunov fluxes at zone interfaces by resolving the fast, Alfvénn, and contact waves."
 The inclusion of these waves in the Godunov flux approximation has been shown to be iniportaut for captuime the correct concentration of magnetic energy. relative to the more dciffusive TLLE aud IILLC approximate solvers (?)..," The inclusion of these waves in the Godunov flux approximation has been shown to be important for capturing the correct concentration of magnetic energy, relative to the more diffusive HLLE and HLLC approximate solvers \citep{Beckwith:2011p4448}."
 Due to the fact that Mara ciuplovs ILES. the (effectively umuerical) viscous aud resistive scales both occur near the eril spacing aud thus the magnetic Praudtl uuuber Pa=RinfRe~ 1.," Due to the fact that Mara employs ILES, the (effectively numerical) viscous and resistive scales both occur near the grid spacing and thus the magnetic Prandtl number $Pm=Rm/Re \sim 1$ ."
" Ποπονα, deviations tn Pin stil occur based on the muuericallv dissipative behavior of the scheme."," However, deviations in $Pm$ still occur based on the numerically dissipative behavior of the scheme."
 Based on qualitative observations. the IILLC' solver ecucrates roughly the sale amount of nuuercal diffusion to the deusitv aud pressure fields as OLED. but substantially wore to the magnetic field.," Based on qualitative observations, the HLLC solver generates roughly the same amount of numerical diffusion to the density and pressure fields as HLLD, but substantially more to the magnetic field."
 This is interpreted as the IILED solver achieving a larger numerical Pin., This is interpreted as the HLLD solver achieving a larger numerical $Pm$.
 It is also important to note that as the wmmerical dissipation goes clown. the chances of euncounteriue robustuess issies (sce 827)) goes up.," It is also important to note that as the numerical dissipation goes down, the chances of encountering robustness issues (see \ref{sec:robustness}) ) goes up."
 Tn this study. Mara uses voluuc-averaged. imuagnetic fields which are stored at cell centers.," In this study, Mara uses volume-averaged magnetic fields which are stored at cell centers."
 The solenoidal constraint V«B=0 (evaluated at cell corners) ds maintained to machine precision using the constrained transport method of ?..," The solenoidal constraint $\nabla \cdot
\vsp{B} = 0$ (evaluated at cell corners) is maintained to machine precision using the constrained transport method of \cite{Toth:2000p2162}."
 Due to the complex nature of the RMIID. equations. robustuess concerns are of eroat importance.," Due to the complex nature of the RMHD equations, robustness concerns are of great importance."
 Iu particular. RMIID. and CRATID codes are known to suffer frou faihwes when inverting Equation 3 in order to obtain the primitive variables from the couserved ones.," In particular, RMHD, and GRMHD codes are known to suffer from failures when inverting Equation \ref{eqn:conserved} in order to obtain the primitive variables from the conserved ones."
" This inversion cau fail when the numerical root finder (c.g, secant or Newton-Rapleson) fails to locate the root due to an insufficiently close initial euess.", This inversion can fail when the numerical root finder (e.g. secant or Newton-Rapheson) fails to locate the root due to an insufficiently close initial guess.
 But it cau also fail when the root corresponds to a state with negative pressure., But it can also fail when the root corresponds to a state with negative pressure.
 Overcoming these numerical limitations has been a iiajor source of effort in evolving the models used in this study., Overcoming these numerical limitations has been a major source of effort in evolving the models used in this study.
 The Mara code has been designed with au exteusivolv tested aud very robust alegorithim for the recovery of pruuitive variables., The Mara code has been designed with an extensively tested and very robust algorithm for the recovery of primitive variables.
 When the failure is strictly uunucerical in nature. it will run through a seres of reset values and inversion relations.," When the failure is strictly numerical in nature, it will run through a series of reset values and inversion relations."
 The default inversion relation used in this study is adapted from ? and in the case of an adiabatic equation of state. mmav be solved using a Newtou-Raphesou iteration in the single uukuowu. Z=phy.," The default inversion relation used in this study is adapted from \cite{Noble:2006p25} and in the case of an adiabatic equation of state, may be solved using a Newton-Rapheson iteration in the single unknown, $Z \equiv \rho h
W^2$."
 The starting value of Z given to the root finder is evaluated frou the primitive variables at the previous time step., The starting value of $Z$ given to the root finder is evaluated from the primitive variables at the previous time step.
" If a suitable solution is not obtained. then the root finder is restarted using Z—v/D2|S2, which becomes the exact solution iu the limiting case of weal magnetic feld aud siall pressure."," If a suitable solution is not obtained, then the root finder is restarted using $Z =
\sqrt{D^2 + S^2}$, which becomes the exact solution in the limiting case of weak magnetic field and small pressure."
 If a solution is still uot obtained. then the procedure is repeated using au inversion relation adapted from ?..," If a solution is still not obtained, then the procedure is repeated using an inversion relation adapted from \cite{Anton:2006p17}."
 which ean solved using a two-dimensional Newtou-Raphesou algorithin for the uknownus Z aud Wo., which can solved using a two-dimensional Newton-Rapheson algorithm for the unknowns $Z$ and $W$.
 Because ucither of these solvers contains the other one's domam of success. there are circumstances where this procedure obtains the solution even when the first solver or starting value fails.," Because neither of these solvers contains the other one's domain of success, there are circumstances where this procedure obtains the solution even when the first solver or starting value fails."
 Towever. there are states whose solution is not obtained by any solver. or whose solution is unacceptable due to negative pressure.," However, there are states whose solution is not obtained by any solver, or whose solution is unacceptable due to negative pressure."
 Under these conditions. the code assumes dt has integrated the conserved quantities iuto iu uuplhlvsical configuration. and requires further safety procedures m order not to crash.," Under these conditions, the code assumes it has integrated the conserved quantities into an unphysical configuration, and requires further safety procedures in order not to crash."
 Thesafety feature we have found to be the most practical is the addition of swall amounts, Thesafety feature we have found to be the most practical is the addition of small amounts
(between the first two epochs).,(between the first two epochs).
 AM Fais=free. the radio core is inferred. to be at the centre: angular separation ~174 mas and position angle ~1307.," At $t_\mathrm{app}=t_\mathrm{rec}$, the radio core is inferred to be at the centre: angular separation $\sim174$ mas and position angle $\sim-130\degr$."
 The radio core was not detected. during any of the VLBI epochs but. since all four observations took place during the X-ray soft state. this is to be expected. according to the unified model of Fenderetal.(2004.2009).," The radio core was not detected during any of the VLBI epochs but, since all four observations took place during the X-ray soft state, this is to be expected, according to the unified model of \citet{fen04, fen09}."
.. Ehe. radio position confirms that its optical counterpart is Suiff--UVO source A (Curranetal.2010)., The radio position confirms that its optical counterpart is -UVOT source A \citep{cur10}.
. If we take the Lux density at the first epoch as the upper limit of S55. then μμuppκDr anc σοςxO.2e<3(0.ἐς)Cos@. in agreement with the jet deceleration scenario on both sides.," If we take the flux density at the first epoch as the upper limit of $S_\mathrm{app}$, then $\frac{S_\mathrm{app}}{S_\mathrm{rec}}\leq5$ and $\beta\cos\theta\leq0.2c < \bar\beta(0,t_\mathrm{rec})\cos\theta$, in agreement with the jet deceleration scenario on both sides."
" The observed. [ux density from the receding jet is deboosted by a factor: 57. where o4,d=(11)""""(1Lyte|dcos8)"," The observed flux density from the receding jet is deboosted by a factor: $\delta_\mathrm{rec}^{3-\alpha}$, where $\delta_\mathrm{rec}=(1-\beta^2)^{0.5}(1+\beta\cos\theta)^{-1}$."
 Because of the jet deceleration. the receding ejecta is less beamed away from our line of sight. and thus it looks relatively brighter.," Because of the jet deceleration, the receding ejecta is less beamed away from our line of sight, and thus it looks relatively brighter."
 Note that the non-detection of the receding jet component at the first epoch may also be because it stavec at an earlier stage (oooO5fapy dE cos9= 0.3) and its Dux density was still much lower than the peak Hus density of the racio Dare., Note that the non-detection of the receding jet component at the first epoch may also be because it stayed at an earlier stage $t_\mathrm{rec}\sim0.5t_\mathrm{app}$ if $\beta\cos\theta=0.3$ ) and its flux density was still much lower than the peak flux density of the radio flare.
 The caveat in the above argument is that component. D. might have not followed the same luminosity evolution mocdoel as component A. It is also possible that the brightening ofthe receding ejecta was due to sudden jet-cloud. interaction. as inthe case of NTI2JI1748 | 228 (οπαςetal.1998:Jonesetal. 2008).," The caveat in the above argument is that component B might have not followed the same luminosity evolution model as component A. It is also possible that the brightening of the receding ejecta was due to sudden jet-cloud interaction, as in the case of XTE $-$ 228 \citep{hje98, mil08}."
. In this paper. we present the results of the first ΝΕΟΙ observations of the new Calactic black hole candidate NTE 223 during its first known outburst.," In this paper, we present the results of the first VLBI observations of the new Galactic black hole candidate XTE $-$ 223 during its first known outburst."
 With EWN and VLBA observations at four epochs in 2010 February. we imaged its radio counterpart at 5 12.," With EVN and VLBA observations at four epochs in 2010 February, we imaged its radio counterpart at 5 GHz."
 We detect an ejected component at the first three epochs and find that its proper motion shows significant deviation from the uniform proper motion mocel and requires a deceleration. rate of 0.34+0.02 mas 7., We detect an ejected component at the first three epochs and find that its proper motion shows significant deviation from the uniform proper motion model and requires a deceleration rate of $0.34\pm0.02$ mas $^{-2}$.
 In the jet deceleration scenario. it has proper motion decelerating from 9.2 mas | at the first epoch to 4.0 mas * at the last epoch.," In the jet deceleration scenario, it has proper motion decelerating from 9.2 mas $^{-1}$ at the first epoch to 4.0 mas $^{-1}$ at the last epoch."
 Lt also shows slow but detectable variation of its transverse size indicating that its expansion is also significantly confined., It also shows slow but detectable variation of its transverse size indicating that its expansion is also significantly confined.
 This is the first time that a Galactic jet is found to be σαςπαν decelerating on the hundred. milliaresecond scale., This is the first time that a Galactic jet is found to be gradually decelerating on the hundred milliarcsecond scale.
 The discovery provides strong evidence for the existence of significant interaction around the jet at an early stage of its evolution., The discovery provides strong evidence for the existence of significant interaction around the jet at an early stage of its evolution.
 In addition to the approaching jet component. we detect another jet feature at the last epoch. which is most likely associated with the receding ejecta.," In addition to the approaching jet component, we detect another jet feature at the last epoch, which is most likely associated with the receding ejecta."
 We infer that the jet deceleration should start at a time much earlier than our VLBI observations using the birth clate (around 2010 February 2) [rom the ATCA radio light curve., We infer that the jet deceleration should start at a time much earlier than our VLBI observations using the birth date (around 2010 February 2) from the ATCA radio light curve.
 Furthermore. we interpret the detection of the receding ejecta as a result of the reduced. Doppler deboosting elfect caused by the jet deceleration on the receding side and give a lower limit of 0.96 for the jet birth speed assuming symmetric jet motion.," Furthermore, we interpret the detection of the receding ejecta as a result of the reduced Doppler deboosting effect caused by the jet deceleration on the receding side and give a lower limit of $0.3c$ for the jet birth speed assuming symmetric jet motion."
 lt has been reported. by Shaposhnikoyetal.(2010). that the distance. estimated. by the spectral-timing correlation scaling technique. is around 3.5 κρο," It has been reported by \citet{sha10} that the distance, estimated by the spectral-timing correlation scaling technique, is around 3.5 kpc."
 Thus. the proper motion observed at the first epoch would correspond to an apparent jet speed of ~0.26. in agreement with our results (but note that the technique is very model dependent)," Thus, the proper motion observed at the first epoch would correspond to an apparent jet speed of $\sim0.2c$, in agreement with our results (but note that the technique is very model dependent)."
 We thank the EVN PC Chair. “P. Venturi and. the VLBA Proposal Selection Committee lor prompt approval of our ToO requests.," We thank the EVN PC Chair, T. Venturi and the VLBA Proposal Selection Committee for prompt approval of our ToO requests."
" eVLIBI developments in trope were supported by the LEC DCO-INISO IPunded Communication Network Developments xojeect. IENPIeS. The National Badio Astronomy Observatory is a lacility of the National Science Foundation operated under Cooperative agreement by Assuciated Universities. lie. The European VLBI Network is a joint lacility of European. Chinese. South Albican and. other radio astronomy institutes Πάσα, by heir national research councils."," e-VLBI developments in Europe were supported by the EC DG-INFSO funded Communication Network Developments project EXPReS. The National Radio Astronomy Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc. The European VLBI Network is a joint facility of European, Chinese, South African and other radio astronomy institutes funded by their national research councils."
understanding of its origin has vet been established. (see Wijers&Pringle 1999)).,understanding of its origin has yet been established (see \pcite{wijers99}) ).
 As noted by Nowak.Wilms&Dove(1999).. the observed long-term. X-ray. variability. of 33394 could be due to a combination of a quasi-steadily precessing disc at large radii with coronal structure changes on small radii.," As noted by \scite{nowak99}, the observed long-term X-ray variability of 339–4 could be due to a combination of a quasi-steadily precessing disc at large radii with coronal structure changes on small radii."
 Hf this is the case. X-ray irradiation can give rise to a precessing disc (Dubusctal.1999). and it might even explain the LS/LIS transition of the source.," If this is the case, X-ray irradiation can give rise to a precessing disc \cite{dubus99} and it might even explain the LS/HS transition of the source."
 Actually. results [rom the aand2D data suggest that the LS/1IS transitions and the outburst-like events during the Ίοπαντα state (as seen in the aand BATSE data) may be due to the same mechanism. although the data is not. well-samiplect.," Actually, results from the and data suggest that the LS/HS transitions and the outburst-like events during the low/hard state (as seen in the and BATSE data) may be due to the same mechanism, although the data is not well-sampled."
 Therefore. it ds possible that such quasi-periocic outbursts might actually bea characteristic precursor to the LS/IIS transition.," Therefore, it is possible that such quasi-periodic outbursts might actually be a characteristic precursor to the LS/HS transition."
 In the next subsection. we will discuss the role of X-ray irradiation in the state transition and also the origin of the long-term X-ray variability.," In the next subsection, we will discuss the role of X-ray irradiation in the state transition and also the origin of the long-term X-ray variability."
" Our oobservations presented here indicate that the source was fainter than the previous observation in. 1999. August (see Ixongetal. 2000)) by a factor 3 in half vear: therefore the source is still in the ""oll, state.", Our observations presented here indicate that the source was fainter than the previous observation in 1999 August (see \pcite{kong00}) ) by a factor 3 in half year; therefore the source is still in the `off' state.
 The energy spectra of both observations are very similar to those obtained in the LS., The energy spectra of both observations are very similar to those obtained in the LS.
" Llenee. there is no doubt that the ""oll, state is an extended LS (see discussions by Koneetal.2000:Corbel 2000)) and we may now refine the source as a transient rather than a persistent source."," Hence, there is no doubt that the `off' state is an extended LS (see discussions by \pcite{kong00,corbel00}) ) and we may now refine the source as a transient rather than a persistent source."
 In fact the luminosity during the ol state is in the upper end of the quiescent Luminosity of typical N-rav. transients (see e.g. Ixongctal. 2000))., In fact the luminosity during the `off' state is in the upper end of the quiescent luminosity of typical X-ray transients (see e.g. \pcite{kong00}) ).
 The results also further confirm that cluring quiescence. significant X-ray variability can be seen.," The results also further confirm that during quiescence, significant X-ray variability can be seen."
 Such variability in quiescence is also found in several black hole transients such as 22023|338 (Wagnerctal.100: Kongetal. 2001)). 1163047 (Parmaretal.LO9Tb) ancl AOGZOOO (Αναetal.1998:Menou 1999).," Such variability in quiescence is also found in several black hole transients such as 2023+338 \pcite{wagner94}; \pcite{kong01}) ), 1630–47 \cite{parmar97} and A0620–00 \cite{asai98,menou99}."
. As noted in and 2.5. 4 underwent a LS/LIS transition in 1998— (see Eig. 3::," As noted in \\ref{sec:339rxte} and \ref{sec:batse}, 339–4 underwent a LS/HS transition in 1998 (see Fig. \ref{fig:batse};"
 also 3339.Bellonietal. 1999)) which differs somewhat from the previous ones observed with the AASAL., also \pcite{belloni99}) ) which differs somewhat from the previous ones observed with the ASM.
" After the source reached the peak of the LIS. it spent nearly 200 days at a constant [ux level of  20 TE//ASM cts/s. Such a flat-topped X-ray. light curve for 33394 is dillerent. from the ""standard! fast-rise exponential decay profile. as observed from various X-ray transicnts (see Chen.Shrader&Livio1997. for various examples of cillerent outburst patterns)."," After the source reached the peak of the HS, it spent nearly 200 days at a constant flux level of $\sim$ 20 /ASM cts/s. Such a flat-topped X-ray light curve for 339–4 is different from the `standard' fast-rise exponential decay profile, as observed from various X-ray transients (see \pcite{chen97} for various examples of different outburst patterns)."
 Phe state transition mechanism for such a persistent source is still a puzzle and observations in different states sugeested that it is most likely. associated with the mass accretion rate (e.g. Méndez&vander 1997))., The state transition mechanism for such a persistent source is still a puzzle and observations in different states suggested that it is most likely associated with the mass accretion rate (e.g. \pcite{mendez97}) ).
 A change in the mass accretion rate can be caused by a thermal instability in the accretion clise (disc instability model. or DIM) and. in fact it is widely accepted that the outbursts in BLSN'Ts are due to this mechanism (e.g. Esin.Lasota&IHlvnes 2000)).," A change in the mass accretion rate can be caused by a thermal instability in the accretion disc (disc instability model, or DIM) and in fact it is widely accepted that the outbursts in BHSXTs are due to this mechanism (e.g. \pcite{esin00}) )."
 Interestinglv. the DIAL was originally developed to describe the outbursts in να novae (DN: see e.g. Cannizzo1993. [or a review) which contain a white dwarf as the compact object: it was suggested. quite carly that it also applies in DIISNTs (e.g. van.ParaclijsWheeler1989)).," Interestingly, the DIM was originally developed to describe the outbursts in dwarf novae (DN; see e.g. \pcite{cannizzo93} for a review) which contain a white dwarf as the compact object; it was suggested quite early that it also applies in BHSXTs (e.g. \pcite{paradijs84,huang89,mineshige89}) )."
 In LAINBs. X-ray irracliation can inlluence the stability of the accretion disc (vanParaclijsL996) and thus a DIM which includes X-ray irraciiation was developed (Dubusetal.1999).," In LMXBs, X-ray irradiation can influence the stability of the accretion disc \cite{paradijs96}
 and thus a DIM which includes X-ray irradiation was developed \cite{dubus99}."
 Ht is worth noting that the recent LIS spectroscopic observations of 3339.4 obtained by Wuetal.(2001) also indicate that the accretion dise is heated by soft. X-ray. irracliation. suggesting that irradiation plays à role in the state transition.," It is worth noting that the recent HS spectroscopic observations of 339–4 obtained by \scite{wu00} also indicate that the accretion disc is heated by soft X-ray irradiation, suggesting that irradiation plays a role in the state transition."
" Irradiation of the outer disc by X-ravs [rom the inner disc gives a stability criterion. [rom which the minimum mass accretion rate for a stable. steady-state accretion in an. X-ray irradiated disc can be expressed as (Dubusetal.1999): where AZ, is the black hole mass. Ado is the mass of the companion. {νε the orbital period in hours."," Irradiation of the outer disc by X-rays from the inner disc gives a stability criterion, from which the minimum mass accretion rate for a stable, steady-state accretion in an X-ray irradiated disc can be expressed as \cite{dubus99}: where $M_1$ is the black hole mass, $M_2$ is the mass of the companion, $P_{orb}$ the orbital period in hours."
" C is à parameter to describe the properties of irradiation through: where 25,4 is the irradiation temperature and c is the Stefan-Doltzmann constant.", $\mathcal{C}$ is a parameter to describe the properties of irradiation through: where $T_{irr}$ is the irradiation temperature and $\sigma$ is the Stefan-Boltzmann constant.
 For a stable X-ray irracliated disc during the LIS. AIeredint can be estimated. [rom the LIS observations by Bellonictal.(1999)... in which they have fitted the spectrum with a power law plus a multi-colour disc blackbody model.," For a stable X-ray irradiated disc during the HS, $\mdot_{crit}^{irr}$ can be estimated from the HS observations by \scite{belloni99}, in which they have fitted the spectrum with a power law plus a multi-colour disc blackbody model."
" By using the derived. inner cise radius. Ry,voos? of 25.40) km (where D4 isthe distance to the source in units of 4 kpe: Zdziarskictal. 1998)) and the blackbody temperature. 2;, of 0.72 keV. we can calculate the bolometric Luminosity of the dise blackbody component as The acerction rate can be estimated: via the relation Lx—056CM,AL/R;,,. which leads to MzL55107 Ότο + if we assume an orbital inclination of about 15° etal.2001). for a 14.8-hr orbital period. (Callananctal. 1992)."," By using the derived inner disc radius, $R_{in}\sqrt{\cos i}$ of $25.4 D_{4}$ km (where $D_4$ isthe distance to the source in units of 4 kpc; \pcite{zdziarski98}) ) and the blackbody temperature, $T_{in}$ of 0.72 keV, we can calculate the bolometric luminosity of the disc blackbody component as The accretion rate can be estimated via the relation $L_X=0.5
GM_1\mdot /R_{in}$, which leads to $\mdot\approx 1.7\times 10^{17}
D_4^2$ g $^{-1}$ if we assume an orbital inclination of about $15^{\circ}$ \cite{wu00} for a 14.8-hr orbital period \cite{callanan92}."
. Assuming an estimated mass of the black hole of and companion star of citecallanan92.. we obtain C=9.4.10. from Eqn.," Assuming an estimated mass of the black hole of \\cite{zdziarski98} and companion star of \\cite{callanan92}, we obtain $\mathcal{C}=9.4\times10^{-4}$ from Eqn."
 1. which is slightly higher than 5101 obtained. bv comparison. with other results in the literature (see Dubusetal. 1999)).," 1, which is slightly higher than $5\times10^{-4}$ obtained by comparison with other results in the literature (see \pcite{dubus99}) )."
 According to the scenario outlined. above. the mass of the compact object is limited to he ~ SAL...," According to the scenario outlined above, the mass of the compact object is limited to be $\sim$ ."
 From the previous LS observations. the estimated mass accretion rate is about 5«LOL & 4 (eg. Wilmsetal. 1999)) which is a factor of 3. below the critical accretion rate for a stable disc and hence the disc may be subject to the DIM in order to give rise to a," From the previous LS observations, the estimated mass accretion rate is about $5\times 10^{16}$ g $^{-1}$ (e.g. \pcite{wilms99,belloni99}) ) which is a factor of 3 below the critical accretion rate for a stable disc and hence the disc may be subject to the DIM in order to give rise to a"
alone —1007 at ~Smes.,along $-100\degr$ at $\sim8{\em mas}$.
 This is probably a poor estimate but the indication is that the magnetic Ποιά ecometry [ies roughly parallel to the jet when the polarisation is high., This is probably a poor estimate but the indication is that the magnetic field geometry lies roughly parallel to the jet when the polarisation is high.
 This is à low declination source not frequently observed., This is a low declination source not frequently observed.
 The only observation obtained gives a [lux density of 1.1 Jy and »olarisation of to better than 3o., The only observation obtained gives a flux density of 1.1 Jy and polarisation of to better than $3\sigma$.
 The position angle of the jet has only been estimated from a VLA map with insullicient resolution close to the core to try to establish he orientation of the magnetico field., The position angle of the jet has only been estimated from a VLA map with insufficient resolution close to the core to try to establish the orientation of the magnetic field.
 Dhis. measurement gives a very Hat spectral index. in broad. agreement with he shallower than average spectral index values obtained on his source by Gear (1994) over three dillerent epochs.," This measurement gives a very flat spectral index, in broad agreement with the shallower than average spectral index values obtained on this source by Gear (1994) over three different epochs."
 ‘The single observation available on this source was obtained with a Εαν density «0.9 Jy. however the polarisation was high at SA with a signal to noise of 3o.," The single observation available on this source was obtained with a flux density $<$ 0.9 Jy, however the polarisation was high at $\sim 8\%$ with a signal to noise of $3\sigma$."
 The polarisation position angle indicates a possibly perpendicular magnetic Ποιά relative to a jet orientation of 45 determined from a low-frequeney (5 Gllz) VEDI map., The polarisation position angle indicates a possibly perpendicular magnetic field relative to a jet orientation of $-45\degr$ determined from a low-frequency (5 GHz) VLBI map.
 This source. has been observed. eight times., This source has been observed eight times.
 The degree of polarisation changes readily between the extremes of θ andσα. averaging to over all epochs.," The degree of polarisation changes readily between the extremes of 0.7 and, averaging to over all epochs."
 The ας variability is very dillerent. since a Large flare at the beginning of the observations (when the Dux increases [ron 1.6 t0 5.8 Jv) is followed by a slow decay to its pre-IHare level over the subsequent three vears.," The flux variability is very different, since a large flare at the beginning of the observations (when the flux increases from 1.6 to 5.8 Jy) is followed by a slow decay to its pre-flare level over the subsequent three years."
 Phe jet is known to be very twisted and it is not. possible to determine its orientation close to the core with any certainty., The jet is known to be very twisted and it is not possible to determine its orientation close to the core with any certainty.
 Lligh-frequency VLBI images show that emerging components follow paths in the range 135 60., High-frequency VLBI images show that emerging components follow paths in the range $-135$ – $-60\degr$.
 Phe polarisation position angle changes bv tens of degrees between epochs., The polarisation position angle changes by tens of degrees between epochs.
 Using a recent value [or (vigi;=55° (obtained at very. high frequency) the magnetic field. is found. within 30 of alignment with the jet on six epochs. but clearly perpendicular to it on one occasion.," Using a recent value for $\theta_{VLBI} = -55\degr$ (obtained at very high frequency) the magnetic field is found within $30\degr$ of alignment with the jet on six epochs, but clearly perpendicular to it on one occasion."
 There is a 907 dillerence in position angle between the first and last polarisation epochs. when the polarisation was close to in both cases but the [Lux measured was 5.8 and. 1.9 Jy. respectively.," There is a $90\degr$ difference in position angle between the first and last polarisation epochs, when the polarisation was close to in both cases but the flux measured was 5.8 and 1.9 Jy respectively."
 There is a single observation on this source with just uncer Ac signal to noise., There is a single observation on this source with just under $3\sigma$ signal to noise.
 Ehe polarisation is with a position anele of 126 and the flux density is 1.16 Jy., The polarisation is with a position angle of $126\degr$ and the flux density is 1.16 Jy.
 The inferred magnetic field orientation is parallel to the jet with 155 (recently observed with VLBI at 22 112)., The inferred magnetic field orientation is parallel to the jet with $\theta_{VLBI} = -155\degr$ (recently observed with VLBI at 22 GHz).
 ‘This source was observed on three occasions over LS months., This source was observed on three occasions over 18 months.
 The polarisation rises [rom 3.7 to while the Lux drops [rom 4.6 to 1.9 Jv. i.e. they are anti-correlated.," The polarisation rises from 3.7 to while the flux drops from 4.6 to 1.9 Jy, i.e. they are anti-correlated."
 The position angle varies within a 50r. range. but the compactness of the core makes the jet. position angle 65° used somewhat uncertain and the magnetic field geometry dillieult to infer.," The position angle varies within a $50\degr$ range, but the compactness of the core makes the jet position angle $65\degr$ used somewhat uncertain and the magnetic field geometry difficult to infer."
 There is only one epoch on this source which is highly polarised at and with a flux density of 1.1 Jy., There is only one epoch on this source which is highly polarised at and with a flux density of 1.1 Jy.
 The magnetic field seems to lie within 25° from the perpendicular to a jet orientation of —166 determined from 5 Cllz maps., The magnetic field seems to lie within $25\degr$ from the perpendicular to a jet orientation of $-166\degr$ determined from 5 GHz maps.
 This low declination source was observed. on four cilferent epochs., This low declination source was observed on four different epochs.
 Lt is very bright. with a flux density ranging from 4.7 to 7 Js. and highly polarised with a mean polarisation level ofO.," It is very bright, with a flux density ranging from 4.7 to 7 Jy, and highly polarised with a mean polarisation level of."
1%.. The position angle changes bv tens of degrees between epochs but the magnetic field remains largely perpenclicular to the jet (only one VLBI map is available but it has been obtained very recently at 43 1111)., The position angle changes by tens of degrees between epochs but the magnetic field remains largely perpendicular to the jet (only one VLBI map is available but it has been obtained very recently at 43 GHz).
 The tus ancl degree of polarisation are uncorrelated., The flux and degree of polarisation are uncorrelated.
 DL Lacertae has been observed six times., BL Lacertae has been observed six times.
 Ht is not brighter han the other BL Lacs with a mean us of LS Jv (varving tween 1 and 2.6 Jv) but it is the most highly. polarised with a mean polarisation ofS., It is not brighter than the other BL Lacs with a mean flux of 1.8 Jy (varying between 1 and 2.6 Jy) but it is the most highly polarised with a mean polarisation of.
1'A.. The changes in flux and percentage polarisation are anti-correlated with r=0.77 at confidence. level., The changes in flux and percentage polarisation are anti-correlated with $r=-0.77$ at confidence level.
 “Phe polarisation position angle is the same on four of the six epochs. thus remaining arecly within a level consistent with a magnetic. field geometry perpendicular to the jet.," The polarisation position angle is the same on four of the six epochs, thus remaining largely within a level consistent with a magnetic field geometry perpendicular to the jet."
 The jet orientation. is well determined from several high-frequency. VLBI maps., The jet orientation is well determined from several high-frequency VLBI maps.
 There are two epochs on this source. taken more than two vears apart. between which the flux density and. the polarisation both drop significantly from 4.5 to 1.1 Jy. and from 3.8 to respectively.," There are two epochs on this source, taken more than two years apart, between which the flux density and the polarisation both drop significantly from 4.5 to 1.1 Jy and from 3.8 to respectively."
 The position angle of the jet is 218 (from a high resolution 100 Gilz VLBI map) and sugeests à perpendicular magnetic. field. when the source Hux is high. deviating from this position by 50° when the emission is low.," The position angle of the jet is $218\degr$ (from a high resolution 100 GHz VLBI map) and suggests a perpendicular magnetic field when the source flux is high, deviating from this position by $50\degr$ when the emission is low."
 This source. has been measured six times., This source has been measured six times.
 The lux. and polarisation lie in the ranges 3.8.6.7 Jv and 29 respectively. and are very. strongly correlated. (=0.94 with a cconfidence level).," The flux and polarisation lie in the ranges 3.8 – 6.7 Jy and 2.9 – respectively, and are very strongly correlated $r=0.94$ with a confidence level)."
 The position angle changes reacily within a 290 range., The position angle changes readily within a $>90\degr$ range.
 The epoch of lowest Ilux and minimum polarisation coincide with a magnetic field aligned with the jet. which has a position angle of 65° ," The epoch of lowest flux and minimum polarisation coincide with a magnetic field aligned with the jet, which has a position angle of $-65\degr$ "
resolution clement. so pushing to very lieh resolution will not lead to much new information.,"resolution element, so pushing to very high resolution will not lead to much new information."
 Alternatively. the central Compton parameter can be estimated by fitting a model (e.g.. the isothermal .} model) to the the observed SZ effect surface brightuess aud extrapolating it to the cluster center (as discussed briefly in 83).," Alternatively, the central Compton parameter can be estimated by fitting a model (e.g., the isothermal $\beta$ model) to the the observed SZ effect surface brightness and extrapolating it to the cluster center (as discussed briefly in 3)."
 This method is widely used to estimate yy (e.g. Carlstrou et al," This method is widely used to estimate $y_0$ (e.g., Carlstrom et al."
 1996: Tlolzaptel et al., 1996; Holzapfel et al.
 1997: (ποσο et al., 1997; Grego et al.
 2000: 2001: Reese et al., 2000; 2001; Reese et al.
 x2000: 2002: Pointecouteau et al., 2000; 2002; Pointecouteau et al.
 2001: 2002: Jones et al., 2001; 2002; Jones et al.
 2002: Ciraiunge et al., 2002; Grainge et al.
 2002)., 2002).
" While thestatistical measurement error on yy for current data is typically only of order 100 μμ @vhich is relatively simall compared to the differences between the models plotted in Figure 2. assuming S,/f, is kuowu). it has vet to be demonstrated that the error (duc to assuniug an incorrect surface brightuess model) for current data is neeheible."," While the measurement error on $y_0$ for current data is typically only of order $100$ $\mu$ K (which is relatively small compared to the differences between the models plotted in Figure 2, assuming $S_{\nu}/f_{\nu}$ is known), it has yet to be demonstrated that the error (due to assuming an incorrect surface brightness model) for current data is negligible."
 If the models eiploved iu the extrapolation provide poor descriptions of the surface brielhtuess profiles. then one would expect the results to vary as a function of mstrinent characteristics (e.g... resolution. field of view).," If the models employed in the extrapolation provide poor descriptions of the surface brightness profiles, then one would expect the results to vary as a function of instrument characteristics (e.g., resolution, field of view)."
 However. a comparison of the results of various studies (which mace use of different instruments: e... telescope.BIMA/OVRO. aud SuZIE) of the same clusters reveals that the agrecinent is quite good. often within oue sigma statistical uncertainty (see Iolzapfel et al.," However, a comparison of the results of various studies (which made use of different instruments; e.g., telescope, and ) of the same clusters reveals that the agreement is quite good, often within one sigma statistical uncertainty (see Holzapfel et al."
 1997. Reese et al.," 1997, Reese et al."
 2002. and Jones et al.," 2002, and Jones et al."
 2002. for exanple).," 2002, for example)."
 Thus. the extrapolation procedure seeuis to be a viable wav cstimating the ceutral Compton parameter at current sensitivity.," Thus, the extrapolation procedure seems to be a viable way estimating the central Compton parameter at current sensitivity."
" By modeling “mock” (future) observations. we also demonstrated that this extrapolation procedure is an accurate way of estimating the true uuderlving values of yy aud 5, of the BBLP02 cluster models (see MeCarthyv et al."," By modeling “mock” (future) observations, we also demonstrated that this extrapolation procedure is an accurate way of estimating the true underlying values of $y_0$ and $S_{\nu}$ of the BBLP02 cluster models (see McCarthy et al."
 2003)., 2003).
" Finally, we point out that part of the motivation for our studv of the S$,yy relations (aud the other scaling relatious involving these quantities) comes frou the fact that observational studies often use either ὃν or yy to characterize a clusters SZ effect but not usually both."," Finally, we point out that part of the motivation for our study of the $S_{\nu}-y_0$ relations (and the other scaling relations involving these quantities) comes from the fact that observational studies often use either $S_{\nu}$ or $y_0$ to characterize a cluster's SZ effect but not usually both."
 For example. integrated SZ effect flux densities are often associated with Iarge-beam sinele-dish experiments. while estimates of the central Compton parameter normally cole from high resolution interferometric observations.," For example, integrated SZ effect flux densities are often associated with large-beam single-dish experiments, while estimates of the central Compton parameter normally come from high resolution interferometric observations."
" Ilowever. if oue is able to determine (0) (160. the SZ effect ""surface brightuess profile) frou: a single dataset hen. strictly speaking. it Euf necessary to use scaling relations to probe the eutropv of the ICAL"," However, if one is able to determine $y(\theta)$ (i.e., the SZ effect “surface brightness” profile) from a single dataset then, strictly speaking, it isn't necessary to use scaling relations to probe the entropy of the ICM."
 For example. he SZ effect surface brightuess profile could be used ogether with the XN-rav surface brightuess profile (if it is known) to determine the true 3-cimenusional cutropy distribution of the eas (with some asstuptions about he geometry of the cluster).," For example, the SZ effect surface brightness profile could be used together with the X-ray surface brightness profile (if it is known) to determine the true 3-dimensional entropy distribution of the gas (with some assumptions about the geometry of the cluster)."
 The disadvantages of this uethod are: (1) it requires X-rav observations. aud (2) it probably cannot be used for the most distant clusters. since the X-ray signal-to-noise ratio falls sharply with increasing +.," The disadvantages of this method are: (1) it requires X-ray observations, and (2) it probably cannot be used for the most distant clusters, since the X-ray signal-to-noise ratio falls sharply with increasing $z$."
 Alternatively. the entropy distribution could be constrained by comparing the observed SZ effect surface brightness profile with theoretically predicted profiles.," Alternatively, the entropy distribution could be constrained by comparing the observed SZ effect surface brightness profile with theoretically predicted profiles."
 This is similar to using scaling relations. since the cutropy distribution is being inferred aud not measured. but with the inuportaut difference that all of the available SZ effect information. g(0). is being used in the comparison.," This is similar to using scaling relations, since the entropy distribution is being inferred and not measured, but with the important difference that all of the available SZ effect information, $y(\theta)$, is being used in the comparison."
 We are in the process of exploring these methods and they will be addressed in detail iu a subsequent paper., We are in the process of exploring these methods and they will be addressed in detail in a subsequent paper.
 For now. we stick with the scaling relatious derived above. which are more readily comparable with available observations.," For now, we stick with the scaling relations derived above, which are more readily comparable with available observations."
" Οτι Iu Figure 3. we present scaling relations between yy aud the total cluster dark matter mass within the radius roy ου, Wrsou)|."," 0.1in In Figure 3, we present scaling relations between $y_0$ and the total cluster dark matter mass within the radius $r_{500}$ [i.e., $M(r_{500})$ ]."
 This is the radius within which the mean dark matter mass density is 500 times the critical deusitv at ; = 0., This is the radius within which the mean dark matter mass density is 500 times the critical density at $z$ = 0.
 The lines hold the same meaning as in Figure 2., The lines hold the same meaning as in Figure 2.
 For clarity. we plot the +=0.2 predictions in the Ieft-haud panel aud the +=1.0 predicts iu the right-hand panel.," For clarity, we plot the $z = 
0.2$ predictions in the left-hand panel and the $z = 1.0$ predicts in the right-hand panel."
 It is appareut that cutropy injection has a substautial effect on the yy—M(rsoo) relation., It is apparent that entropy injection has a substantial effect on the $y_0-M(r_{500})$ relation.
 Both the normalization and the steepuess of the relation are modified., Both the normalization and the steepness of the relation are modified.
 First. the normnalizatious nuüplv that. for a cluster of given mass at a giveu redshift. cutropy injection tends to diminish the streneth of wo.," First, the normalizations imply that, for a cluster of given mass at a given redshift, entropy injection tends to diminish the strength of $y_0$."
 For example. at 2=0.2. a cluster with Mirsgg)29«LOMAL. will have yyz ος absence of an eutropy floor. but the central Compton parameter is ouly half this value for au eutropy floor of AgcmBO keV ο”.," For example, at $z = 0.2$, a cluster with $M(r_{500}) \approx 9 \times 10^{14} M_{\odot}$ will have $y_0 \approx 4 
\times 10^{-4}$ in the absence of an entropy floor, but the central Compton parameter is only half this value for an entropy floor of $K_0 \approx 430$ keV $^2$."
 This correspouds to a difference of neatly 10604A at 30 CIIz (Gvhich is large compared to the typical statistical measurement error of 106020048: Reese et al., This corresponds to a difference of nearly $1060 \mu K$ at 30 GHz (which is large compared to the typical statistical measurement error of $100-200 \mu K$; Reese et al.
 2000: 2002)., 2000; 2002).
" Of course. there will also be a measurement error associated with the cluster mass,"," Of course, there will also be a measurement error associated with the cluster mass."
" Typically, X-rav-determuined masses have associated statistical uncertainties of about (Nevalainen et al."," Typically, X-ray-determined masses have associated statistical uncertainties of about (Nevalainen et al."
 2000: Fiuoguenuov ct al., 2000; Finoguenov et al.
 2001)., 2001).
 Therefore. based on Figure 3. it should be possible to coustraiu Ay to within about +£L00 keV cu? with current data.," Therefore, based on Figure 3, it should be possible to constrain $K_0$ to within about $\pm 100$ keV $^2$ with current data."
 Physically. the diminution in the normalization yy relation can be understood since eutropy jection ercatly reduces the pressure of the cluster ooOgas ucar the," Physically, the diminution in the normalization $y_0-M(r_{500})$ relation can be understood since entropy injection greatly reduces the pressure of the cluster gas near the"
necessarily isothermal.,necessarily isothermal.
 We solve this equation numerically. using a 4th order RungeIxutta method with the boundary condition αμ)=0. corresponding to zero torque at the inner edge.," We solve this equation numerically, using a 4th order Runge–Kutta method with the boundary condition $\mathrm{d}W/\mathrm{d}r(r_\mathrm{in})=0$, corresponding to zero torque at the inner edge."
 The amplitude of this linear. solution niv be fixed by specifving the value Woy=M(0) at the inner boundary. be. at the marginally stable orbit: this corresponds to the (small) inclination of the inner edge of the disc with respect to the equator of the black hole.," The amplitude of this linear solution may be fixed by specifying the value $W_0=W(r_\mathrm{in})$ at the inner boundary, i.e., at the marginally stable orbit; this corresponds to the (small) inclination of the inner edge of the disc with respect to the equator of the black hole."
 X typical solution is shown in Fig. 2.., A typical solution is shown in Fig. \ref{warpg}.
 Phe warp has an oscillatory »haviour. as found by Ivanov.&IHlarionov(LOOT). with he wavelength increasing with radius. consistent with the ocal dispersion relation.," The warp has an oscillatory behaviour, as found by \cite{ivanovillarionov1997}, with the wavelength increasing with radius, consistent with the local dispersion relation."
 This non-monotonic behaviour of he inclination contrasts with the BardeenPetterson elfec (Bardeen&Petterson1975).. which was derived. using an incorrect equation for the warp.," This non-monotonic behaviour of the inclination contrasts with the Bardeen–Petterson effect \citep{bardeenpetterson1975}, which was derived using an incorrect equation for the warp."
 We would normally expec Wir) to tend to a constant value at large r. corresponding o the inclination of the outer part of the disc with respec o the equator of the black hole.," We would normally expect $W(r)$ to tend to a constant value at large $r$, corresponding to the inclination of the outer part of the disc with respect to the equator of the black hole."
 Unfortunately this is no rue of the approximate equation (27)). which does no 1old. accurately at large r because the wavelength becomes comparable to the radius.," Unfortunately this is not true of the approximate equation \ref{eqg}) ), which does not hold accurately at large $r$ because the wavelength becomes comparable to the radius."
 However. since we are interestec in the interaction of the warp with waves that propagate in the inner disc. this is not expected to significantly alfec the final results.," However, since we are interested in the interaction of the warp with waves that propagate in the inner disc, this is not expected to significantly affect the final results."
 We defer to a second paper a more realistic treatment of the propagation of the warp into the inner par of the disc., We defer to a second paper a more realistic treatment of the propagation of the warp into the inner part of the disc.
 The non-linearities in the basic equations (1)) ancl (2)) oovide couplings between the dillerent. linear modes. of he system., The non-linearities in the basic equations \ref{motion1}) ) and \ref{energy1}) ) provide couplings between the different linear modes of the system.
 We are interested in those couplings that lead o amplification of the trapped. modes., We are interested in those couplings that lead to amplification of the trapped modes.
 The basic idea of he excitation mechanism in warped disces (kato2004) is that the warp interacts with a wave in the disc (the rapped r mode) giving rise to an intermediate mode., The basic idea of the excitation mechanism in warped discs \citep{katowarp2004} is that the warp interacts with a wave in the disc (the trapped r mode) giving rise to an intermediate mode.
 This intermediate mode can then couple with the warp to feed ick on the original oscillations (see Fig. 3)).," This intermediate mode can then couple with the warp to feed back on the original oscillations (see Fig. \ref{diagram}) ),"
 resulting in erowth of the latter., resulting in growth of the latter.
 For the r mode to be excited. it needs to gain energy in his coupling.," For the r mode to be excited, it needs to gain energy in this coupling."
 Since the warp has null frequency. its energy is essentially zero and so the energy exchanges only happen αἱοσα the r and the intermediate modes and the disc.," Since the warp has null frequency, its energy is essentially zero and so the energy exchanges only happen between the r and the intermediate modes and the disc."
 11 is widely agreed. and certainly true in the short-wavelength limit. although a general proof is lacking. that a mode that propagates inside its corotation radius has negative energy. Le. the total energy. of the disce is reduced in the presence of the wave. which is possible because the disc is rotating.," It is widely agreed, and certainly true in the short-wavelength limit, although a general proof is lacking, that a mode that propagates inside its corotation radius has negative energy, i.e., the total energy of the disc is reduced in the presence of the wave, which is possible because the disc is rotating."
 On the other hand. an axisvmmetric wave. such as the r mode. or one that propagates outside its corotation radius. has positive energy.," On the other hand an axisymmetric wave, such as the r mode, or one that propagates outside its corotation radius, has positive energy."
 Suppose that. through coupling with the warp. the r mode generates an intermediate wave that propagates inside its corotation radius and therefore has negative enerev.," Suppose that, through coupling with the warp, the r mode generates an intermediate wave that propagates inside its corotation radius and therefore has negative energy."
 In the process of generating this wave. the r mode gains energy ancl is amplified.," In the process of generating this wave, the r mode gains energy and is amplified."
 For sustained growth of the r mode. the intermediate wave must be damped so that its negative energy is continually replenished by the r mode. (," For sustained growth of the r mode, the intermediate wave must be damped so that its negative energy is continually replenished by the r mode. ("
Phe damping process itselt draws positive energy from the rotation of the disc.),The damping process itself draws positive energy from the rotation of the disc.)
 Therefore a dissipation term should. be included. in the equations for the intermediate mode., Therefore a dissipation term should be included in the equations for the intermediate mode.
 We choose to damp this wave locally at à rate JQ. where 3 is a dimensionless parameter.," We choose to damp this wave locally at a rate $\beta\Omega$, where $\beta$ is a dimensionless parameter."
 The origin of this term djs not discussed. here but if we interpret it as some type of viscous dissipation or [friction in the disc we expect the intermediate mode. which propagates in a larger region in the disc. to be more allected by it than the r mode. as the latter is trapped in a small region. am has a simpler racial structure.," The origin of this term is not discussed here but if we interpret it as some type of viscous dissipation or friction in the disc we expect the intermediate mode, which propagates in a larger region in the disc, to be more affected by it than the r mode, as the latter is trapped in a small region, and has a simpler radial structure."
 Also. the intermediate mode approaches its corotation radius (or the mareinally stable orbit). where it is expected to be absorbed. and this elfec is implicitly. included in the intermediate mode equations when the friction term is included.," Also, the intermediate mode approaches its corotation radius (or the marginally stable orbit), where it is expected to be absorbed, and this effect is implicitly included in the intermediate mode equations when the friction term is included."
 Vherefore. we neglec the dissipation term in the equations for the r mode.," Therefore, we neglect the dissipation term in the equations for the r mode."
 The erowth rate that we obtain for the trapped mode shoulc be compared with estimates of its damping rate due to turbulent viscosity., The growth rate that we obtain for the trapped mode should be compared with estimates of its damping rate due to turbulent viscosity.
 For the coupling to occur the waves need to propagate in the same region in the disc and the parameters a and m [or the 3 modes need to follow some basic coupling rules. where the subscripts It. Wane LI refer to r mode. warp and intermediate mode quantities. respectively.," For the coupling to occur the waves need to propagate in the same region in the disc and the parameters $\omega$ and $m$ for the 3 modes need to follow some basic coupling rules, where the subscripts R, W and I refer to r mode, warp and intermediate mode quantities, respectively."
 These rules follow from the quadratic nature of the non-linearitites in the, These rules follow from the quadratic nature of the non-linearitites in the
 7, .
220817) The huuinosities associated with these energies can be obtained by imultiphiug them by the accretion rate. with an additional factor f=1.55 for the thermal/compressional lunünositv.," The luminosities associated with these energies can be obtained by multiplying them by the accretion rate, with an additional factor $f=1.75$ for the thermal/compressional luminosity."
" The radius is typically a «105 απ, πο the accretion euergv is roughly two orders of maguitude larger than the thermal or nuclear energy available duriug the accretion phase."," The radius is typically a $\E{8}$ cm, so the accretion energy is roughly two orders of magnitude larger than the thermal or nuclear energy available during the accretion phase."
 However. the accretion Iuuinosityv from a disk is variable and ereatlv reduced in disk quiescence.," However, the accretion luminosity from a disk is variable and greatly reduced in disk quiescence."
 Thus. it is still possible to observe the luminosity produced frou the interior of the accreted cuvelope. as demonstrated by Townsley&Bildsten(2003) in thei work on relating effective temperature. τμ. to AL.," Thus, it is still possible to observe the luminosity produced from the interior of the accreted envelope, as demonstrated by \cite{tb03} in their work on relating effective temperature, $T_{\rm eff}$ , to $\dot{M}$."
" Figue 10 shows the surface bhuunmositv. ignorius the accretion huuinositv. aud effective. teniperature as a function of tie for a LOA/.. WD with Af=105A,vro ὃν X4=10 2. and metallicities of 0.1. 0.2. 0.5. 1.0. 2.0. 5.0 Z... increasing from right to left."," Figure \ref{fig:lumz} shows the surface luminosity, ignoring the accretion luminosity, and effective temperature as a function of time for a $1.0 \ M_\odot $ WD with $ \dot{M}=10^{-8} \ \smpy$ , $X_3=10^{-4}$ , and metallicities of 0.1, 0.2, 0.5, 1.0, 2.0, 5.0 $Z_\odot$, increasing from right to left."
 The 0.5 Z.. model is interesting m that there is euouch carbon to power significant unclear luminosity. but nof enough to trieecr a nova.," The 0.5 $Z_\odot$ model is interesting in that there is enough carbon to power significant nuclear luminosity, but not enough to trigger a nova."
 For this case. the compressional ΠΠ is «0.5L. and Tig«5.5104 K for ~50% of the nova evele.," For this case, the compressional luminosity is $<0.5 \ L_\odot$ and $T_{\rm eff}<5.5\E{4}$ K for $\sim50\%$ of the nova cycle."
 Four thousand vears after the onset of accretion. carbon is buried aud depleted. causing the iunuinositv from inside the euvelope to double within a span of oulv 1000 vr.," Four thousand years after the onset of accretion, carbon is burned and depleted, causing the luminosity from inside the envelope to double within a span of only 1000 yr."
 The burning of !N beeins ~1000 vr after C-depletion. aud the hunuinositv aud effective emperature vise to 1.3.Εν and 8<104 K just prior to he CNO cevele-trigeered ignition.," The burning of $^{14}$ N begins $\sim 1000$ yr after C-depletion, and the luminosity and effective temperature rise to $1.3 \ L_\odot$ and $8\E{4}$ K just prior to the CNO cycle-triggered ignition."
 Chaneine NY; also has an effect on the surface muinositv and Tuy., Changing $X_3$ also has an effect on the surface luminosity and $T_{\rm eff}$.
" Figure 11 shows the surface uunmnositw and effective temperature— as in Figure 10.. mit with fixed metallicity Z=Z. and varving ""He mass Yactions of 0.1. 0.2, 0.5. 1.0. 2.0, 5.0.10.2. increasing Youn right to left."," Figure \ref{fig:lumx3} shows the surface luminosity and effective temperature as in Figure \ref{fig:lumz}, but with fixed metallicity $Z=Z_\odot$ and varying $^3$ He mass fractions of 0.1, 0.2, 0.5, 1.0, 2.0, $5.0\E{-3}$, increasing from right to left."
 Hore. the interesting case is V3=10.? which las cuough ?Ie to produce siguificaut encre. but uot enough to trigeer the nova.," Here, the interesting case is $X_3=10^{-3}$, which has enough $^3$ He to produce significant energy, but not enough to trigger the nova."
" The surface hpnuuinosity reaches L. aud Tig rises to 6.5«10! AS after only 1500 yr of accretion. but the surface then cools as the envelope succeeds in depleting the acciunulated Πο, "," The surface luminosity reaches $L_\odot$ and $T_{\rm eff}$ rises to $6.5\E{4}$ K after only 1500 yr of accretion, but the surface then cools as the envelope succeeds in depleting the accumulated $^3$ He."
The sface brightens when C-buruiug conuueuces. aud again dinis when carbou is depleted. finally rising to 2L. aud 8<10t Is when !N beeius buruiug.," The surface brightens when C-burning commences, and again dims when carbon is depleted, finally rising to $2 \ L_\odot$ and $8\E{4}$ K when $^{14}$ N begins burning."
 Motivated bv uncertainties in classical mova core-musing mechauisuis aud the lack of metal enliancenmeuts in sole nova ejecta. we have quantified the effects of coluposition on nova ienitions (sce Fig. 5))," Motivated by uncertainties in classical nova core-mixing mechanisms and the lack of metal enhancements in some nova ejecta, we have quantified the effects of composition on nova ignitions (see Fig. \ref{fig:mignvsz}) )"
 under the asstuuption that the wnderling material is C-poor aud diffusion thus uuiiportaut. as appropriate for accretors with laree helimu buffers or O/Ne cores.," under the assumption that the underlying material is C-poor and diffusion thus unimportant, as appropriate for accretors with large helium buffers or O/Ne cores."
 We have found liat for carbou mass fractions <2«10.7. 22€ is depleted and converted to t!N without releasing enough heat o trigger a nuclear instability.," We have found that for carbon mass fractions $\lesssim 2\E{-3}$, $^{12}$ C is depleted and converted to $^{14}$ N without releasing enough heat to trigger a nuclear instability."
 The laver coutinucs to acerete until £N cau captive protons. leading to a nova rigeered by the full CNO exele and an ignition mass arecr than the carbon-ienited case.," The layer continues to accrete until $^{14}$ N can capture protons, leading to a nova triggered by the full CNO cycle and an ignition mass larger than the carbon-ignited case."
 Thus. the ignition nass dnereases by a factor of ~3 as the ietallicity is decreased from 5.0Z.. to 0.1Z..," Thus, the ignition mass increases by a factor of $\sim 3$ as the metallicity is decreased from $5.0 \ Z_\odot$ to $0.1 \ Z_\odot$."
" The critical carbon nass fraction is neur-solar and is thus relevant to sub-solar metallicity systems as well as systems with evolved secondaries that have wudergone CNO processing of °C o HN, We have also examined the effect of accreted Πο (see Fie. 6))."," The critical carbon mass fraction is near-solar and is thus relevant to sub-solar metallicity systems as well as systems with evolved secondaries that have undergone CNO processing of $^{12}$ C to $^{14}$ N. We have also examined the effect of accreted $^3$ He (see Fig. \ref{fig:mignvsx3}) ),"
 which can reach mass fractions of E4102 as an evolved donors interior is uucovered by mass transfer., which can reach mass fractions of $4\E{-3}$ as an evolved donor's interior is uncovered by mass transfer.
 For X4>24107. ?He L?IHo reactions trigger novae witl Afiey afactor of ~3 times sinaller than the C-trigecredao Case.," For $X_3 \gtrsim 2\E{-3}$, $^3$ $^3$ He reactions trigger novae with $M_{\rm ign} $ afactor of $ \sim 3 $ times smaller than the C-triggered case."
" The dependence of Mj, on accreted composition will affect population-averaged nova rates: naively. ligh- environments would have nova rates higher by a factor of ~3 than svstenis with sub-solar moetallicities (suchas novae iu globular clusters: Shafter&Quimby 2007)) or evolved donors that have undergone CNO"," The dependence of $M_{\rm ign}$ on accreted composition will affect population-averaged nova rates: naively, high-metallicity environments would have nova rates higher by a factor of $\sim 3$ than systems with sub-solar metallicities (suchas novae in globular clusters; \citealt{sq07}) ) or evolved donors that have undergone CNO"
Fe Κα emission line (Dewangan et 22004. Young et 22007).,"Fe $\alpha$ emission line (Dewangan et 2004, Young et 2007)."
 In stellar mass systems. it is not possible to test disk- connections on the period defined by the QPO. although the oscillation might be tied to jet production.," In stellar mass systems, it is not possible to test disk-jet connections on the period defined by the QPO, although the oscillation might be tied to jet production."
 In supermassive black holes. however. this timescale 1s accessible.," In supermassive black holes, however, this timescale is accessible."
 For a black hole of 7«107M. like that in M81* (Devereux et 22003). monitoring every 2-4 weeks can sample the corresponding timescale.," For a black hole of $7\times 10^{7}~{\rm M}_{\odot}$ like that in M81* (Devereux et 2003), monitoring every 2–4 weeks can sample the corresponding timescale."
 In this paper. we present contemporaneous X-ray and radio observations of M81* made using and the VLA. with visits separated by approximately 20 days.," In this paper, we present contemporaneous X-ray and radio observations of M81* made using and the VLA, with visits separated by approximately 20 days."
 The observations and data reduction methods are described in Section 2., The observations and data reduction methods are described in Section 2.
 Our analysis and results are presented in Section 3., Our analysis and results are presented in Section 3.
 We do not find a clear correlation between radio and X-ray emission in M81*. though a small number of simultaneous points were obtained and span a factor of approximately two in X-ray flux.," We do not find a clear correlation between radio and X-ray emission in M81*, though a small number of simultaneous points were obtained and span a factor of approximately two in X-ray flux."
 These results are discussed in Section 4., These results are discussed in Section 4.
 We observed M81* on ten occasions usingChandra., We observed M81* on ten occasions using.
 Each observation achieved a total exposure of approximately 15 ksec (see Table 1)., Each observation achieved a total exposure of approximately 15 ksec (see Table 1).
 In order to minimize photon pile-up in the zeroth order ACIS image. the HETGS was inserted into the light path in each case.," In order to minimize photon pile-up in the zeroth order ACIS image, the HETGS was inserted into the light path in each case."
" The ACIS chips were operated in ""FAINT"" mode.", The ACIS chips were operated in “FAINT” mode.
 We used CIAO version 4.0.2 in processing the data., We used CIAO version 4.0.2 in processing the data.
" First-order dispersed spectra from the MEG and HEG were split from the standard ""pha2"" file. and associated instrument. response files were constructed."," First-order dispersed spectra from the MEG and HEG were split from the standard “pha2” file, and associated instrument response files were constructed."
" The first-order MEG spectra and responses were then added using the CIAO tool sspectra"": the first-order HEG spectra and responses were eggrating.added in the same way.", The first-order MEG spectra and responses were then added using the CIAO tool spectra”; the first-order HEG spectra and responses were added in the same way.
" The zeroth-order ACIS spectra and responses were generated using the CIAO tool ""psextract"".", The zeroth-order ACIS spectra and responses were generated using the CIAO tool “psextract”.
 In each case. a circular region was used to extract the source flux and a radially—offset annular region was used to extract the background flux.," In each case, a circular region was used to extract the source flux and a radially–offset annular region was used to extract the background flux."
 All spectra were grouped to require at least 20 counts per bin using the FTOOL “erppha™. in order to ensure the validity of 47 statistics.," All spectra were grouped to require at least 20 counts per bin using the FTOOL “grppha”, in order to ensure the validity of $\chi^{2}$ statistics."
 The VLA also observed M81* on ten occasions (see Table 1)., The VLA also observed M81* on ten occasions (see Table 1).
 Useful data were obtained on seven occasions that coincide with the X-ray observations., Useful data were obtained on seven occasions that coincide with the X-ray observations.
 All observations were obtained at 8.4 GHz., All observations were obtained at 8.4 GHz.
" The first three coincident exposures were obtained in the “A” configuration (achieving a typical angular resolution of approximately 0.37). while the last four were obtained in the ""B configuration (achieving a typical angular resolution of approximately 17)."," The first three coincident exposures were obtained in the “A” configuration (achieving a typical angular resolution of approximately 0.3”), while the last four were obtained in the “B” configuration (achieving a typical angular resolution of approximately 1”)."
 Standard compact calibrator sources were used to calibrate phase and amplitude variations. and to set the overall amplitude scale.," Standard compact calibrator sources were used to calibrate phase and amplitude variations, and to set the overall amplitude scale."
 The average flux density measured in each exposure ts reported in Table |., The average flux density measured in each exposure is reported in Table 1.
 The X-ray spectra were analyzed using XSPEC version 12.4 (Arnaud 1996)., The X-ray spectra were analyzed using XSPEC version 12.4 (Arnaud 1996).
 Spectral fits were made in the 10.0 keV band., Spectral fits were made in the 0.5--10.0 keV band.
 All of the errors reported in this work are Ia confidence errors., All of the errors reported in this work are $1\sigma$ confidence errors.
 In calculating luminosity values. distances were assumed to be absolute. and uncertainties in luminosity were derived from the flux uncertainties.," In calculating luminosity values, distances were assumed to be absolute, and uncertainties in luminosity were derived from the flux uncertainties."
 We initially made separate fits to the zeroth-order. combined MEG. and combined HEG spectra.," We initially made separate fits to the zeroth-order, combined MEG, and combined HEG spectra."
 In all direct fits. the equivalent neutral hydrogen column density drifted towards zero. which is unphysical.," In all direct fits, the equivalent neutral hydrogen column density drifted towards zero, which is unphysical."
" A value of 4.1«10°""em is expected along this line of sight (Dickey Lockman 1990). but this value is too low to be constrained directly in the MEG spectra obtained."," A value of $4.1\times 10^{20}~{\rm cm}^{-2}$ is expected along this line of sight (Dickey Lockman 1990), but this value is too low to be constrained directly in the MEG spectra obtained."
 For consistency. then. the expected value was fixed in all fits.," For consistency, then, the expected value was fixed in all fits."
 All of the spectra were acceptably fit QUfrx1.0. where » is the number of degrees of freedom in the fit) with a simple power-law model (see Figure 1).," All of the spectra were acceptably fit $\chi^{2}/\nu \leq 1.0$, where $\nu$ is the number of degrees of freedom in the fit) with a simple power-law model (see Figure 1)."
 The spectrum of M81* is likely more complex. mostly owing to local diffuse emission (Young et 22007). but a simple power-law is an acceptable fit to the modest spectra obtained in our observations.," The spectrum of M81* is likely more complex, mostly owing to local diffuse emission (Young et 2007), but a simple power-law is an acceptable fit to the modest spectra obtained in our observations."
 The zeroth-order spectra suffer from photon pile-up. and are not robust.," The zeroth-order spectra suffer from photon pile-up, and are not robust."
 Particularly in the last two observations. where the flux is higher. the best-fit power-law photon index was found to be harder.," Particularly in the last two observations, where the flux is higher, the best-fit power-law photon index was found to be harder."
 This is consistent with multiple low-energy X-rays being detected as single high energy. photons., This is consistent with multiple low-energy X-rays being detected as single high energy photons.
 Moreover. the data/model ratio in each spectrum shows an increasing positive trend with energy.," Moreover, the data/model ratio in each spectrum shows an increasing positive trend with energy."
 The HEG spectra contain many fewer photons than the MEG spectra. and were found to be of little help in constraning the souce flux or spectral index.," The HEG spectra contain many fewer photons than the MEG spectra, and were found to be of little help in constraning the souce flux or spectral index."
 We therefore restriced our flux analysis to the combined first-order MEG spectra., We therefore restriced our flux analysis to the combined first-order MEG spectra.
 The second spectrum listed in Table, The second spectrum listed in Table
by N-body simulations can be modified by the inclusion of barvonie gas. mostlv in the inner part of a galaxy where its density is increased owing to barvonic infall. in such a manner that the density profile may be changed to an isothermal one (e.g. Barnes LOST).,"by N-body simulations can be modified by the inclusion of baryonic gas, mostly in the inner part of a galaxy where its density is increased owing to baryonic infall, in such a manner that the density profile may be changed to an isothermal one (e.g. Barnes 1987)."
 Ht is also noted that the adoption of the NEW profile does not allect our results. as the inner density. profile in concern is dominated. by the bulge and disk components.," It is also noted that the adoption of the NFW profile does not affect our results, as the inner density profile in concern is dominated by the bulge and disk components."
 A dwarl ealaxy orbiting in this galactic potential is then represented by self-eravitating particles following a Wing model. where the central density. central. velocity dispersion. ancl core radius are given as AM. . 18.1 km + ancl (56 kpe. respectively.," A dwarf galaxy orbiting in this galactic potential is then represented by self-gravitating particles following a King model, where the central density, central velocity dispersion, and core radius are given as 0.3 $M_\odot$ $^{-3}$, 18.1 km $^{-1}$, and 0.56 kpc, respectively."
 In adclition. a particle with the mass of 5.1 PAL. represcning w Con is placed at the center of the galaxy.," In addition, a particle with the mass of $5 \times 10^6$ $_\odot$ representing $\omega$ Cen is placed at the center of the galaxy."
 ‘This setting vields the total mass of the system as Adj.)=5.79107 M.., This setting yields the total mass of the system as $M_{tot}=5.79 \times 10^8$ $_\odot$.
 On the other hand. the total mass of stars in this ivpothetical galaxy is estimated from the mean metallicity of stars in & Cen (CFe/l]p ~ 1.6). combined with the metallicity-Iuminosity relation for the Local Group dwarls (Cótté et al.," On the other hand, the total mass of stars in this hypothetical galaxy is estimated from the mean metallicity of stars in $\omega$ Cen $\langle$ $\rangle\sim-1.6$ ), combined with the metallicity-luminosity relation for the Local Group dwarfs (Côtté et al."
 200€) and the mass-to-light ratio (assuming AL~ 4). vielding ωρες~10)J ," 2000) and the mass-to-light ratio (assuming $M/L \sim 4$ ), yielding $M_{stars}\sim 10^7$ $_\odot$."
‘Thus. the mass of he simulated. galaxy is largely dominated by clark matter.," Thus, the mass of the simulated galaxy is largely dominated by dark matter."
 The model galaxy is represented. by a collection of 101 yarticles ancl the sell-eravity is calculated in terms of a multiple expansion of the internal potential to fourth order (White 1983: Zaritsky White LOSS)., The model galaxy is represented by a collection of $10^4$ particles and the self-gravity is calculated in terms of a multiple expansion of the internal potential to fourth order (White 1983; Zaritsky White 1988).
 This modeldwarf is disrupted by Galactic tides in the course of its orbital motion. whereas its dense core isexpected o survive and follow w Cens orbit.," This model dwarf is disrupted by Galactic tides in the course of its orbital motion, whereas its dense core is expected to survive and follow $\omega$ Cen's orbit."
 We examine two representatiyeofwit orbits for the progenitor. models and 2: model 1 follows he current orbit of: Cen. whereas 1for model 2. we calculate an orbit back to the past over ~2 Car from its current »osition and velocity taking into account dvnamical friction and then set a progenitor galaxy on its non-decaving orbit.," We examine two representative orbits for the progenitor, models 1 and 2: model 1 follows the current orbit of $\omega$ Cen, whereas for model 2, we calculate an orbit back to the past over $\sim 2$ Gyr from its current position and velocity taking into account dynamical friction and then set a progenitor galaxy on its non-decaying orbit."
 These two models provide us with satisfactory information on the generic properties of a tically disrupted progenitor and we postulate that the realistic nature of their debris is midway between these model predictions., These two models provide us with satisfactory information on the generic properties of a tidally disrupted progenitor and we postulate that the realistic nature of their debris is midway between these model predictions.
 We w Con's orbit. based on the current distance 2=5.3+0.5 MNnikpe from the Sun. position (4.6)=(3097.15). proper (yt.cos0.ps)m(DOSEO35.3.57ο mas +. and heliocentric racial velocity ey...=232.540.7 kms (Dinescu. Cirard. van Altena 1999. DCA).," We calculate $\omega$ Cen's orbit, based on the current distance $D=5.3 \pm 0.5$ kpc from the Sun, position $(l,b)=(309^\circ,15^\circ)$, proper motion $(\mu_\alpha \cos\delta,\mu_\delta) =(-5.08\pm0.35, -3.57\pm0.34)$ mas $^{-1}$, and heliocentric radial velocity $v_{los} = 232.5 \pm 0.7$ km $^{-1}$ (Dinescu, Girard, van Altena 1999, DGvA)."
 This orbit for mock1 lis characterized by frequent clisk crossings with a period of tan=OS.10 ve. retrograde motion. and apo ancl pericentric distances (rape.μον)=(6.4.1.1) kpe.," This orbit for model 1 is characterized by frequent disk crossings with a period of $\tau_{orb} = 0.8 \times 10^8$ yr, retrograde motion, and apo and pericentric distances $(r_{apo},r_{peri}) = (6.4, 1.1)$ kpc."
" For model 2. weobtain 7,4=1.5«105 vrand (riisri)=(11.3.3.0) kpe."," For model 2, we obtain $\tau_{orb} = 1.5 \times 10^8$ yr and $(r_{apo},r_{peri}) = (11.3, 3.0)$ kpc."
 1n both experiments. we place a progenitore ogalaxy at apocenter to maximize its survival chances.," In both experiments, we place a progenitor galaxy at apocenter to maximize its survival chances."
 We plot the spatial distribution of the tidallv disrupted debris in Figure 1., We plot the spatial distribution of the tidally disrupted debris in Figure 1.
 Upper(middle) panel shows model 1 (mocdel 2) after the 1.37 (1.86) vr orbital excursion o£ 106 progenitor galaxy., Upper (middle) panel shows model 1 (model 2) after the 1.37 (1.86) Gyr orbital excursion of the progenitor galaxy.
 Lower panel shows the orbit of the ealaxy center., Lower panel shows the orbit of the galaxy center.
 A rosette-like feature of the debris becomes steady after about eight. orbital periods., A rosette-like feature of the debris becomes steady after about eight orbital periods.
 Model 1 results in more compact distribution than mocel 2. which rellects 1e dilference in orbital radii.," Model 1 results in more compact distribution than model 2, which reflects the difference in orbital radii."
 Figure 2 shows the velocity listributions of the debris particles in exlindrical coordinates Ppl.02).," Figure 2 shows the velocity distributions of the debris particles in cylindrical coordinates $(v_R,v_\phi,v_z)$."
" As is evident. both mocels provide essentially re same debris kinematics: most. notorious is a sharply peaked ος, distribution⋠⋠⋠ at. —O kni s arising Pfrom a retrograde orbit of a progenitor."," As is evident, both models provide essentially the same debris kinematics: most notorious is a sharply peaked $v_\phi$ distribution at $\sim -100$ km $^{-1}$, arising from a retrograde orbit of a progenitor."
 These kinematics suggest iu the dilference in model 1l and 2 resides only in the spatial extent of the debris., These kinematics suggest that the difference in model 1 and 2 resides only in the spatial extent of the debris.
 We analyze the kinematics of both the simulated debris ancl other Galactic stars generated randomly by a Monte Carlo method., We analyze the kinematics of both the simulated debris and other Galactic stars generated randomly by a Monte Carlo method.
" The metal-poor halo is modeled as a Ilattened spheroidl p55-xCRF|τς}2BEDin. where q is an axis ratio ranging).0.50.7. anisotropic velocity ellipsoid (ay...)=(154.121.96 )km sL. and small mean rotation. ελ,= sms 1 as found for halo stars with η]2 hear he Sun (CD)."," The metal-poor halo is modeled as a flattened spheroid $\rho \propto
(R^2 + z^2/q^2)^{-3.5/2}$, where $q$ is an axis ratio ranging 0.55-0.7, anisotropic velocity ellipsoid $(\sigma_R,\sigma_\phi,\sigma_z)=(154,121,96)$ km $^{-1}$, and small mean rotation $\langle v_\phi \rangle= 24$ km $^{-1}$ , as found for halo stars with $<-2$ near the Sun (CB)."
 Thin and thick disks are .Jed as exptBER)sect(2far) (thick)where Ry=3.53. diskskpe and τω3 (1) for thin{ disk.," Thin and thick disks are modeled as $\rho \propto
\exp(-R/R_d) \sec^2 (z/z_d)$ , where $R_d=3.5$ kpc and $z_d=0.3$ (1) kpc for thin (thick) disk."
 Both rotate at 200 uns (46.50.Kyehaving velocity ellipsoids thickof{(34225.20) km and .35) )kms + for thin and disks. respectively (CD).," Both disks rotate at 200 km $^{-1}$, having velocity ellipsoids of $(34,25,20)$ km $^{-1}$ and $(46,50,35)$ km $^{-1}$ for thin and thick disks, respectively (CB)."
 The relative fraction of each component is fixed using observed local densities near the Sun. in such a manner that he halo anc thick-disk densities at D«1 kpe are 0.2% and 2% of the thin-disk density. respectively (Yamagata Yoshii 1992)," The relative fraction of each component is fixed using observed local densities near the Sun, in such a manner that the halo and thick-disk densities at $D<1$ kpc are 0.2 and 2 of the thin-disk density, respectively (Yamagata Yoshii 1992)."
 In our model of ω Cen’s progenitor galaxy. the self- particles represent both stars and dark matter particles.," In our model of $\omega$ Cen's progenitor galaxy, the self-gravitating particles represent both stars and dark matter particles."
 We note that the correct conversion ofthe mass of the simulated svstem into that of the presumed stellar, We note that the correct conversion ofthe mass of the simulated system into that of the presumed stellar
ietalbrich white dwarfs.,metal-rich white dwarfs.
 All accretion rates were re- following Farilictal.2009 using recent results on inctal diffusion timescales and convective envelope masses for both DAZ and DBZ white cwarfs. including nüxed cases (IKoesterx20092:: D. Ἱνουνίο 2009. private conumaiunication): parameters for the newly observed Cycle 5 stars are listed in Table &..," All accretion rates were re-calculated following \citealt{far09a} using recent results on metal diffusion timescales and convective envelope masses for both DAZ and DBZ white dwarfs, including mixed cases \citealt{koe09a}; D. Koester 2009, private communication); parameters for the newly observed Cycle 5 stars are listed in Table \ref{tbl8}."
 Plotted as newly discovered disks are 3253. | 0716. 1630. aud tentatively 079.," Plotted as newly discovered disks are $-$ 3253, $+$ 0746, $-$ 1630, and tentatively $-$ 079."
 To these are added three white cowarfs orbited by both dust and gaseous metals (Brinkworthetal. 2009:: ο Brinkworth 2008. private comuumnication). while 16321177 (observed by Emihietal.2008 and listed as DAZ in Farilietal. 2009)) has been removed from this erowing sanuple. as it is not metallined (Zuckerman 2003).," To these are added three white dwarfs orbited by both dust and gaseous metals \citealt{bri09}; ; C. Brinkworth 2008, private communication), while $+$ 177 (observed by \citealt{far08} and listed as DAZ in \citealt{far09a}) ) has been removed from this growing sample, as it is not metal-lined \citep{zuc03}."
. The new additions to the plots continue the treud that mctal accretion rates dAL/dt23<10S ges+ are those that are most Likely to exhibit infrared excess (Farilietal. 2009):: the sole exception is G166-58.," The new additions to the plots continue the trend that metal accretion rates $dM/dt \ga 3\times 
10^8$ $^{-1}$ are those that are most likely to exhibit infrared excess \citep{far09a}; the sole exception is G166-58."
 Two of the disks discovered in this work orbit relatively cool white dwarts with cooling ages of 0.73 and 0.9Cyr., Two of the disks discovered in this work orbit relatively cool white dwarfs with cooling ages of 0.73 and Gyr.
 These are solmewhat older than the previous averageC cooling age of white dwarts with disks as found in(2009).. aud may iudicate relative longevity for orbiting dust.," These are somewhat older than the previous average cooling age of white dwarfs with disks as found in, and may indicate relative longevity for orbiting dust."
 Viscous spreading among solids may allow optically thick dust vines at white dwarfs to persist for at least MM. timescales. possibly orders of magnitude longer. analogous to plauctary rings.," Viscous spreading among solids may allow optically thick dust rings at white dwarfs to persist for at least Myr timescales, possibly orders of magnitude longer, analogous to planetary rings."
 The inner radius of Satis rings represcuts the furthest iuner extent of its spreading over a period up to a few Car (Espositoctal.2008)., The inner radius of Saturn's rings represents the furthest inner extent of its spreading over a period up to a few Gyr \citep{esp08}.
. For dust riugs at white dwarfs. eas drag should decrease the viscous timescale by orders of magnitudo. but the gas coutent im ost white dwarf disks is presently uncoustrained.," For dust rings at white dwarfs, gas drag should decrease the viscous timescale by orders of magnitude, but the gas content in most white dwarf disks is presently unconstrained."
 At the inner edee of a typical dust disk. rapid sublimation of dust will produce significant eas; as nay erain-erain collisions.," At the inner edge of a typical dust disk, rapid sublimation of dust will produce significant gas, as may grain-grain collisions."
 Such viscosity euliancenmenuts are necessary fo reproduce the dust acerction rates interred for stars with short photospheric metal lifetimes., Such viscosity enhancements are necessary to reproduce the dust accretion rates inferred for stars with short photospheric metal lifetimes.
4. Dust rings at white dwarts are unlikely to be as simple as mmocdeled here: cale triplet cutission lies at three motal-polluted stars can be approximately reproduced by orbiting disks. but there are asvuuuctries which indicate deviatious from the mocel (Causickeetal.2008).," Dust rings at white dwarfs are unlikely to be as simple as modeled here; calcium triplet emission lines at three metal-polluted stars can be approximately reproduced by orbiting disks, but there are asymmetries which indicate deviations from the model \citep{gan08}."
.. Yet it is noteworthy that the simple model works rather well eiven ouly three free parameters. two of which the ner and outer radius — are plysically well-inotivated.," Yet it is noteworthy that the simple model works rather well given only three free parameters, two of which – the inner and outer radius – are physically well-motivated."
 Significant orbital cnerev must be shed to transform. an asteroid on a perturbed eccentric orbit into a disk contained eutirelv within the Roche huit of a white dwarf.," Significant orbital energy must be shed to transform an asteroid on a perturbed, eccentric orbit into a disk contained entirely within the Roche limit of a white dwarf."
 The primordial eucrgv iav dissipate partially via collisions which initially widen the disk (or perhaps uot iu the case of narrow ries). vet an cfiicicut wav to discard orbital energy is within a large fragment of the tidally disrupted parent body.," The primordial energy may dissipate partially via collisions which initially widen the disk (or perhaps not in the case of narrow rings), yet an efficient way to discard orbital energy is within a large fragment of the tidally disrupted parent body."
 Surviving fraguienuts iav be responsible for warps invoked iu some disk models. such as 232362 and 556 (Juraetal.2009a.2007b).," Surviving fragments may be responsible for warps invoked in some disk models, such as 362 and 56 \citep{jur09a,jur07b}."
. Satuusrngduoon Daphuis — 8kkum in diuneter auk conrparable to a small asteroid causes lkkn vertical disturbances in the A ring: LOO —times lavecr than the height of the uuperturbed rugs (Weissetal.2009)., Saturn's ring-moon Daphnis – km in diameter and comparable to a small asteroid – causes km vertical disturbances in the A ring; 100 times larger than the height of the unperturbed rings \citep{wei09}.
. The more substantial Mimas. line just outside Saturu's main vines (and its Roche lint) with a iiass nearly identical to the total mass ofthe vines. is responsiblefor the Cassini Division between the A and B vines.," The more substantial Mimas, lying just outside Saturn's main rings (and its Roche limit) with a mass nearly identical to the total mass of the rings, is responsiblefor the Cassini Division between the A and B rings."
 A Mimas-analoe, A Mimas-analog
Nordini et al. (,Nardini et al. (
2009). and Zheng et al. (,"2009), and Zheng et al. ("
2009) [rom the data analysis.,2009) from the data analysis.
 The quantities of neutral gas in the host galaxies can be obtained [rom (he optical spectra., The quantities of neutral gas in the host galaxies can be obtained from the optical spectra.
 By (he measurement of Ίωνα absorption. Fvnbo et al. (," By the measurement of $\alpha$ absorption, Fynbo et al. ("
"2009) have established one sample in which 33 values of neutral hydrogen cohunn density 2, ave derived.","2009) have established one sample in which 33 values of neutral hydrogen column density $N_{H,opt}$ are derived."
 The range of these values is [rom LOMem7 to LOem2 (see Fig., The range of these values is from $10^{17}cm^{-2}$ to $10^{23}cm^{-2}$ (see Fig.
 10 of Fynbo et al., 10 of Fynbo et al.
" 2009). while the tyne distribution of [N,,,4,; may extend to the higher column densities."," 2009), while the true distribution of $N_{H,opt}$ may extend to the higher column densities."
 On the other hand. the damped Lya svstem with the neutral hydrogen number exceeding 2x10?emi/7 has the possibility of star formation to form a protogalaxy (see the simulations by Pontzen et al.," On the other hand, the damped $\alpha$ system with the neutral hydrogen number exceeding $2\times 10^{20}cm^{-2}$ has the possibility of star formation to form a protogalaxy (see the simulations by Pontzen et al."
 2009 recently): thus. it could be the GRD host.," 2009 recently); thus, it could be the GRB host."
" But the thin cloud with the smaller neutral hvdrogen values might be intervening along the line of sight between the observer and the GRB place: this kind of thin cloud with the column densitv less than ~10?""ej7 might not be related to the GRB host.", But the thin cloud with the smaller neutral hydrogen values might be intervening along the line of sight between the observer and the GRB place; this kind of thin cloud with the column density less than $\sim 10^{20}cm^{-2}$ might not be related to the GRB host.
 Also. in this paper. we assume that GRD hosts are rich in neutral gas.," Also, in this paper, we assume that GRB hosts are rich in neutral gas."
" Therefore. we only select the Nyy.) values larger than L0""""em? and compare them with the corresponding X-ray absorption Nyy, values."," Therefore, we only select the $N_{H,opt}$ values larger than $10^{20}cm^{-2}$ and compare them with the corresponding X-ray absorption $N_{H,x}$ values."
 We find the relation between X-ray absorption and optical neutral gas shown in Fig. 9:, We find the relation between X-ray absorption and optical neutral gas shown in Fig. \ref{f6}:
" /og.N4,,=(0.492:0.04)/09.N2 0.9)."," $log N_{H,x}=(0.49\pm0.04)log N_{H,opt}+(11.3\pm 0.9)$ ."
 The linear correlation coefficient is 0.58 with the probability 0.001., The linear correlation coefficient is 0.58 with the probability 0.001.
 Through this weak relationship. it is Likely to find the trace of possible cosmic evolution of neutral gas μοι similar to the evolution of X-ray Ny; in. Fig. 8..," Through this weak relationship, it is likely to find the trace of possible cosmic evolution of neutral gas $N_{H,opt}$, similar to the evolution of X-ray $N_{H,x}$ in Fig. \ref{f5}."
" Suppose that GRBs are the unbiased tracers of star formation at high redshift. (his possible Niro), distribution mav eive an interpretation to the observations of IHE gas evolution by Prochaska Wolle (2009)."," Suppose that GRBs are the unbiased tracers of star formation at high redshift, this possible $N_{H,opt}$ distribution may give an interpretation to the observations of HI gas evolution by Prochaska Wolfe (2009)."
 llowever. we caution (hat the relation between A-rayv absorption ancl optical neutral gas may have larger uncertainties. due to the limited redshilt range [rom 2 to 3.," However, we caution that the relation between X-ray absorption and optical neutral gas may have larger uncertainties, due to the limited redshift range from 2 to 3."
J) A complete sample is requirecl to investigate (his relation in the future., A complete sample is required to investigate this relation in the future.
 From our model. we see that dust absorption ον is the function of SFR and metallicity.," From our model, we see that dust absorption $A_v$ is the function of SFR and metallicity."
 Combining the effects of both SER and metallicity. we obtain the redshift distribution of dust absorption shown in Fig. 10..," Combining the effects of both SFR and metallicity, we obtain the redshift distribution of dust absorption shown in Fig. \ref{f7}."
 We see that the Av values from our model are slightly increasing wilh redshift., We see that the Av values from our model are slightly increasing with redshift.
 From the data of Schady et al. (, From the data of Schady et al. (
2010). we do not find any prominent evidence ol d. variation.,"2010), we do not find any prominent evidence of $A_v$ variation."
 But it was claimed by. IXann et al. (, But it was claimed by Kann et al. (
2007) that the sl. value decreases with increasing redshilt.,2007) that the $A_v$ value decreases with increasing redshift.
 ILere. after the comparison between the two data sets given by Schack et al. (," Here, after the comparison between the two data sets given by Schady et al. ("
2010) and Wann et al. (,2010) and Kann et al. (
"2007). we find that the two cata sets are not consistent with each other: some ;1, values of the same burst have large difference.","2007), we find that the two data sets are not consistent with each other: some $A_v$ values of the same burst have large difference."
 The details are listed in the caption to Fig. 10.., The details are listed in the caption to Fig. \ref{f7}.
 At high redshift. the values of Kann et al. (," At high redshift, the values of Kann et al. ("
2007) are lower than those of Schady et al. (,2007) are lower than those of Schady et al. (
2010). while at low recdshift. the values of Kann et al. (,"2010), while at low redshift, the values of Kann et al. ("
2007) are larger than those of Schady et al. (,2007) are larger than those of Schady et al. (
2010).,2010).
 In general. low-mass stars take a long time to evolve to asvinplolic giant branch (AGB) phase and to produce the dust: thus. AGB population only dominate (he dust production at local universe.," In general, low-mass stars take a long time to evolve to asymptotic giant branch (AGB) phase and to produce the dust; thus, AGB population only dominate the dust production at local universe."
 It is suggested (hat at high-redshift the dust factory is supernova explosion., It is suggested that at high-redshift the dust factory is supernova explosion.
 However. recently. AGB population has been found to be the," However, recently, AGB population has been found to be the"
models of 2 and. ?..,models of \citet{Lagache04} and \citet{Chary01}.
 This justifies our use of the Schechter and power-law functional forms as our ‘best’ and ‘maximal’ estimates. respectively.," This justifies our use of the Schechter and power-law functional forms as our `best' and `maximal' estimates, respectively."
 The lensecl source count estimates in this paper at high Dux densities should be considered to be known local (e.g. LRALS)) galaxies. unless otherwise stated.," The lensed source count estimates in this paper at high flux densities should be considered to be known local (e.g., ) galaxies, unless otherwise stated."
 We will see that. as expected. highly amplified sources make an increasing contribution at brighter I[uxes.," We will see that, as expected, highly amplified sources make an increasing contribution at brighter fluxes."
 If an object with flux density S is amplified through lensing by a factor pi. the object is observed with Dux density (2).," If an object with flux density $S$ is amplified through lensing by a factor $\mu$, the object is observed with flux density $\Sobs = \mu S$ \citep{Schneider92}."
 For à given observed [ux Sa. we run through a range of magnifications. find the corresponding S. and assign pGONcls to be the contribution to the lensed differential source counts from that ye.," For a given observed flux $\Sobs$, we run through a range of magnifications, find the corresponding $S$, and assign $p(\mu) \d N/\d S$ to be the contribution to the lensed differential source counts from that $\mu$."
 In this way we can express the cdilferential source counts after lensing as For different values of Sin. we plot the contribution to the above integral as a function of jr in Fig. 3.., In this way we can express the differential source counts after lensing as For different values of $\Sobs$ we plot the contribution to the above integral as a function of $\mu$ in Fig. \ref{fig:z2-amp-combo}.
 We do this specifically for sources at >=2. as well as showing a separate panel with results normalised by the total contribution fron all ji," We do this specifically for sources at $z=2$, as well as showing a separate panel with results normalised by the total contribution from all $\mu$."
 In the normalised case. the line can be interpreted as a local p(y) for that Soi.," In the normalised case, the line can be interpreted as a local $p(\mu)$ for that $\Sobs$."
 Although we have onlv. plotted results for z=2. other redshift. choices look qualitatively similar.," Although we have only plotted results for $z=2$, other redshift choices look qualitatively similar."
 We can also examine the contribution to the source counts coming from various magnifications by specifving a minimum fr that we might be interested in., We can also examine the contribution to the source counts coming from various magnifications by specifying a minimum $\mu$ that we might be interested in.
 The fractional contribution to cLΑαδν From feμμ is found simply by changing the lower limit on the integral in equation 9. from Ü to fis. and normalising by the full dNdS.," The fractional contribution to $\d N/\d\Sobs$ from $\mu>\mu_\rmn{min}$ is found simply by changing the lower limit on the integral in equation \ref{eqn:dNdSobs} from $0$ to $\mu_\rmn{min}$, and normalising by the full $\d N/\d\Sobs$."
 This is plotted in Fig. 4.., This is plotted in Fig. \ref{fig:z2-mmu}.
 It can be seen that at high flux densities a large portion of sources will have significant aniplifications., It can be seen that at high flux densities a large portion of sources will have significant amplifications.
 Lensing PDFs and α Να function can be combined at arbitrary redshifts in the range 0.5<z«5.7 (where ? quantify their lensing simulation).," Lensing PDFs and a $\d N/\d S$ function can be combined at arbitrary redshifts in the range $0.5 < z < 5.7$ (where \citet{Hilbert07}
 quantify their lensing simulation)."
 The cdilferential source counts of the lensecl populations are found: via equation 9 and shown in Fig., The differential source counts of the lensed populations are found via equation \ref{eqn:dNdSobs} and shown in Fig.
 5. [or a range of source redshifts., \ref{fig:num-combo} for a range of source redshifts.
 One can, One can
a photometric survey of stars in the Wepler field-ot-,a photometric survey of stars in the Kepler field-of-view.
" However. Brownctal.(2011) state that the KIC stellar parameters are reliable for Sun-like stars. mt are ""untrustworthnv for stars with τω less than 3750 Iv."," However, \citet{Brown2011} state that the KIC stellar parameters are reliable for Sun-like stars, but are “untrustworthy” for stars with $T_{\rm eff}$ less than 3750 $K$."
 The most reliable estiamtes of M. dwarf masses iid radii are derived by combining measured stellarπιοΊος with reliable mass-huninosity rolatious (e.g.Delfosseetal.2000). amd mass-racius relations predicted wi stellar evolutionary models (c.g.Daraffeetal.1998).. often with corrections fo account for discrepancies vetween ameasured ancl predicted radi (c.g.Torres 2007).," The most reliable estiamtes of M dwarf masses and radii are derived by combining measured stellarluminosities with reliable mass-luminosity relations \citep[e.g.][]{Delfosse2000} and mass-radius relations predicted by stellar evolutionary models \citep[e.g.][]{Baraffe1998}, often with corrections to account for discrepancies between measured and predicted radii \citep[e.g.][]{Torres2007}."
. Unfortuuatelv. the low-mass RKOIs do uot have ο measurements which are necessary to estimate stellar luminosity aud hence mass aud radius with these uethods.," Unfortunately, the low-mass KOIs do not have parallax measurements, which are necessary to estimate stellar luminosity and hence mass and radius with these methods."
 Near-infrared spectroscopy offers ἃ more robust uethod for deteriuniug physical paramcters for low-nass stars that lack parallax measurements., Near-infrared spectroscopy offers a more robust method for determining physical parameters for low-mass stars that lack parallax measurements.
 The A- (2.0 to 2.1 qun) contains several useful spectral diagnostics includiug coutimmim regions sensitive to 1η (Coveyetal.2010). aud absorption features that are sensitive to stellar metallicity (Rojas-Avalaetal.2010. 2012).. with minimal scusitivity to interstellar reddening.," The $K$ -band (2.0 to 2.4 $\mu m$ ) contains several useful spectral diagnostics including continuum regions sensitive to $T_{\rm eff}$ \citep{Covey2010} and absorption features that are sensitive to stellar metallicity \citep{Rojas2010, Rojas2012}, , with minimal sensitivity to interstellar reddening."
 Iu this letter we report Tig and |M/TI] acasuvements of Sl low-mass KOISs using A-baud spectroscopy., In this letter we report $T_{\rm eff}$ and [M/H] measurements of 84 low-mass KOIs using $K$ -band spectroscopy.
 We interpolated Tig aud [M/II| outo the Dartinouth evolutionary isochrones (Dotteretal.2008:Feiden 2011).. which reproduce measurements from optical long-baseline interferometry COLDI) relatively well (see Figure 2)) aud contain a large spread of metallicity erid-points. as required for reliable interpolation of stellar piranueters;," We interpolated $T_{\rm eff}$ and [M/H] onto the Dartmouth evolutionary isochrones \citep{Dotter2008, Feiden2011}, which reproduce measurements from optical long-baseline interferometry (OLBI) relatively well (see Figure \ref{fig:models}) ) and contain a large spread of metallicity grid-points, as required for reliable interpolation of stellar parameters."
 We report interpolated stellar masses aud radii of the low-ass IKOlIs. and recaleulate the planetary paralcters based on the transit parameters in Dorucki (2011).," We report interpolated stellar masses and radii of the low-mass KOIs, and recalculate the planetary parameters based on the transit parameters in \citet{Borucki2011}."
. Observations were carricd out with the TripleSpec Spectrograph at the Palomar Observatory 200-inch Tale Telescope., Observations were carried out with the TripleSpec Spectrograph at the Palomar Observatory 200-inch Hale Telescope.
 TripleSpec is a near-intrared slit-spectrograph covering 1.0 to 2.5 yan simultancously with a resolution of 2700 (IIlerteretal.2008)., TripleSpec is a near-infrared slit-spectrograph covering 1.0 to 2.5 $\mu m$ simultaneously with a resolution of 2700 \citep{Herter2008}.
. Two positions on the slit. A and D. were used for cach target. and exposures were taken in an ABBA pattern with 60 second integration times at cach position.," Two positions on the slit, A and B, were used for each target, and exposures were taken in an ABBA pattern with 60 second integration times at each position."
 Multiple ABBA sets were taken and combined uutil each spectra had a media channel signal-to-noise of at least 60., Multiple ABBA sets were taken and combined until each spectrum had a median per-channel signal-to-noise of at least 60.
 For tellurie calibration we used SIMDAD to ideutifv a erid of AOV stars in the [Kepler field-ofview. aud developed an observing sequence such that each IKOI observation has a correspouding AOV star observation taken within LO minutes and with au airmass difference less than 0.1.," For telluric calibration we used SIMBAD to identify a grid of A0V stars in the Kepler field-of-view, and developed an observing sequence such that each KOI observation has a corresponding A0V star observation taken within 40 minutes and with an airmass difference less than 0.1."
 We observed all of the KOTs with IKIC-ascribed effective temperatures less than 1100. A. over 7 nights iu June of 2011., We observed all of the KOIs with KIC-ascribed effective temperatures less than 4400 $K$ over 7 nights in June of 2011.
 Of the 7I IKOIs. | appeared to be double objects in the TipleSpec slit viewer with separations of less than 6 areseconds (roughly the size of the Kepler Spacecraft’s point spread function): NOT 326. KOT 611. KOI 219 and WOT 51.," Of the 74 KOIs, 4 appeared to be double objects in the TripleSpec slit viewer with separations of less than 6 arcseconds (roughly the size of the Kepler Spacecraft's point spread function): KOI 326, KOI 641, KOI 249 and KOI 51."
 IKOI 667 consisted of 5 objects within a 6 arcsecond radius., KOI 667 consisted of 5 objects within a 6 arcsecond radius.
 These objects are not included in this survey., These objects are not included in this survey.
 Iu August of 2011 we observed 15 additional NOIs with RKIC-ascribed Tig ereater than 100. A. but with colors indicative of low-mass stars: (7.I)>OT. Gei)20.3. or (gor)>LO.," In August of 2011 we observed 15 additional KOIs with KIC-ascribed $T_{\rm eff}$ greater than 4400 $K$, but with colors indicative of low-mass stars: $(J-K)>0.7$, $(r-i)>0.3$, or $(g-r)> 1.0$."
 We include 13 of these IKOIs in this letter. as their spectra revealed CO features indicative of low-mass stars.," We include 13 of these KOIs in this letter, as their spectra revealed CO features indicative of low-mass stars."
 In total. we obtained spectra of 82 IKOIs.," In total, we obtained spectra of 82 KOIs."
 The spectranal wereysis extracted using a version of the Spextool program modified for the Palomar TripleSpec Spectrograph (Cushingetal.20014.:ALCushing.pri-vatecommmluication2011)..," The spectra were extracted using a version of the Spextool program modified for the Palomar TripleSpec Spectrograph \citep[][; M. Cushing, private communication 2011]{Cushing2004}."
 Spextool accepts data in ABBA format., Spextool accepts data in ABBA format.
 The package within Spextool accepts spectra of AQ stars and compares them to a model spectruii of Vega to identify and remove telluric absorption lues in a target spectrum (Vaccaoetal.2003)., The package within Spextool accepts spectra of A0 stars and compares them to a model spectrum of Vega to identify and remove telluric absorption lines in a target spectrum \citep{Vacca2003}.
 Figure 1 plots three example spectra. with templates of simular spectra type aud relevant spectral features mdicated.," Figure \ref{spectra_plot} plots three example spectra, with templates of similar spectra type and relevant spectral features indicated."
 The templates are taken from the IRTF Spectral Library (Cushingetal.2005:Ravner20049).," The templates are taken from the IRTF Spectral Library \citep{Cushing2005, Rayner2009}."
. One star in our sample has a K-band spectru consistent with a eiut star. sueeestiug that the observed helt curve is due to a stellar. rather than plauetary. colmpanion. or that the transit signal is due to au unresolved bleud with au eclipsing binary.," One star in our sample has a K-band spectrum consistent with a giant star, suggesting that the observed light curve is due to a stellar, rather than planetary, companion, or that the transit signal is due to an unresolved blend with an eclipsing binary."
 KOI 977 shows weal Na I aud Ca I absorption aud strong CO absorption. which qualitatively match IRTF template spectra of οσα stars but not dwarf stars of the same spectral type.," KOI 977 shows weak Na I and Ca I absorption and strong CO absorption, which qualitatively match IRTF template spectra of giant stars but not dwarf stars of the same spectral type."
 The spectrum is included in Figure L.. with a eiut aud dwarf template for comparison.," The spectrum is included in Figure \ref{spectra_plot}, with a giant and dwarf template for comparison."
 To measure Tig and [M/TI] of the remaining dwarfs. we measured three spectral indices frou the I-baud spectra: the equivalent widths of the Na I and Ca I lines. at 2.210 aud 2.260 san respectively. aud an index describing the change in flux between three 0.02 jan-wide bauds dominated bv water opacityceutered at 2.215. 2.370. 2.080 pan called the IT20-K2 index.," To measure $T_{\rm eff}$ and [M/H] of the remaining dwarfs, we measured three spectral indices from the K-band spectra: the equivalent widths of the Na I and Ca I lines, at 2.210 and 2.260 $\rm \mu m$ respectively, and an index describing the change in flux between three 0.02 $\rm \mu m$ -wide bands dominated by water opacity–centered at 2.245, 2.370, 2.080 $\rm \mu m$ –called the $\rm H_2O$ -K2 index."
 Rojas-Avalaet.al.(2012) describe the measurement of the Na I aud Ca I equivalent widths. introduce H5O-R2 iudex. aud derive relations between the spectral indices aud Digg. overall unctallicity C([NLTE[) and IKITM spectral type.," \citet{Rojas2012} describe the measurement of the Na I and Ca I equivalent widths, introduce $\rm H_2O$ -K2 index, and derive relations between the spectral indices and $T_{\rm eff}$, overall metallicity ([M/H]) and KHM spectral type."
 Dreflv. we uote that the Rojas-Avalaetal.(2012) ΛΙ relation was calibrated eimpiricallvy using nearby AI dwartsΠΙ with EF. € or K-type binary companions that have SPOCS |M/TI] measurements (Valenti&Fischer 20053.," Briefly, we note that the \citet{Rojas2012} [M/H] relation was calibrated empirically using nearby M dwarfs with F, G or K-type binary companions that have SPOCS [M/H] measurements \citep{Valenti2005}."
. Moetallicitv measurements for stars earlier than MO (Fgzo04000 I) represent au extrapolation of the AL dwarf [MIT] calibration., Metallicity measurements for stars earlier than M0 $T_{\rm eff} \gtrsim 4000$ K) represent an extrapolation of the M dwarf [M/H] calibration.
 The Tig is calculated by interpolating the[MIT] measurement aud H;O-K2 iudex onto a theoretical surface of Tig versus |M/TI] aud IT4O- index calculated to the D'T-settl. late-type model spectra of Allardetal. (2011). , The $T_{\rm eff}$ is calculated by interpolating the[M/H] measurement and $\rm H_2O$ -K2 index onto a theoretical surface of $T_{\rm eff}$ versus [M/H] and $\rm H_2O$ -K2 index calculated to the BT-settl late-type model spectra of \citet{Allard2011}. .
To validate the Tig mcasurcment method. we compare Z;g measurements by Rojas-Avalaetal.(2012) to measurements from optical long-baseliue interferometry in Figure 3.. Paucl A. The Tig surface is verv uuctallicity iuseusitive (< 10 offsets cue to metallicity effects) for 3200<Tig 3900.," To validate the $T_{\rm eff}$ measurement method, we compare $T_{\rm eff}$ measurements by \citet{Rojas2012} to measurements from optical long-baseline interferometry in Figure \ref{comparison}, , Panel A. The $T_{\rm eff}$ surface is very metallicity insensitive $<10 $ K offsets due to metallicity effects) for $3200<T_{\rm eff}< 3900$ ,"
ALAXIPOL employs a three mirror. f Gregorian telescope with a 1.3 m oll-axis parabolic primary mirror.,"MAXIPOL employs a three mirror, f/1 Gregorian telescope with a 1.3 m off-axis parabolic primary mirror."
 “Phe elliptical secondary and tertiary reimaeing mirrcxs (21 and 18 em in diameter. respectively) are held at [liquid helium temperatures inside the receiver to reduce radiative oacding on the bolomoeters.," The elliptical secondary and tertiary reimaging mirrors (21 and 18 cm in diameter, respectively) are held at liquid helium temperatures inside the receiver to reduce radiative loading on the bolometers."
 To keep the instrumental polarization properties of the telescope constant. all telescope mirrors are fixed with res|oct to each other for all Light from the sky is reimaged toa 4x4 array of horns at the focal plane.," To keep the instrumental polarization properties of the telescope constant, all telescope mirrors are fixed with respect to each other for all Light from the sky is reimaged to a 4x4 array of horns at the focal plane."
 Observations are made in bands centered on 140 (11 anc| 420 Gllz (Av= 30 112)., Observations are made in bands centered on 140 GHz and 420 GHz $\Delta\nu \simeq$ 30 GHz).
 Phe twelve 140 Ον pl10toómoeters are optimized to measure the CMD and the four 420 (1 pheXometers are used to monitor foreground dust contamination., The twelve 140 GHz photometers are optimized to measure the CMB and the four 420 GHz photometers are used to monitor foreground dust contamination.
 The 10 FWHIAL Gaussian beam shape for the 140 (11 pheXometers is defined by a smooth walled. single-mode conical horn and a cold Lyot stop: the 420 Cllz photometers employ multi-mode Winston horns.," The $^\prime$ FWHM Gaussian beam shape for the 140 GHz photometers is defined by a smooth walled, single-mode conical horn and a cold Lyot stop; the 420 GHz photometers employ multi-mode Winston horns."
 Phe bolometers are maintained at 100 mix by an acliabatic demagnetization refrigerator 0] and a 300 mlx ο refrigerator., The bolometers are maintained at 100 mK by an adiabatic demagnetization refrigerator \cite{hagman} and a 300 mK $^{3}$ He refrigerator.
 A photograph of the focal plane and a cross-sectional view of the inside of the ervostat is shown in Figure L.., A photograph of the focal plane and a cross-sectional view of the inside of the cryostat is shown in Figure \ref{fig:cryostat}.
 For reference. t10 LWP and wire-ericl> polarizer that will be cliscussecl in Section 2.2. are mounted near the Lyot stop and at the focal plane. respectively.," For reference, the HWP and wire-grid polarizer that will be discussed in Section \ref{sec:polarimeter} are mounted near the Lyot stop and at the focal plane, respectively."
 The telescope was focused. before Light bx mapping tje detector beams with a chopped. LOO Watt halogen filament inmaged at infinity by a 38 inch on-axis parabolie mirror.," The telescope was focused before flight by mapping the detector beams with a chopped, 100 Watt halogen filament imaged at infinity by a 38 inch on-axis parabolic mirror."
 An absorptive. 0.75r inch thick plug of IE2eccosorb. MET10 was inserted in the optical path at the intermediate focus of the telescope to attenuate the intensity of warm loads during lab measurements.," An absorptive, 0.75 inch thick plug of Eccosorb MF110 was inserted in the optical path at the intermediate focus of the telescope to attenuate the intensity of warm loads during lab measurements."
 This attenuator was anti-rellection. (AR) coated with a 0.015 inch thick sheet of etehed Teflon., This attenuator was anti-reflection (AR) coated with a 0.015 inch thick sheet of etched Teflon.
 The calculat«cd transmission was ~ at 140 Giz 9)., The calculated transmission was $\sim$ at 140 GHz \cite{eccosorb}.
 A schematic of the NLAXIPOL insrument can be seen in Figure 2.., A schematic of the MAXIPOL instrument can be seen in Figure \ref{fig:gondola}.
 This illustration shows the pavloac without sun shielding so the telescope. receiver and attitude control subsystems are visible.," This illustration shows the payload without sun shielding so the telescope, receiver and attitude control subsystems are visible."
 Before Light. sun shielding was installed to protect all subsystems from solar radiation during davtime observations and to shield the telescope and receiver from spurious signals caused by sunlight and HI transmitters.," Before flight, sun shielding was installed to protect all subsystems from solar radiation during daytime observations and to shield the telescope and receiver from spurious signals caused by sunlight and RF transmitters."
 Phe ballling was made of Celotex aluminized. foam shecting and was painted white on all sun and earth-facing surfaces., The baffling was made of Celotex aluminized foam sheeting and was painted white on all sun and earth-facing surfaces.
 We selected a white paint pigmented with TiO» because this material has low solar absorptivity (~10%)) and. high infrared. emissivity (~90%)) a combination that compensated for the loss of convective cooling in the low-pressure balloon environment by providing aclequate radiative cooling., We selected a white paint pigmented with $_{2}$ because this material has low solar absorptivity $\sim$ ) and high infrared emissivity $\sim$ ) – a combination that compensated for the loss of convective cooling in the low-pressure balloon environment by providing adequate radiative cooling.
 This sun shield design successtulls maintained all instrument subsystenis within nominal temperature specifications during the claytime portion of the flight., This sun shield design successfully maintained all instrument subsystems within nominal temperature specifications during the daytime portion of the flight.
 A Iarge aluminum ground shield (also not. illustratec) was mounted to the inner frame to shield the main bearn of the telescope from terrestrial emission., A large aluminum ground shield (also not illustrated) was mounted to the inner frame to shield the main beam of the telescope from terrestrial emission.
 ‘Telescope attitude was feedback controlled., Telescope attitude was feedback controlled.
 “Phe Ilswheel mounted at the top of the gondola moved the telescope in azimuth while the elevation angle was adjusted with a Lincro actuator arm that nodeed the entire inner frame., The flywheel mounted at the top of the gondola moved the telescope in azimuth while the elevation angle was adjusted with a linear actuator arm that nodded the entire inner frame.
 A second motor mounted at the very top of the gondola further assisted in moving the telescope in azimuth by torquing the pavload against he balloon cabling., A second motor mounted at the very top of the gondola further assisted in moving the telescope in azimuth by torquing the payload against the balloon cabling.
 Phe azimuth feedback-Ioop relied on evros and a magnetometer and the elevation feedback relied. or ια l16-bit optical encoder., The azimuth feedback-loop relied on gyros and a magnetometer and the elevation feedback relied on a 16-bit optical encoder.
 The magnetometer was calibrated before Uieht: the olfset was measured to within a degree and the non-linearity was mapped and stored in a ookup table that was used by the on-board. pointing computer curing 1ight for making fine corrections., The magnetometer was calibrated before flight; the offset was measured to within a degree and the non-linearity was mapped and stored in a lookup table that was used by the on-board pointing computer during flight for making fine corrections.
 A second. looku> table was also implemented to account for variations in the magnetic field of the Earth as a function of longitude and [atitcle., A second lookup table was also implemented to account for variations in the magnetic field of the Earth as a function of longitude and latitude.
 Poining reconstruction for data analysis relies on the position of a reference star in one of two boresight Cohu 4910 CCD cameras., Pointing reconstruction for data analysis relies on the position of a reference star in one of two boresight Cohu 4910 CCD cameras.
" Phe camera used during daytime observaions was filtered. with a 695 nm Schott. glass filter ancl fitted with a 510 mnm Promaster Spectrum 7 reflective lens that orovided a 0.727 by 0.55"" field-of-view: the unfiltered nighttime camera used a 50r mm Fujinon lens that provided a 7.17 w 5.507 field-of-view.", The camera used during daytime observations was filtered with a 695 nm Schott glass filter and fitted with a 500 mm Promaster Spectrum 7 reflective lens that provided a $^{\circ}$ by $^{\circ}$ field-of-view; the unfiltered nighttime camera used a 50 mm Fujinon lens that provided a $^{\circ}$ by $^{\circ}$ field-of-view.
" Pixel size for the davtime ancl nighttime COnieras was 0.084 by 0.069"" and 0.84! by 0.69% respectivev."," Pixel size for the daytime and nighttime cameras was $^\prime$ by $^\prime$ and $^\prime$ by $^\prime$, respectively."
 Phe small field-of-view and the filter on the davtime camera were necessary to improve the ratio of star to sky brightness., The small field-of-view and the filter on the daytime camera were necessary to improve the ratio of star to sky brightness.
 With the combination. we detected stars of apparent visual maenitude 2 at balloon altitude.," With the combination, we detected stars of apparent visual magnitude 2 at balloon altitude."
 The two cameras and the telescope were aligned. before Hight to within a quarter of a degree., The two cameras and the telescope were aligned before flight to within a quarter of a degree.
 3olometer data and housekeeping signals were multiplexed into a single data stream that was telemetered to fixed erouncd stations during light., Bolometer data and housekeeping signals were multiplexed into a single data stream that was telemetered to fixed ground stations during flight.
 Phese signals were monitored in real time to ensure nominal operation of the instrument and because the crvogenic system needed to be manually eveled., These signals were monitored in real time to ensure nominal operation of the instrument and because the cryogenic system needed to be manually cycled.
 A new on-board data recorder was added to the experiment for the ALANIPOL-1 flight after à NASA data transmitter failecl curing NLAXIPOL-0., A new on-board data recorder was added to the experiment for the MAXIPOL-1 flight after a NASA data transmitter failed during MAXIPOL-0.
 The data recorder. which was custom designed and built by the Weizmann Institute of Science in Israel. consists of NIM modules cach containing an Altera FPGA chip and up to 128 Intel 16 MB flash memory chips.," The data recorder, which was custom designed and built by the Weizmann Institute of Science in Israel, consists of NIM modules each containing an Altera FPGA chip and up to 128 Intel 16 MB flash memory chips."
 With an uncompressed serial data rate of 160 kbps each module is capable of storing 28.4 hours of data., With an uncompressed serial data rate of 160 kbps each module is capable of storing 28.4 hours of data.
 Individual modules can be daisy-chained to each other to increase the total recording capacity., Individual modules can be daisy-chained to each other to increase the total recording capacity.
 Phe FPGA chip reads the incoming data stream and when it detects à. pre-programmed frame structure it controls the storage of the data on the memory chips., The FPGA chip reads the incoming data stream and when it detects a pre-programmed frame structure it controls the storage of the data on the memory chips.
 I. also controls the post-IHight export of the data from the memory chips into a computer through a standard parallel port., It also controls the post-flight export of the data from the memory chips into a computer through a standard parallel port.
 Power consumption during steady-state data recording is less than 0.25 Watt (at 5r V). and modules that are idle require only about 0.05 Watt.," Power consumption during steady-state data recording is less than 0.25 Watt (at 5 V), and modules that are idle require only about 0.05 Watt."
 For ALANIPOL-1 two mocdules containing 96 memory chips provided a total recording capacity of 42.6 hours., For MAXIPOL-1 two modules containing 96 memory chips provided a total recording capacity of 42.6 hours.
 Approximately 28 hours of pre-flight. ascent and at-Doat data were recorded.," Approximately 28 hours of pre-flight, ascent and at-float data were recorded."
 ALANIPOL analyzed the polarization of the millimeter-wave sky. with a rotating HWI and fixed wire-grid polarizer., MAXIPOL analyzed the polarization of the millimeter-wave sky with a rotating HWP and fixed wire-grid polarizer.
 While this technique is à well-known standard in astronomy. it is the first implementation in a CMD experiment.," While this technique is a well-known standard in astronomy, it is the first implementation in a CMB experiment."
 The strategy is, The strategy is
The black hole. candidate: 33394 was discovered: by Alarkert οἱ al. (,The black hole candidate 339–4 was discovered by Markert et al. (
1973) with the OSO7 satellite and was soon noted for its similarity in N-ravs to the classical black hole candidate NX.1 (Market ct al.,1973) with the OSO–7 satellite and was soon noted for its similarity in X-rays to the classical black hole candidate X–1 (Market et al.
 1973: Macjima et al., 1973; Maejima et al.
 1984: Dolan et al., 1984; Dolan et al.
 1987)., 1987).
 Phe source exhibits aperiodic and quasi-periodic modulations on time scales spanning from milliseconds to vears over a wide range of wavelengths., The source exhibits aperiodic and quasi-periodic modulations on time scales spanning from milliseconds to years over a wide range of wavelengths.
 It spends most of the time in the so-called. X-ray low state (LS) which has a power-law spectrum with spectral index oc1.5.2 (Ricketts 1983: \lacjima ct al., It spends most of the time in the so-called X-ray low state (LS) which has a power-law spectrum with spectral index $\alpha \sim 1.5-2$ (Ricketts 1983; Maejima et al.
 1984) and strong rms) band-Iimited noise (Nowak et al., 1984) and strong rms) band-limited noise (Nowak et al.
 1999: Belloni et al., 1999; Belloni et al.
 1999)., 1999).
 In the high state (LIS). it becomes brighter (in the 210 keV band) and exhibits an ultra-soft. spectral Component plus à steeper power-law (Alacjima ct al.," In the high state (HS), it becomes brighter (in the 2–10 keV band) and exhibits an ultra-soft spectral component plus a steeper power-law (Maejima et al."
 1984: Belloni et al., 1984; Belloni et al.
 1999). while the temporal variability is only a ew percent remis (CGrebeney et al.," 1999), while the temporal variability is only a few percent rms (Grebenev et al."
 1993: Belloni et al., 1993; Belloni et al.
 1999)., 1999).
 lt also shows a very high state (VIIS: Mivamoto et al., It also shows a very high state (VHS; Miyamoto et al.
 1991) with broad band noise of rms and 310 Lz quasi-»eriodic oscillations (QPOs) seen in its fast time variability. out. with a higher X-ray luminosity than in the LIS.," 1991) with broad band noise of rms and 3–10 Hz quasi-periodic oscillations (QPOs) seen in its fast time variability, but with a higher X-ray luminosity than in the HS."
 Recently. an intermediate state (LS) was reported by Ménndez and van der Ixlis (10997) and its spectral ancl timing properties are similar to the VIIS but with a much lower luminosity.," Recently, an intermediate state (IS) was reported by Ménndez and van der Klis (1997) and its spectral and timing properties are similar to the VHS but with a much lower luminosity."
 Finally. an ο state has also been reported (see Markert οἱ al.," Finally, an `off' state has also been reported (see Markert et al."
 1973: Motch et al., 1973; Motch et al.
 1985: Hovaisky ct al., 1985; Ilovaisky et al.
 1986: Asai et al., 1986; Asai et al.
 1998). in which the X-ray fast time variability is consistent with that seen in the LS (Ménndez van der Wlis 1997) while the energy. spectrum (power law with à of 1.52) is similar to the LS but with a 2I10keV. flux which is ~10 times lower or even fainter than in the LS.," 1998), in which the X-ray fast time variability is consistent with that seen in the LS (Ménndez van der Klis 1997) while the energy spectrum (power law with $\alpha$ of 1.5–2) is similar to the LS but with a 2–10keV flux which is $\sim 10$ times lower or even fainter than in the LS."
 lt has already been suspected that the oll” state is in fact a weak LS (see e.g. van der Ixlis 1995)., It has already been suspected that the `off' state is in fact a weak LS (see e.g. van der Klis 1995).
 A summary. of the cliflerent states and their properties is given in Table 1., A summary of the different states and their properties is given in Table 1.
 The optical counterpart. of. 3339.4. was identified bv Doxsev et abl (, The optical counterpart of 339–4 was identified by Doxsey et al. (
1970) as à V—18 blue star. but subsequent observations showed that it exhibited à. wide range of variability from V.—15.4 to 20.2 (Motch et al.,"1979) as a $V\sim 18$ blue star, but subsequent observations showed that it exhibited a wide range of variability from $V=15.4$ to 20.2 (Motch et al."
 1985: Corbet et al., 1985; Corbet et al.
 LOST) in its N-rav. LS and coll” state. while 1—16Is (Aloteh οἱ al.," 1987) in its X-ray LS and `off' state, while $V=16-18$ (Motch et al."
 1955) in the X-ray LIS., 1985) in the X-ray HS.
 Simultaneous optical/X-ray observations also showed a remarkable anti-correlation in the optical ancl soft. X-ray (36 keV) fluxes during a transition [rom X-ray LS to LIS (Aloteh et al., Simultaneous optical/X-ray observations also showed a remarkable anti-correlation in the optical and soft X-ray (3–6 keV) fluxes during a transition from X-ray LS to HS (Motch et al.
 1985). the cause of which is unknown.," 1985), the cause of which is unknown."
 However. llovaisksy et al. (," However, Ilovaisky et al. ("
1986) showed that there are times when the optical flux can be correlated. with the N-rav. Iuminositv.,1986) showed that there are times when the optical flux can be correlated with the X-ray luminosity.
 A possible orbital period of 14.8 hr from optical photometry was reported by Callanan ct al. (, A possible orbital period of 14.8 hr from optical photometry was reported by Callanan et al. (
1992).,1992).
 At present. there," At present, there"
"terminate at (he ""maximum mass point.","terminate at the ""maximum mass"" point."
 Comparing (he results for static ancl rotating stars. il is seen clearly that the rapid rotation increases noticeablv the mass (hat can be supported against collapse while lowering the central density of (he maximnunm-mass configuration.," Comparing the results for static and rotating stars, it is seen clearly that the rapid rotation increases noticeably the mass that can be supported against collapse while lowering the central density of the maximum-mass configuration."
 This is what one should expect. since. as already mentioned. rotation stabilizes the star against the gravitational pull providing an extra (centrifugal) repulsion.," This is what one should expect, since, as already mentioned, rotation stabilizes the star against the gravitational pull providing an extra (centrifugal) repulsion."
 The rotational effect on the mass-raclius relation is illustrated in Fig., The rotational effect on the mass-radius relation is illustrated in Fig.
 3 where the gravitational mass is given as a [function of the circumferential radius., 3 where the gravitational mass is given as a function of the circumferential radius.
 For rapid rotation at the IxXepler Irequency. a mass increase up to ~17% (Table 3) is obtained. depending on the EOS.," For rapid rotation at the Kepler frequency, a mass increase up to $\sim 17\%$ (Table 3) is obtained, depending on the EOS."
 The equatorial radius increases bv several kilometers. while the polar radius decreases by. several kilometers (see Fie.," The equatorial radius increases by several kilometers, while the polar radius decreases by several kilometers (see Fig."
 +) leading to an overall oblate shape of the rotating star., 4) leading to an overall oblate shape of the rotating star.
 Table 2 summarizes the properties (masses. radii and central energy densities) of the maximum-nmass nonrotating neutron star configurations.," Table 2 summarizes the properties (masses, radii and central energy densities) of the maximum-mass nonrotating neutron star configurations."
 Our studies on the effect of rapid rotation on (he upper mass limits for the four EOSs considered in the present paper are presented in Table 3., Our studies on the effect of rapid rotation on the upper mass limits for the four EOSs considered in the present paper are presented in Table 3.
 Dr each case the upper mass limit is attained [or à model at the mass-shedding limit where v=νε. with central density ~15%. below that of the static model with the largest mass.," In each case the upper mass limit is attained for a model at the mass-shedding limit where $\nu=\nu_k$, with central density $\sim
15\%$ below that of the static model with the largest mass."
 These findings are consistent with those by Friedinanοἱal.(1984). and Stereioulas&Friedman(1995)...," These findings are consistent with those by \citet{1984Natur.312..255F}
 and \citet{Stergioulas:1994ea}."
 Table 3 also provides an estimate of the upper limiting rotation rate of a neutron star., Table 3 also provides an estimate of the upper limiting rotation rate of a neutron star.
 In general. softer EOSs permit larger rotational [recuencies since the resulting stellar models are more centrally condensed (seee.e..Friedinanetal.1984)..," In general, softer EOSs permit larger rotational frequencies since the resulting stellar models are more centrally condensed \citep[see
e.g.,][]{1984Natur.312..255F}."
 In the last column of Table 3 we show the Ixepler frequencies computed via (he empirical relation proposed by Friedmanοἱal. (1989).., In the last column of Table 3 we show the Kepler frequencies computed via the empirical relation proposed by \citet{Friedman:1989}. .
 The uncertainty of Eq. (15)), The uncertainty of Eq. \ref{eq.15}) )
 is ~LOM. (see [or an improved version of theempirical lormula)., is $\sim 10\%$ (see \cite{HZ1989} for an improved version of theempirical formula).
 At the time of constructing, At the time of constructing
The purpose of this Appendix ts first to carry out the integration set out in Eq. (5)),The purpose of this Appendix is first to carry out the integration set out in Eq. \ref{m}) )
 and second to give an expression valid in the Klein-Nishina regime for the total spectrum emitted by a single electron scattering a mono-energetic beam of photons (Eq. 7))., and second to give an expression valid in the Klein-Nishina regime for the total spectrum emitted by a single electron scattering a mono-energetic beam of photons (Eq. \ref{thomson}) ).
 The fraction of scattered photons per time. energy and steradian is given by Eq. (6)).," The fraction of scattered photons per time, energy and steradian is given by Eq. \ref{thom}) ),"
 which can be expanded using Eqs. (1--5)), which can be expanded using Eqs. \ref{e}- \ref{m}) )
" where primed (unprimed) quantities are measured in the electron (observer) frame. 4). wo=cos9"". uio=cosP. ur,=cos6, ete."," where primed (unprimed) quantities are measured in the electron (observer) frame, $\mu_{\Theta'}\equiv\cos\Theta'=\mu'\mu'_1+\sin\theta'\sin
\theta'_1 \cos(\phi'_1-\phi')$ , $\mu'=\cos\theta'$, $\mu_0=\cos\theta_0$, $\mu'_0=\cos\theta'_0$ etc."
" Re-arranging the last Dirac and performing the three integrations yields The integration over Q, to obtain the full spectrumxi of radiation emitted by the electron is simplified if y;>>1.", Re-arranging the last Dirac and performing the three integrations yields The integration over $\Omega_1$ to obtain the full spectrum of radiation emitted by the electron is simplified if $\gamma_{\rm e}\gg1$.
" In that case. equivalent to saying the outgoing photon is emitted along theBui direction of electron motion when y, >> |."," In that case, which is equivalent to saying the outgoing photon is emitted along the direction of electron motion when $\gamma_{\rm e}\gg1$."
Thelast Dirae can then be rewritten as a function of μι; The integration over πρQ; is now straightforward. giving for the total spectrum: Relativistic kinematics gives the domain of variation of the scattered photon energy ej 1n the observer frame.," The last Dirac can then be rewritten as a function of $\mu_1$: The integration over $\Omega_1$ is now straightforward, giving for the total spectrum: Relativistic kinematics gives the domain of variation of the scattered photon energy $\epsilon_1$ in the observer frame."
 The maximum €. and minimum e. energies in the spectrum are :, The maximum $\epsilon_{+}$ and minimum $\epsilon_{-}$ energies in the spectrum are :
2003).,.
. Therefore. this corona could be the reason that millisecoucd burst oscillatious are uot observed at 87<2 in Figure L..," Therefore, this corona could be the reason that millisecond burst oscillations are not observed at $S_Z \lesssim 2$ in Figure \ref{fig:cc}."
 The two persistent uullisecoud pulsars provide the ouly exceptions to the above troeud: all of the burst oscillations. aud all of the N-rav bursts; ire observed in the hard portion of the Z-track on the color-color diagram that corresponds to low AL.," The two persistent millisecond pulsars provide the only exceptions to the above trend: all of the burst oscillations, and all of the X-ray bursts, are observed in the hard portion of the Z-track on the color-color diagram that corresponds to low $\dot{M}$."
 The main proposed difference between the millisecond pulsars and the other bursters is that the former have strouger magnetic fields (Chakrabartyctal.2003)., The main proposed difference between the millisecond pulsars and the other bursters is that the former have stronger magnetic fields \citep{cha03}.
. It as therefore plausible that oscillations are observed at low inferred AL for the pulsars because magnetic effects either cuhance the oscillation auplitudes. or supress a scattering corona.," It is therefore plausible that oscillations are observed at low inferred $\dot{M}$ for the pulsars because magnetic effects either enhance the oscillation amplitudes, or supress a scattering corona."
 The properties of X-ray bursts also ave correlated with he accretion rate onto the neutron star. but iu a mauner hat additionally depends ou the neutron star rotation vate.," The properties of X-ray bursts also are correlated with the accretion rate onto the neutron star, but in a manner that additionally depends on the neutron star rotation rate."
 We expect the burst properties to be determined by wo factors (Fujimoto.Wanawa.&Mivaji1981:Fushiki&Lamb1987:Bildsten 2000).," We expect the burst properties to be determined by two factors \citep[Fujimoto, Hanawa, \& Miyaji 1981;][]{fl87,bil00}."
 First. as AL inercases. he temperature at the burning laver also increases. and lus the coluun density of helimm required to trigger a must generally decreases.," First, as $\dot{M}$ increases, the temperature at the burning layer also increases, and thus the column density of helium required to trigger a burst generally decreases."
 Therefore. if the accretion is spherically sviuuetric; X-ray bursts that occur at high AT should be weaker.," Therefore, if the accretion is spherically symmetric, X-ray bursts that occur at high $\dot{M}$ should be weaker."
 Second. there is a competition between how quickly a sufficient column deusitv is accunmlated such that hein burning is unstable. and how quickly lvdrogen in the accreted material can be stably fused into heli.," Second, there is a competition between how quickly a sufficient column density is accumulated such that helium burning is unstable, and how quickly hydrogen in the accreted material can be stably fused into helium."
 The local accretion rate per unit area (01) actually drives the competition. but the accretion is generally assumed to occur with spherical svuuuetiy.," The local accretion rate per unit area $\dot{m}$ ) actually drives the competition, but the accretion is generally assumed to occur with spherical symmetry."
 At values of ij thought to correspond to the lower eud of those commonly observed roni bursting LAINBs (equivalent to elobal accretion rates of 0.01MpgqacM0.05 Aa). ID is burned iuto Io faster han it can be accreted. so the N-vav burst occurs from mire Πο fuel.," At values of $\dot{m}$ thought to correspond to the lower end of those commonly observed from bursting LMXBs (equivalent to global accretion rates of $0.01\dot{M}_{\rm Edd} < \dot{M} < 0.05\dot{M}_{\rm Edd}$ ), H is burned into He faster than it can be accreted, so the X-ray burst occurs from pure He fuel."
" Αν increases (0.05AaazMx Mgaa). IT is accreted faster than it can be burned into Ie. so the mrsts occur from mixed IE/IHe The X-aav burst properties change between these regenues because the relative amount of II aud Ie in the ""el determines how rapidly the uuclear energy is released during the bursts."," As $\dot{m}$ increases $0.05\dot{M}_{\rm Edd} < \dot{M} < \dot{M}_{\rm Edd}$ ), H is accreted faster than it can be burned into He, so the bursts occur from mixed H/He The X-ray burst properties change between these regimes because the relative amount of H and He in the fuel determines how rapidly the nuclear energy is released during the bursts."
 Welimm burns via a strong triple-a process that releases energv quickly. so the low-; Ie bursts are more likely to exhibit radius expansion.," Helium burns via a strong $\alpha$ process that releases energy quickly, so the $\dot{m}$ He bursts are more likely to exhibit radius expansion."
 Iu contrast. IT serves to iioderate the We burning at the start of the burst. and ouly burns through a slow rp-capture process onto the products of Πο burning at the cud of a burst.," In contrast, H serves to moderate the He burning at the start of the burst, and only burns through a slow $rp$ -capture process onto the products of He burning at the end of a burst."
 Therefore. the highli-ib» mixed ΗΕΠο bursts should last longer. and be less likely to exhibit radius expausion.," Therefore, the $\dot{m}$ mixed H/He bursts should last longer, and be less likely to exhibit radius expansion."
" As n result of these two effects; as the global accretion rate onto the neutron star increases, the N-rav bursts should become weaker aud less likely to exhibit radius expansion."," As a result of these two effects, as the global accretion rate onto the neutron star increases, the X-ray bursts should become weaker and less likely to exhibit radius expansion."
 This is the case for the slow rotators. but in the fast rotators the bursts areκο to exhibit radius expansion at high AJ.," This is the case for the slow rotators, but in the fast rotators the bursts are to exhibit radius expansion at high $\dot{M}$."
 One possible explanation for this is that the accretion is not spherically πιοο (Bildsten2000)., One possible explanation for this is that the accretion is not spherically symmetric \citep{bil00}.
. Iu. particular. if the local accretion rate Gi) as the elobal rate (AD) increases in the fast rotators. then the ligh-iy bursts with radius expansion could occur at low AL. and the low-ib bursts without radius expansion could occur at high AT.," In particular, if the local accretion rate $\dot{m}$ ) as the global rate $\dot{M}$ ) increases in the fast rotators, then the $\dot{m}$ bursts with radius expansion could occur at low $\dot{M}$, and the $\dot{m}$ bursts without radius expansion could occur at high $\dot{M}$."
 Th contrast. the change in burst properties iu the slow rotators appears consistent with a local i; that increases as the global AT does.," In contrast, the change in burst properties in the slow rotators appears consistent with a local $\dot{m}$ that increases as the global $\dot{M}$ does."
" It is then possible that the rotation rate of the neutron star influences how accreted material spreads over its surface. either through a lower effective surface eravity or a stronger Coriolis force in the fast rotators.: (οι,Bild-sten2000:Spitkovskyetal. 2002)."," It is then possible that the rotation rate of the neutron star influences how accreted material spreads over its surface, either through a lower effective surface gravity or a stronger Coriolis force in the fast rotators.; \citep[e.g.][]{bil00,slu02}."
. Although we have determined observationally that the rotation rate of a neutron star influences how the properties of thermonuclear X-ray burst change with the accretion rate onto the neutron star. it is not clear what causes the observed correlations.," Although we have determined observationally that the rotation rate of a neutron star influences how the properties of thermonuclear X-ray burst change with the accretion rate onto the neutron star, it is not clear what causes the observed correlations."
 Further progress should be made bv studyius this sample of sources to see how the burst time scales.peak fluxes. fiuences. aud reciueuce times chauge with the accretion rates.," Further progress should be made by studying this sample of sources to see how the burst time scales,peak fluxes, fluences, and recurrence times change with the accretion rates."
"Can the two type of SED fits. explaining the Balmer break with ""old"" stars or by the presence of emission lines. be distinguished observationally?","Can the two type of SED fits, explaining the Balmer break with “old” stars or by the presence of emission lines, be distinguished observationally?"
" For the SED fit of 23-6714 shown in reftig,ed| forexample. thecurrentS FR~ 800yr!."," For the SED fit of 23-6714 shown in \\ref{fig_sed1} for example, the current SFR $\sim 800$."
. From this the expected fflux is of the order 3.x107!'em™.. clearly beyond the reach of current spectroscopic facilities at~ 4.5um.," From this the expected flux is of the order $3.\times 10^{-16}$, clearly beyond the reach of current spectroscopic facilities $\sim$ 4.5."
.. This object being the brightest of the sample it appears that direct detections of the restframe optical emission lines need to await future facilities. such as the JWST.," This object being the brightest of the sample it appears that direct detections of the restframe optical emission lines need to await future facilities, such as the JWST."
 In principle deep photometry at longer wavelengths. where the expected strength of emission lines decreases. could help to disentangle the two solutions.," In principle deep photometry at longer wavelengths, where the expected strength of emission lines decreases, could help to disentangle the two solutions."
" However. the current 5.8 and 8.0 ddata from Spitzer is not deep enough to rule out the ""old SEDs. as shown for the two brightest objects 23-6714 and 31-2185."," However, the current 5.8 and 8.0 data from Spitzer is not deep enough to rule out the “old” SEDs, as shown for the two brightest objects 23-6714 and 31-2185."
 Although the spectrum of the ~700 Myr old population displays an excess with respect to the 5.8 and/or 8.0 ffluxes. its significance is <2c at best.," Although the spectrum of the $\sim 700$ Myr old population displays an excess with respect to the 5.8 and/or 8.0 fluxes, its significance is $\la 2 \sigma$ at best."
 Future. deeper observations should be able to set firmer limits on this issue.," Future, deeper observations should be able to set firmer limits on this issue."
 In any case we note that neglecting nebular emission for high-z galaxies with ongoing massive star formation — as testified bby their eemission often used to confirm their redshift spectroscopically — is physically inconsistent and its effect on the determination of their physical parameters should be taken into account., In any case we note that neglecting nebular emission for $z$ galaxies with ongoing massive star formation — as testified by their emission often used to confirm their redshift spectroscopically — is physically inconsistent and its effect on the determination of their physical parameters should be taken into account.
 As shown above. the inclusion of nebular emission (lines and continuum) can alter the physical properties of galaxies determined from their SED.," As shown above, the inclusion of nebular emission (lines and continuum) can alter the physical properties of galaxies determined from their SED."
" For example. considering just the average properties of our best fit solutions for the 10 objects fitted here we obtain = 400 Myr. M,=1.1x10M... and Ay=0.16 for standard Bruzual Charlot spectral templates when metallicity is also varied."," For example, considering just the average properties of our best fit solutions for the 10 objects fitted here we obtain $\overline{t_\star} \approx$ 400 Myr, $\overline{M_\star}=1.1 \times 10^{10}$, and $\overline{A_V}=0.16$ for standard Bruzual Charlot spectral templates when metallicity is also varied."
" Including nebular emission we obtain 7x 120 Myr. M,=7.9x10°M... and a higher extinction Ay=0.34."," Including nebular emission we obtain $\overline{t_\star} \approx$ 120 Myr, $\overline{M_\star}=7.9 \times 10^{9}$, and a higher extinction $\overline{A_V}=0.34$."
" Although indicative of the trend obtained with nebular emission. the reader should be aware that large deviations are obtained and uncertainties should be treated properly refs, iscuss))."," Although indicative of the trend obtained with nebular emission, the reader should be aware that large deviations are obtained and uncertainties should be treated properly \\ref{s_discuss}) )."
 Clearly the most striking result is that the average age of the z=6 galaxies may be decreased typically by à factor ~3 compared to earlier studies (cf.22)..," Clearly the most striking result is that the average age of the $z\approx 6$ galaxies may be decreased typically by a factor $\sim 3$ compared to earlier studies \cite[cf.][]{Yan06,Eyles07}."
 Translated to the average redshift of formation zj;4; this would imply a shift from ο7 9., Translated to the average redshift of formation $z_{\rm form}$ this would imply a shift from $z_{\rm form} \sim$ 9.
 to 6.6. assuming =5.873 for the average redshift of our sample EEO07).," to 6.6, assuming $\overline{z}=5.873$ for the average redshift of our sample E07)."
 For only one out of 10 objects do we find best fit ages of >200 Myr. aa formation redshift beyond 7.," For only one out of 10 objects do we find best fit ages of $>200$ Myr, a formation redshift beyond 7."
 |f true. this means in particular that the contribution of the galaxies currently observed at z~6 to cosmic retonisation must have been negligible at 27. in contrast to the finding of E07.," If true, this means in particular that the contribution of the galaxies currently observed at $z \sim 6$ to cosmic reionisation must have been negligible at $z \ga 7$, in contrast to the finding of E07."
 The average extinction of Ay~0.34 mag ~ 0.08) we find. corresponding to a UV attenuation by a factor ~2.2 1s also worth noticing.," The average extinction of $A_V \sim 0.34$ mag $\sim$ 0.08) we find, corresponding to a UV attenuation by a factor $\sim 2.2$ is also worth noticing."
 With standard spectral templates. neglecting nebular emission. the average is Ay=0.16.," With standard spectral templates, neglecting nebular emission, the average is $A_V=0.16$."
 This result indicates that dust attenuation may be larger than previously thought for Lyman Break galaxies at this redshift., This result indicates that dust attenuation may be larger than previously thought for Lyman Break galaxies at this redshift.
 Por comparison. ? assume a UV attenuation factor of ~ 1.5 at zo6.," For comparison, \citet{Bouwens08} assume a UV attenuation factor of $\sim$ 1.5 at $z \sim 6$."
 df true and representative for the population of z~6 galaxies. it may imply an upward revision of the SFR density by ~50%.," If true and representative for the population of $z \sim 6$ galaxies, it may imply an upward revision of the SFR density by $\sim 50$."
. Finally stellar masses and hence the estimated stellar mass density of i-drop galaxies may also be slightly affected., Finally stellar masses and hence the estimated stellar mass density of i-drop galaxies may also be slightly affected.
" Also already mentioned. the average stellar mass of our sample M, is ~30 lower when we include nebular emission in the SED fits and leave the other assumptions unchanged."," Also already mentioned, the average stellar mass of our sample $\overline{M_\star}$ is $\sim 30$ lower when we include nebular emission in the SED fits and leave the other assumptions unchanged."
 The main objective of this work has been to examine the effect of nebular emission (ines and continua) on the derivation of physical parameters of z~6 galaxies through broad band SED fits., The main objective of this work has been to examine the effect of nebular emission (lines and continua) on the derivation of physical parameters of $z \sim 6$ galaxies through broad band SED fits.
 Although our work clearly indicates that this effect can significantly alter derived properties such as galaxy ages and extinction. we do not (yet) aim at determining absolute values.," Although our work clearly indicates that this effect can significantly alter derived properties such as galaxy ages and extinction, we do not (yet) aim at determining absolute values."
 Indeed both uncertainties/difficulties in the photometry as well as uncertainties and degeneracies in the models remain. and their effect needs to be quantified properly., Indeed both uncertainties/difficulties in the photometry as well as uncertainties and degeneracies in the models remain and their effect needs to be quantified properly.
 For example. the differences obtained using photometry from two sources (revisions) has been illustrated above.," For example, the differences obtained using photometry from two sources (revisions) has been illustrated above."
 Stellar ages and other quantities may also vary if different. more complex SF histories. different extinction. laws (possibly also. extinction differences between stars and the gas). different IMFs are allowed im the models.," Stellar ages and other quantities may also vary if different, more complex SF histories, different extinction laws (possibly also extinction differences between stars and the gas), different IMFs are allowed in the models."
 Last but not least. our and other SED fitting tools are also limitec by their dependence on stellar tracks. which are (?).. (?)..," Last but not least, our and other SED fitting tools are also limited by their dependence on stellar tracks, which are \citep{Vazquez07}, \citep[cf.]{Maraston06}
 \citep{Debarros09}."
scattered flux. uniform models unclerestimate e;,"scattered flux, uniform models underestimate $a$."
 At small uniform models overestimate à., At small uniform models overestimate $a$.
 A hierarchical model in which the star happens to be within a chunp has smaller errors., A hierarchical model in which the star happens to be within a clump has smaller errors.
 Adding a constant density decreases the errors αἱ e» ] so that the average a is ~ 0.5., Adding a constant density decreases the errors at $\sim$ 1 so that the average $a$ is $\sim $ 0.5.
 OF course. there is no error in the unrealistic case of the density being uniform throughout.," Of course, there is no error in the unrealistic case of the density being uniform throughout."
 An important question is. how reliably can we determine a(A) from observations of a parücular nebula covering a sienilicant range in wavelengths. implving differing opacities or optical depths?," An important question is, how reliably can we determine $a(\lambda)$ from observations of a particular nebula covering a significant range in wavelengths, implying differing opacities or optical depths?"
 In this case. we might hope to lind optical properties of the grains accurately even with uniform models.," In this case, we might hope to find optical properties of the grains accurately even with uniform models."
 Figure 7 shows the results for (wo wavelengths [or which the extinctions (both and 74)) differ bx a factor of 2: for example. AA)) if the extinction law is the Galactic average (Cardelli. Clavton. Mathis 1989) and Ry=AMM/LAGB)—AA)3.1.," Figure \ref{fig7}
 shows the results for two wavelengths for which the extinctions (both and ) differ by a factor of 2: for example, ) if the extinction law is the Galactic average (Cardelli, Clayton, Mathis 1989) and $R_V=A(V)/[A(B)-A(V)]=3.1$."
" The figure displays the ratio of a(,—22) to a(m==1). both albedos derived from uniform models. for the same hierarchical model as shown in Figure L.."," The figure displays the ratio of $a$ 2) to $a$ =1), both albedos derived from uniform models, for the same hierarchical model as shown in Figure \ref{fig1}."
 The abscissa in the ligure is for the lower extinction., The abscissa in the figure is for the lower extinction.
" All points in the figure were computed with the α = 0.6. so the ""true"" value of the albedo ratio is unity."," All points in the figure were computed with the $a$ = 0.6, so the “true” value of the albedo ratio is unity."
" We see that uniform models have difficulty in predicting even relative values of αλ). since the spread in a(n,=2)/a(n,1) is ~40%.."," We see that uniform models have difficulty in predicting even relative values of $a(\lambda)$, since the spread in $a(\tau_0=2)/a(\tau_0=1)$ is $\sim$."
 The figure gives the impression that if is >1.5. uniform models would derive a greater albedo at short wavelengths (say. AA)) than at longer Που a factor of 2 in the extinction).," The figure gives the impression that if is $\ge1.5$, uniform models would derive a greater albedo at short wavelengths (say, ) than at longer for a factor of 2 in the extinction)."
 For the opposite is [ο to be the case (Murtliv οἱ al., For the opposite is found to be the case (Murthy et al.
 1993: Witt et al., 1993; Witt et al.
 1993)., 1993).
 From Figure 7 we would conclude that a( <a( AA)) is robust. since Wilt οἱ al. (," From Figure \ref{fig7} we would conclude that $a$ $<a$ ) is robust, since Witt et al. ("
1993) suggest. AA)) ~(1.5 2).,1993) suggest ) $\sim$ (1.5 – 2).
 However. models with the star embedded in a clump have the albedo ratios shift from mostly <1 to mostly >1 at AA)) ~2.5. considerably greater than shown in the figure.," However, models with the star embedded in a clump have the albedo ratios shift from mostly $<1$ to mostly $>1$ at ) $\sim2.5$, considerably greater than shown in the figure."
 These models would produce an illusory result. af AA)) <a( AA))., These models would produce an illusory result $a$ ) $<a$ ).
 The basic problem is that the importance of a clump depends upon its optical depth., The basic problem is that the importance of a clump depends upon its optical depth.
 Ifa clump becomes optically thin. its existence no longer very signilicant to the radiative transport.," If a clump becomes optically thin, its existence no longer very significant to the radiative transport."
 Sinularly. iucreasine the optical depth of an optically thick chump has a limited effect on (he radiation.," Similarly, increasing the optical depth of an optically thick clump has a limited effect on the radiation."
 While the density structure within the ISM is independent of the, While the density structure within the ISM is independent of the
Strohmayer (2002) performed a coherent timing analysis of and data and found that the 569 sec period in IX 11914|24 was spinning up at a rate of 8253.1007 Hz/s 101 s/s).,Strohmayer (2002) performed a coherent timing analysis of and data and found that the 569 sec period in RX J1914+24 was spinning up at a rate of $\pm3\times10^{-18}$ Hz/s $\times10^{-12}$ s/s).
 This was later refinec with the acelition ol cata to 7.00.810D είς (23.107 s/s. Strohmaver 2O004a)," This was later refined with the addition of data to $\pm0.8\times10^{-18}$ Hz/s $\times10^{-12}$ s/s, Strohmayer 2004a)."
 This result. was consistent with he expected spin up rate if the systems was being driven entirely by gravitational radiation and the secondary star das a LOW Ll., This result was consistent with the expected spin up rate if the systems was being driven entirely by gravitational radiation and the secondary star has a low mass.
 The spin-up in both RX JOSOG|15 and RAN J1914|24 las been met with some degree of scepticism (ee Wouclt Warner 2003)., The spin-up in both RX J0806+15 and RX J1914+24 has been met with some degree of scepticism (eg Woudt Warner 2003).
 To show the effect of the spin-up in UX J1914]24 we show in Figure 2. theος.UA. (which. were extracted from the archive and reduced in the same manner as Strohmayer 20042) and data folded on the constant period term (αμ) of Strohmaver. (2004a).," To show the effect of the spin-up in RX J1914+24 we show in Figure \ref{period} the, (which were extracted from the archive and reduced in the same manner as Strohmayer 2004a) and data folded on the constant period term $\nu_{o}$ ) of Strohmayer (2004a)."
 Phe phase of the rise to niaximum Εαν has continued to arrive progressively earlier since the epoch of the data: it is clear that WA J1914|24 is spinning up., The phase of the rise to maximum flux has continued to arrive progressively earlier since the epoch of the data: it is clear that RX J1914+24 is spinning up.
 Lt is highly unlikely that the true period is a side-band power peak Strohmaver (20042)., It is highly unlikely that the true period is a side-band power peak Strohmayer (2004a).
 For the phased light curves shown in Figure 2. we estimated the start ofthe bright phase by noting the phase at which the count rate rose by a significant level compared o the previous phase bin., For the phased light curves shown in Figure \ref{period} we estimated the start of the bright phase by noting the phase at which the count rate rose by a significant level compared to the previous phase bin.
 We show how the start of the sight phase varies over time in Figure 3. together with an estimate of the uncertainty in the phase., We show how the start of the bright phase varies over time in Figure \ref{pdot} together with an estimate of the uncertainty in the phase.
 We fitted. these points with a constant 2 term: the best [it is shown as a solid line in Figure 2.., We fitted these points with a constant $\dot{P}$ term: the best fit is shown as a solid line in Figure \ref{period}.
 We note that the shift in. phase or à constant 2 term is not linear over time (see Cropper et al 2004)., We note that the shift in phase for a constant $\dot{P}$ term is not linear over time (see Cropper et al 2004).
 We find a spin-up rate of 3.174010.1072 s/s 10*° Hz/s)., We find a spin-up rate of $\pm0.10\times10^{-12}$ s/s $\times10^{-17}$ Hz/s).
 Because the errors on the phase of he rise to the bright phase are not strictly lo errors. we determined. the error using a bootstrap method.," Because the errors on the phase of the rise to the bright phase are not strictly $\sigma$ errors, we determined the error using a bootstrap method."
 This gave a larger error than the formal error to the fit (0.07.10.E? s/s)., This gave a larger error than the formal error to the fit $\times10^{-12}$ s/s).
 Phe spin-up rate which we determine is slightly greater than the spin-up rate derived by Strohmaver (20042)., The spin-up rate which we determine is slightly greater than the spin-up rate derived by Strohmayer (2004a).
 We also show the long term X-ray intensity of tX. J1914]24 in the lower panel of Figure 3..., We also show the long term X-ray intensity of RX J1914+24 in the lower panel of Figure \ref{pdot}.
 Fhese. values were obtained by converting the peak X-ray. count rate in each epoch to the equivalentROSAL LRT count rate (since most observations were determined using this instrument)., These values were obtained by converting the peak X-ray count rate in each epoch to the equivalent HRI count rate (since most observations were determined using this instrument).
 We used the spectral parameters shown Table 2. ancl the tool (Mukai 1993) to convert the count rate in other detectors to that expected for theROSAL," We used the spectral parameters shown Table \ref{fits}
 and the tool (Mukai 1993) to convert the count rate in other detectors to that expected for the HRI."
 LR We note that the largest residual to the fit occurred after the peak X-ray [lux had. decreased. [rom a short epoch of enhanced X-ray emission., We note that the largest residual to the fit occurred after the peak X-ray flux had decreased from a short epoch of enhanced X-ray emission.
 However. the deviation is small. (1.90). so therefore probably not significant.," However, the deviation is small, $\sigma$ ), so therefore probably not significant."
 Ramsay et al (2000) showed that the peak of the X-ray and optical band emission were offset in phase., Ramsay et al (2000) showed that the peak of the X-ray and optical band emission were offset in phase.
 Since their and / band data were taken within 3 months of each other. this phase olfset is still valid despite the spin-up of the system.," Since their and $I$ band data were taken within 3 months of each other, this phase offset is still valid despite the spin-up of the system."
 Yo fully utilise the ata [rom the two observations (cf ‘Table 1). we combined the EPIC pn event files from the two epochs (we also did this for EPIC MOS2).," To fully utilise the data from the two observations (cf Table 1), we combined the EPIC pn event files from the two epochs (we also did this for EPIC MOS2)."
 We used single ancl double events for the pn spectrum. and single to cquacdruple events for the MOS spectrum. and those events with PLAG=0.," We used single and double events for the pn spectrum, and single to quadruple events for the MOS spectrum, and those events with FLAG=0."
 We ereated response and auxillary files using the SAS tasks andrmfgen., We created response and auxillary files using the SAS tasks and.
 Phe spectra are shown in Figure 4.., The spectra are shown in Figure \ref{spec_fig}.
 Ramsay et al (2000) Πίος X-ray spectra taken usingROSA and.A using an absorbed blackbocdy model., Ramsay et al (2000) fitted X-ray spectra taken using and using an absorbed blackbody model.
 We used. (Arnaud 1996) to fit the spectra. binned them so there was a minimum. of 50 counts per min and we applied a low energy. eut-olf of 0.25keV. The spectra are not well fitted using such a model.," We used (Arnaud 1996) to fit the spectra, binned them so there was a minimum of 50 counts per min and we applied a low energy cut-off of 0.25keV. The spectra are not well fitted using such a model."
 Various other models were fitted including: an absorbed two emperature blackbody model: a thermal plasma model with varving clement abundances: the previous model together with a blackbody: and an emission model together with a ine inabsorption around 0.55keV. The only model which even approached a reasonable fit (4511.63. 79 dol in EPIC on) was an absorbed: blackhocky plus broad. Gaussian line in emüssion at 0.59 keV: this was an unexpected. result (the fit to the EPIC) MOS2 spectra was poorer giving . 29 dof the response of the MOS cameras has," Various other models were fitted including; an absorbed two temperature blackbody model; a thermal plasma model with varying element abundances; the previous model together with a blackbody; and an emission model together with a line in around 0.55keV. The only model which even approached a reasonable fit 1.63, 79 dof in EPIC pn) was an absorbed blackbody plus broad Gaussian line in emission at 0.59 keV: this was an unexpected result (the fit to the EPIC MOS2 spectra was poorer giving 2.20, 29 dof – the response of the MOS cameras has"
mmean opacity (??7.anclreferencestherein)...,"mean opacity \citep[][and references therein]{dixon2009, mcquinn2009c}."
 Lhe redshift at which Z5(2) begins to decline. as well as the peak value of wwill depend primarily on the evolution of wwith recdshilt.," The redshift at which $T_0(z)$ begins to decline, as well as the peak value of will depend primarily on the evolution of with redshift."
 Finally. we note that our values of ab 2294 Sare consistent with a broad range of hydrogen reionization scenarios.," Finally, we note that our values of at $z \simeq 4-5$ are consistent with a broad range of hydrogen reionization scenarios."
 At 2296. lis expected to decrease rapidly following rrejonization due to the combined. elfects of acdiabatic expansion and Compton cooling (?2)..," At $z > 6$, is expected to decrease rapidly following reionization due to the combined effects of adiabatic expansion and Compton cooling \citep{me1994,huignedin1997}. ."
 Leven for rreionization at z10. however. the asymptotic value of aat 2d—5 is expected to be —SO000 KIN (2)... which is consistent with our measurements.," Even for reionization at $z
\gtrsim 10$, however, the asymptotic value of at $z \simeq 4-5$ is expected to be $\sim$ K \citep{huihaiman2003}, which is consistent with our measurements."
 This temperature assumes a dquasar-ike background. with a=1.5. while a stellar spectrum with à—3.0 would lower ον —20 per cent (2)..," This temperature assumes a quasar-like background with $\alpha = 1.5$, while a stellar spectrum with $\alpha = 3.0$ would lower by $\sim$ 20 per cent \citep{huihaiman2003}."
 On the other hand. our estimates of the temperature at z2»4 may be too high by roughly he same amount for the reasons described in Section 7?..," On the other hand, our estimates of the temperature at $z > 4$ may be too high by roughly the same amount for the reasons described in Section \ref{sec:thermal_history}."
 Given the uncertainties in both the measurements. and he thermal models. therefore. we conclude. that our empoerature measurements at z24 do not obviously require a recent (2< 10) temperature boost from rrejonization.," Given the uncertainties in both the measurements and the thermal models, therefore, we conclude that our temperature measurements at $z > 4$ do not obviously require a recent $z < 10$ ) temperature boost from reionization."
 This in in contrast to earlier. works (77) owed. on temperatures. estimates that were considerably igher than the ones measured here.," This in in contrast to earlier works \citep{theuns2002b,huihaiman2003} based on temperatures estimates that were considerably higher than the ones measured here."
 We note that the rapid cooling expected after reionization completes means that our measurements are also consistent with rrejonization at z«10 for a wide variety of reionization redshifts and. initial temperature boosts (727)..," We note that the rapid cooling expected after reionization completes means that our measurements are also consistent with reionization at $z < 10$ for a wide variety of reionization redshifts and initial temperature boosts \citep{huihaiman2003,furoh2008b}."
 2 recently demonstrated. that. local constraints on the timing of hydrogen reionization can also be placed [from temperature measurements in the near-zones of z~6 QSOs., \citet{bolton2010} recently demonstrated that local constraints on the timing of hydrogen reionization can also be placed from temperature measurements in the near-zones of $z \sim 6$ QSOs.
 Combining near-zone measurements with our results from the forest should in the future provide a powerful means of studying the history hydrogen reionization., Combining near-zone measurements with our results from the forest should in the future provide a powerful means of studying the history hydrogen reionization.
 We have produced new measurements of the IGM emperature over 2.0xz48S., We have produced new measurements of the IGM temperature over $2.0 \le z \le 4.8$.
 Phese are the most sensitive such measurements to date. and the first in the general IGM at z>4.5.," These are the most sensitive such measurements to date, and the first in the general IGM at $z > 4.5$."
 Our analysis uses a new statistic. the curvature. o characterize the small-scale structure of the forest in a manner that is highly sensitive to. the IGAL temperature while allowing systematic errors to. he ellectively minimized.," Our analysis uses a new statistic, the curvature, to characterize the small-scale structure of the forest in a manner that is highly sensitive to the IGM temperature while allowing systematic errors to be effectively minimized."
 We further focus on measuring the emperature at an optimal overdensity for cach recshilt. where the temperature is nearly a one-to-one function of he mean absolute curvature.," We further focus on measuring the temperature at an optimal overdensity for each redshift, where the temperature is nearly a one-to-one function of the mean absolute curvature."
 Phe temperature at the mean density is then inferred from our mumeasurements for dillerent values of the slope for the emperature-densitv relation., The temperature at the mean density is then inferred from our measurements for different values of the slope for the temperature-density relation.
 Our main conclusions are as ollows: Two kev challenges remain to fully characterizing the thermal evolution of the IGM over this redshift range. at least in a volume-averaged sensed.," Our main conclusions are as follows: Two key challenges remain to fully characterizing the thermal evolution of the IGM over this redshift range, at least in a volume-averaged sensed."
 “Phe most pressing issue is to establish the shape of the temperature-density relation and its evolution with redshift., The most pressing issue is to establish the shape of the temperature-density relation and its evolution with redshift.
 This will clarify the evolution of du(z) at z«4. and allow us to identify the end of rreionization based on the redshift at which ppeaks.," This will clarify the evolution of $T_{0}(z)$ at $z < 4$, and allow us to identify the end of reionization based on the redshift at which peaks."
 The second. challenge is to determine the amount of Jeans smoothing in the IGM. which can possibly be constrained using lines of sight towards pairs of QSOs (7)..," The second challenge is to determine the amount of Jeans smoothing in the IGM, which can possibly be constrained using lines of sight towards pairs of QSOs \citep{peeples2010b}."
 We leave both of these tasks for the future. butnote that the measurements presented here provide a basis for making further progress.," We leave both of these tasks for the future, butnote that the measurements presented here provide a basis for making further progress."
 For example. our values may be combined with Voigt. profile measurements to determine oover 272.d.," For example, our values may be combined with Voigt profile measurements to determine over $z \sim 2-4$."
 Ultimately one would like to characterize not onlv the evolution of the volume-averaged thermal state of the IGM but also the evolution in its scatter., Ultimately one would like to characterize not only the evolution of the volume-averaged thermal state of the IGM but also the evolution in its scatter.
 A key characteristic of rreionization should bethat growing bbubbles around QSOs introduce. spatial Ductuations in both ionization and temperature (222?7?).. even if these are partially damped: by. radiative transfer. effects," A key characteristic of reionization should bethat growing bubbles around QSOs introduce spatial fluctuations in both ionization and temperature \citep{sokasian2002, gleser2005, lai2006, paschos2007, furoh2008b}, , even if these are partially damped by radiative transfer effects"
"within the Kuiper Belt (30—50 AU), while others could make even closer approaches.","within the Kuiper Belt $30-50$ AU), while others could make even closer approaches."
" If each white dwarf we observe has a distributions of asteroids with the geometry of the Oort Cloud, we have a good chance to detect transits withKep/er monitoring of even a single white dwarf."," If each white dwarf we observe has a distributions of asteroids with the geometry of the Oort Cloud, we have a good chance to detect transits with monitoring of even a single white dwarf."
" The Kuiper Belt itself has a scattered component in which the average orbital inclination angle, i. is 12. while individual orbits can be even more inclined. ("," The Kuiper Belt itself has a scattered component in which the average orbital inclination angle, $i,$ is $12^o,$ while individual orbits can be even more inclined. ("
See Sheppard 2006 and references therein.),See Sheppard 2006 and references therein.)
" Therefore. if the asteroid systems of white dwarts have a geometry similar to that of the scattered Kuiper Belt, and if we could monitor 10—15 white dwarts, we would have a good chance that the equivalent of the scattered ""Kuiper"" Belt of one or more of them was inclined toward our line of sight."," Therefore, if the asteroid systems of white dwarfs have a geometry similar to that of the scattered Kuiper Belt, and if we could monitor $10-15$ white dwarfs, we would have a good chance that the equivalent of the scattered “Kuiper” Belt of one or more of them was inclined toward our line of sight."
" Assuming a spherical distribution, we can quantify the probability of detecting a transit as à function ofa as follows."," Assuming a spherical distribution, we can quantify the probability of detecting a transit as a function of $a$ as follows."
 We compute the number of asteroids that would have to have periastrons at a particular value of@ in order to have a probability near unity of detecting the transit., We compute the number of asteroids that would have to have periastrons at a particular value of $a$ in order to have a probability near unity of detecting the transit.
" Given em, then for a= (0.1, 1. 10. 100) AU, the number of such asteroids is (1.9x10°, 1.9x107, 1.9x10°, 1.9x 105)."," Given $(R_{\rm wd} + R_{\rm ast}) = 8 \times 10^8$ cm, then for $a=$ (0.1, 1, 10, 100) AU, the number of such asteroids is $1.9\times 10^3$ , $1.9\times 10^4$, $1.9\times 10^5$, $1.9\times 10^6$ )."
 These numbers assume that the interval Τ during which continuous observations occur spans the orbital periods of the asteroids., These numbers assume that the interval $T$ during which continuous observations occur spans the orbital periods of the asteroids.
" In fact, 7 will be longer than P, for small orbital apastrons, and the scaling above holds for those separations."," In fact, $T$ will be longer than $P_{orb}$ for small orbital apastrons, and the scaling above holds for those separations."
" Even better, for T »nPqy.ntransits will be detected: confidence that the photometric dips were caused by transits can therefore be high, just as multiple planetary transits enhance confidence in the discovery of planets."," Even better, for $T>n\, P_{\rm orb}$, $n$ transits will be detected; confidence that the photometric dips were caused by transits can therefore be high, just as multiple planetary transits enhance confidence in the discovery of planets."
" For wider separations, however, the probability is reduced by a factor 7/P,,4. For circular orbits, the number of asteroids needed to have a transit detection probability near unity ts It T=P, fora2LAU. then the numbers of asteroids needed to ensure detection for ¢ (0.1. 1. 10, 100) AU are (1.9x107, 1.9x107, 6.0x107, 1.9x 10!),"," For wider separations, however, the probability is reduced by a factor $T/P_{\rm orb}.$ For circular orbits, the number of asteroids needed to have a transit detection probability near unity is If $T=P_{\rm orb}$ for $a=1\, AU,$ then the numbers of asteroids needed to ensure detection for $a=$ (0.1, 1, 10, 100) AU are $1.9\times 10^3$, $1.9\times 10^4$, $6.0\times 10^7$, $1.9\times 10^{11}$ )."
" These numbers are modest enough to suggest that, unless the white dwarf systems are depleted in 100-km class asteroids relative to the solar system,Kepler can discover asteroid transits by monitoring a handful of white dwarfs."," These numbers are modest enough to suggest that, unless the white dwarf systems are depleted in $100$ -km class asteroids relative to the solar system, can discover asteroid transits by monitoring a handful of white dwarfs."
" Furthermore, these numbers above are small enough that a null result would represent ameaningful limit."," Furthermore, these numbers above are small enough that a null result would represent ameaningful limit."
 The characteristics of transit light curves are related to the properties of the, The characteristics of transit light curves are related to the properties of the
Let us choose barveentric Carlesian coordinates CX.Y.Z) such that the plane of the ealaxv is (he X—Y plane and (he X axis is directed. towiudls (he center of the &alaxy.,"Let us choose barycentric Cartesian coordinates $(X,Y,Z)$ such that the plane of the galaxy is the $X-Y$ plane and the $X$ axis is directed towards the center of the galaxy."
 The axis Y is aligned with the direction of the galactic rotation., The axis $Y$ is aligned with the direction of the galactic rotation.
 Let us introduce spherical coordinates (r./.b) such that the unit basis vectors of the two coordinate svstems are related bv the following equations (Binney&Alerrifielcl1998) where (he angular coordinates (/.b) are the galactic longitude and Iatitude respectively. and the direction to the galactic center is given bv the unit vector ey.," Let us introduce spherical coordinates $(r,l,b)$ such that the unit basis vectors of the two coordinate systems are related by the following equations \citep{binney}
 where the angular coordinates $(l,b)$ are the galactic longitude and latitude respectively, and the direction to the galactic center is given by the unit vector $\vec e_X$."
 The geometric direction to a star al coordinates (r./.b) is given bv the unit vector ον=K.," The geometric direction to a star at coordinates $(r,l,b)$ is given by the unit vector $\vec e_r\equiv \vec K$."
 Let us assume that at a eiven epoch fy the orbital velocity of (he Sun around the ealactie center is =(Vy.VyMZ) and the acceleration is 1=(ely.Ay.ely).," Let us assume that at a given epoch $t_0$ the orbital velocity of the Sun around the galactic center is $\vec{V}=(V_X,V_Y,V_Z)$ and the acceleration is $\vec{A}=(A_X,A_Y,A_Z)$."
 Observed direction to the star measured. by a moving observer dillers from its (rue eeonmetric position. NK. because of the stellar aberration that is proportional at each instant of time to the velocity of the observer with respect to the static Irae.," Observed direction to the star measured by a moving observer differs from its true geometric position, $\vec K$, because of the stellar aberration that is proportional at each instant of time to the velocity of the observer with respect to the static frame."
" Since the velocity. vector ol the solar svstem is not constant due to the galactocentrie acceleration. (he aberration angle changes progressively,"," Since the velocity vector of the solar system is not constant due to the galactocentric acceleration, the aberration angle changes progressively."
 Therefore. al another epoch /. the star is seen by a fictitious observer located at the solar barveenter in the direction. . given by where ~x denotes the Euclidean. cross-procduct of two 3-vectors. e is the speed of light in vacuum. /p is the initial epoch. / is the lime of observation. ancl we assume that the contribution of the first and higher order derivatives of acceleration is negligible.," Therefore, at another epoch $t$, the star is seen by a fictitious observer located at the solar barycenter in the direction, $\vec k$, given by where $``\times""$ denotes the Euclidean cross-product of two 3-vectors, $c$ is the speed of light in vacuum, $t_0$ is the initial epoch, $t$ is the time of observation, and we assume that the contribution of the first and higher order derivatives of acceleration is negligible."
 The first two terms in the right-hand side of Eq. (4)), The first two terms in the right-hand side of Eq. \ref{2}) )
 ave constant if one neglects the secular parallax (Binnev&Merrifield1998).. which is irrelevant for quasars and will not be considered in this paper.," are constant if one neglects the secular parallax \citep{binney}, which is irrelevant for quasars and will not be considered in this paper."
 The third term in (he right-hand side of Eq. (4)), The third term in the right-hand side of Eq. \ref{2}) )
 vields the secular aberration that makes all objects bevond the boundaries of the solar svstem to move in the skv with a proper motion, yields the secular aberration that makes all objects beyond the boundaries of the solar system to move in the sky with a proper motion
(Weintraubetal..2000:Bary2003:Biter2007) require high (~ 1000 K) disk gas temperatures.,"\citep{weintraub00, bary03, bitner07} require high $\sim$ 1000 K) disk gas temperatures."
 This warm disk gas may also have been detected through CO emission (Najitàetal.2005:Blake&Boogert.2004:Brittainal...2007) and in fine-structure transitions of after t0nization by EUV/X-ray photons. as further discussed below.," This warm disk gas may also have been detected through CO emission \citep{najita03, blake04, brittain07} and in fine-structure transitions of ] after ionization by EUV/X-ray photons, as further discussed below."
 As a most important by-product of ionization and heating. chemistry is driven across temperature gradients in the disks. especially in UV shielded regions (Aikawa&Herbst.1999).," As a most important by-product of ionization and heating, chemistry is driven across temperature gradients in the disks, especially in UV shielded regions \citep{aikawa99}."
. Furthermore. disk photoevaporation due to X-ray or EUV heating of the inner disk has attracted increasing attention (Alexanderetal..2004:Ercolanoetal..2008;Gorti 2009).," Furthermore, disk photoevaporation due to X-ray or EUV heating of the inner disk has attracted increasing attention \citep{alexander04, ercolano08, gorti09}."
. Short-wavelength disk irradiation is thus of central interest for our understanding of the ionization of cireumstellar disks. their heating and chemical processing. disk instabilities. and photoevaporation and therefore the long-term evolution of disks.," Short-wavelength disk irradiation is thus of central interest for our understanding of the ionization of circumstellar disks, their heating and chemical processing, disk instabilities, and photoevaporation and therefore the long-term evolution of disks."
 All these mechanisms obviously affect the process of planet formation., All these mechanisms obviously affect the process of planet formation.
 Glassgoldetal.(2007) proposed that the mid-infrared u]] fine-structure transition at. 12.81jmmn is a tracer of gas requiring X-ray irradiation of the disk., \citet{glassgold07} proposed that the mid-infrared ] fine-structure transition at $\mu$ m is a tracer of gas requiring X-ray irradiation of the disk.
 Because the first ionization potential of Ne is high (21.6 eV). its photoionization indeed requires EUV or X-ray photons.," Because the first ionization potential of Ne is high (21.6 eV), its photoionization indeed requires EUV or X-ray photons."
" In the X-ray ionization model of Glassgoldetal.(2007) Ne is ionized by K-shell absorption. requiring photons with energies of at least 0.9 keV. As the same X-rays also heat the upper layers of the gas disk to several 1000 K (Glassgoldetal..2004).. u]] fine-structure transitions with excitation temperatures of =1000 K are produced over a scale height of warm gas of 10!""—107? em."," In the X-ray ionization model of \citet{glassgold07}
 Ne is ionized by K-shell absorption, requiring photons with energies of at least 0.9 keV. As the same X-rays also heat the upper layers of the gas disk to several 1000 K \citep{glassgold04}, ] fine-structure transitions with excitation temperatures of $\approx 1000$ K are produced over a scale height of warm gas of $10^{19}-10^{20}$ $^{-2}$."
 An alternative model (Hollenbach&Gorti2009.. see also Gorti&Hollenbach 2008)) proposes tonization of the disk surface by EUV radiation. producing an Hu--like highly-ionized region in which the |Neu]| line is formed.," An alternative model \citealt{hollenbach09}, see also \citealt{gorti08}) ) proposes ionization of the disk surface by EUV radiation, producing an -like highly-ionized region in which the ] line is formed."
 The [Neu]] transition thus appears to be an ideal tracer both of disk gas and of the environmental impact of high-energy stellar radiation. u], The ] transition thus appears to be an ideal tracer both of disk gas and of the environmental impact of high-energy stellar radiation. ]
] emissivities were also calculated for shocks in molecular clouds. e.g. such às those forming when a protostellar jet rams into the surrounding molecular gas (e.g.. Hollenbach&McKee 1989)).," emissivities were also calculated for shocks in molecular clouds, e.g. such as those forming when a protostellar jet rams into the surrounding molecular gas (e.g., \citealt{hollenbach89}) )."
 The Infrared Space Observatory (ISO) indeed detected the u]|] line feature in the T Tau system (vandenAnckeretal..1999); the authors attributed this line to shocks in the outflows of the T Tau system., The Infrared Space Observatory (ISO) indeed detected the ] line feature in the T Tau system \citep{ancker99}; the authors attributed this line to shocks in the outflows of the T Tau system.
 The u]] line has also been seen in Herbig-Haro objects (Neufeld 20006)., The ] line has also been seen in Herbig-Haro objects \citep{neufeld06}.
. The advent of the henceforth: Werneretal. 2004)) has renewed interest in the [Nen]] feature from pre-main sequence stars as theoretical calculations of X-ray irradiated circumstellar disks (Glassgoldetal.2007) predicted easy detection by the Infrared Spectrometer (IRS: Houcketal. 2004))., The advent of the henceforth; \citealt{werner04}) ) has renewed interest in the ] feature from pre-main sequence stars as theoretical calculations of X-ray irradiated circumstellar disks \citep{glassgold07} predicted easy detection by the Infrared Spectrometer (IRS; \citealt{houck04}) ).
 Follow-up calculations have confirmed and deepened these initial predictions (Meijerinketal...2008)... also including EUV irradiation (Hollenbach&Gorti.2009).," Follow-up calculations have confirmed and deepened these initial predictions \citep{meijerink08}, also including EUV irradiation \citep{hollenbach09}."
.. First. successful detections of the [Nen]] line by the IRS were reported by Pascuceietal.(2007).. Espaillatetal.(2007).. Lahuisetal.(2007).. and Ratzkaetal.(2007) from a variety of pre-main sequence objects with circumstellar disks and partly also with protostellar envelopes.," First successful detections of the ] line by the IRS were reported by \citet{pascucci07}, \citet{espaillat07}, \citet{lahuis07}, and \citet{ratzka07} from a variety of pre-main sequence objects with circumstellar disks and partly also with protostellar envelopes."
 Initial attempts were made to relate the observed [Net]] fluxes to stellar or disk properties. but the mostly small samples yielded ambiguous results.," Initial attempts were made to relate the observed ] fluxes to stellar or disk properties, but the mostly small samples yielded ambiguous results."
 Pascuecietal.(2007) claimed a correlation between the luminosity in the [Neiu]] line. Zjnejj. and the stellar X-ray luminosity.Ly. but the sample contained only four detections distributed in Lx over a mere 0.2 dex. which corresponds to the usual range of variability of stellar coronal X-ray emission.," \citet{pascucci07} claimed a correlation between the luminosity in the ] line, $L_{\rm [Ne~II]}$, and the stellar X-ray luminosity,$L_{\rm X}$, but the sample contained only four detections distributed in $L_{\rm X}$ over a mere 0.2 dex, which corresponds to the usual range of variability of stellar coronal X-ray emission."
 A correlation with the stellar mass aceretion rate. Mj. was not found.," A correlation with the stellar mass accretion rate, $\dot{M}_{\rm acc}$, was not found."
" Conversely. Espaillatetal.(2007). suggested Linej| to. be correlated with M, but not with Ly. but again. the data sampleand the dynamic range in Ly were small (=0.8 dex)."," Conversely, \citet{espaillat07} suggested $L_{\rm [Ne~II]}$ to be correlated with $\dot{M}_{\rm acc}$ but not with $L_{\rm X}$ but again, the data sampleand the dynamic range in $L_{\rm X}$ were small $\approx 0.8$ dex)."
 Finally. Lahuisetal.(2007) considered u]] line production both as à consequence of X-ray or EUV irradiation of disks and of shock formation on the disks themselves.," Finally, \citet{lahuis07} considered ] line production both as a consequence of X-ray or EUV irradiation of disks and of shock formation on the disks themselves."
 They noticed that the measured Linojj correspond well to predictions made by Glassgoldetal.(2007)., They noticed that the measured $L_{\rm [Ne~II]}$ correspond well to predictions made by \citet{glassgold07}.
. Ground-based observations of the [Neu]] line allow for much higher spectral resolving power. uncovering the kinematics of the emitting regions.," Ground-based observations of the ] line allow for much higher spectral resolving power, uncovering the kinematics of the emitting regions."
 Herezegetal.(2007) observed a relatively narrow profile m the near-pole-on system TW Hya. interpreting the emission as coming either from the disk surface (1.9... from a region where the gas is gravitationally bound). or a slow photoevaporative flow from à disk region that allows for escaping gas flows given a sufficiently high gas temperature. u]," \citet{herczeg07} observed a relatively narrow profile in the near-pole-on system TW Hya, interpreting the emission as coming either from the disk surface (i.e., from a region where the gas is gravitationally bound), or a slow photoevaporative flow from a disk region that allows for escaping gas flows given a sufficiently high gas temperature. ]"
] emission from such flows was indeed modeled by Alexander(2008).. their results being consistent with the observations.," emission from such flows was indeed modeled by \citet{alexander08}, their results being consistent with the observations."
 Additional support came from high spectral resolution observations of transition disks: three objects of this class showed line profiles and blue-shifts consistent with those predicted from photoevaporative flows (Pascucet&Sterzik.2009)., Additional support came from high spectral resolution observations of transition disks; three objects of this class showed line profiles and blue-shifts consistent with those predicted from photoevaporative flows \citep{pascucci09}.
.. Further. ground-based observations of AA Tau and GM Aur with the TEXES instrument clearly show that most of their [Nemn]|] emission is consistent with a disk origin although the lower fluxes. when compared withSpitzer results. suggest that there is also an extended component (not recovered by the narrow slit of TEXES).," Further ground-based observations of AA Tau and GM Aur with the TEXES instrument clearly show that most of their ] emission is consistent with a disk origin although the lower fluxes, when compared with results, suggest that there is also an extended component (not recovered by the narrow slit of TEXES)."
 Alternatively. the sources may be time-variable. or the line ts spectrally unusually broad (Najitaetal...2009).," Alternatively, the sources may be time-variable, or the line is spectrally unusually broad \citep{najita09}."
 A rather unexpected twist came with the observation of the T Taurt system with the VLT at high spectral resolution (vanBoekeletal..2009)., A rather unexpected twist came with the observation of the T Tauri system with the VLT at high spectral resolution \citep{boekel09}.
. Here. the | emission region was clearly extended (by several areseconds) and showed line broadening and line shifts (by up to 126 km s!) compatible with the jets of the T Tauri system. previously observed in a similar fashion in u]] and [Otr]] by Bóhm&Solf (1994)..," Here, the ] emission region was clearly extended (by several arcseconds) and showed line broadening and line shifts (by up to 126 km $^{-1}$ ) compatible with the jets of the T Tauri system, previously observed in a similar fashion in ] and ] by \citet{boehm94}. ."
 The interpretation favored shock-induced |Neu]| formation.but X-ray irradiation by the star and consequent ionization of the jet material remains a viable option (vanBoekelet 2009)..," The interpretation favored shock-induced ] formation,but X-ray irradiation by the star and consequent ionization of the jet material remains a viable option \citep{boekel09}. ."
 This scenario (stellar X-rays tonizing the jet. and [Nen]] emission forming in the jet itself) has recently been suggested on theoretical grounds as well (Shangetal.. 2010)..," This scenario (stellar X-rays ionizing the jet, and ] emission forming in the jet itself) has recently been suggested on theoretical grounds as well \citep{shang10}. ."
sstate up to the sstate (following the terminology of P06). in order to compare the model parameter evolution during the spectral transition.,"state up to the state (following the terminology of P06), in order to compare the model parameter evolution during the spectral transition."
 We remind the reader that Z sources. unlike atolls. have never been detected in the oor state. possibly because they never reach low enough accretion rate levels.," We remind the reader that Z sources, unlike atolls, have never been detected in the or state, possibly because they never reach low enough accretion rate levels."
 Thus the sstate and the spectral transition to the sstate is a complete pattern of the spectral evolution for à Z source., Thus the state and the spectral transition to the state is a complete pattern of the spectral evolution for a Z source.
 We briefly remind the reader about the main characteristics of the mmodel (?):: the total emerging spectrum is given by where the first and second terms of the right-hand side represent the seed BB-like photon spectrum and its convolution with the system Green's. function (Comptonized spectrum). respectively.," We briefly remind the reader about the main characteristics of the model \citep{farinelli08}: the total emerging spectrum is given by where the first and second terms of the right-hand side represent the seed BB-like photon spectrum and its convolution with the system Green's function (Comptonized spectrum), respectively."
 The factor 1/(1|4) is the fraction of the seed photon radiation directly seen by the Earth observer. whereas the factor 1/(1|A) is the fraction of the seed photon radiation up-scattered by the Compton cloud.," The factor $1/(1+A)$ is the fraction of the seed photon radiation directly seen by the Earth observer, whereas the factor $A/(1+A)$ is the fraction of the seed photon radiation up-scattered by the Compton cloud."
" Free parameters of the model are the BB seed photons color temperature. kT, and normalization Cy. the plasma temperature. KT... the logarithm of the illuminating factor A. log(.A)."," Free parameters of the model are the BB seed photons color temperature, $_{s}$ and normalization $C_N$, the plasma temperature, $_{e}$, the logarithm of the illuminating factor A, $\log(A)$."
 Moreover. the information on the efficiency of the Comptonization is given by two parameters. o and à.," Moreover, the information on the efficiency of the Comptonization is given by two parameters, $\alpha$ and $\delta$."
 The parameter o (photon index T=a| 1) indicates an overall Comptonization efficiency related to an observable quantity in the photon spectrum of the data., The parameter $\alpha$ (photon index $\Gamma=\alpha+1$ ) indicates an overall Comptonization efficiency related to an observable quantity in the photon spectrum of the data.
 The lower the © parameter (spectrum extending to higher energies) the higher the efficiency. re. the higher energy transfer from hot electrons to soft seed photons.," The lower the $\alpha$ parameter (spectrum extending to higher energies) the higher the efficiency, i.e. the higher energy transfer from hot electrons to soft seed photons."
 The à-parameter provides information about the efficiency of bulk BC with respect that to TC., The $\delta$ -parameter provides information about the efficiency of bulk BC with respect that to TC.
 Both o and à are closely related (see Fig., Both $\alpha$ and $\delta$ are closely related (see Fig.
 | in F08) and only in the case of a clear cut-off in the hard X-ray tail data is it possible to have a precise estimate of the efficiency of both Each continuum of the five spectra of ((two HB and three NB) has been fitted using the sum of two ccomponents., 1 in F08) and only in the case of a clear cut-off in the hard X-ray tail data is it possible to have a precise estimate of the efficiency of both Each continuum of the five spectra of (two HB and three NB) has been fitted using the sum of two components.
 In Tab., In Tab.
 | we report the results of our spectral analysis., \ref{tab_fit} we report the results of our spectral analysis.
 In all cases. the first ccomponent [pure TC (6=0)] provides a significant contribution in the total emerging spectrum.," In all cases, the first component [pure TC $\delta$ =0)] provides a significant contribution in the total emerging spectrum."
 Most of the source energetic budget (60-70 %)) is. in fact. determined by this component which is described (a well-known feature of Z-class LMXBs) by Comptonization of ~ 0.3 keV photons off cool electrons (&T;.~ 3-5 keV) of high optical depth environment (7~ 10 or 5 depending," Most of the source energetic budget (60-70 ) is, in fact, determined by this component which is described (a well-known feature of Z-class LMXBs) by Comptonization of $\sim$ 0.3 keV photons off cool electrons $\kte \sim$ 3-5 keV) of high optical depth environment $\tau \sim$ 10 or 5 depending"
horizontal branches it becomes clear that horizontal branch morphology cannot solve the 115 problem.,horizontal branches it becomes clear that horizontal branch morphology cannot solve the $\beta$ problem.
 Likewise. inclusion of blue stragglers will not help. even if (here was evidence lor a strong modulation of blue straggler lrequency. with metallicity. which there is not 2005).," Likewise, inclusion of blue stragglers will not help, even if there was evidence for a strong modulation of blue straggler frequency with metallicity, which there is not \citep{sand05}."
. Emission fill-in of the Dalmer lines due to hydrogen recombination lines [rom external nebulae is probably ruled out. except for case like (he planetary nebula in M5.," Emission fill-in of the Balmer lines due to hydrogen recombination lines from external nebulae is probably ruled out, except for case like the planetary nebula in M5."
 Stellar activity in individual cluster stars seems to be the only surviving mechanism that has good evidence., Stellar activity in individual cluster stars seems to be the only surviving mechanism that has good evidence.
 llowever. even being generous. the cool tail of M dwarls do not appear to be able to generate enough [ιν to cause the modulation in Io needed.," However, even being generous, the cool tail of M dwarfs do not appear to be able to generate enough flux to cause the modulation in $\beta$ needed."
 The remaining stellar source is fIaring in AI giants., The remaining stellar source is flaring in M giants.
 These stars are bright enough. and inherently stochastic in nature. which seems to fit the observations of clusters that scatter to low IL? rather randomly.," These stars are bright enough, and inherently stochastic in nature, which seems to fit the observations of clusters that scatter to low $\beta$ rather randomly."
 Finally. it remains a long-shot possibility that abundance ratios in O or the noble gases can cause isochrone temperature cdrilis severe enough to affect the IL? problem.," Finally, it remains a long-shot possibility that abundance ratios in O or the noble gases can cause isochrone temperature drifts severe enough to affect the $\beta$ problem."
 The authors gratefully. acknowledge support [rom National Science Foundation grants AST-0307487 and AST-0346341., The authors gratefully acknowledge support from National Science Foundation grants AST-0307487 and AST-0346347.
We would like to stress that. since the initialfinal lnass relationship is almost flat in the low mass regime and the massmain sequence lifetime relationship is verv steep. any attempt to derive ages of individual field white dwarts from the position of low mass white dawarfs in the colormagnitude diagram is subject to potentially large uncertainties. since anv small error im the determunuation of the white ciwarf mass translates iuto a huge relative error on its total age.,"We would like to stress that, since the initial–final mass relationship is almost flat in the low mass regime and the mass–main sequence lifetime relationship is very steep, any attempt to derive ages of individual field white dwarfs from the position of low mass white dwarfs in the color–magnitude diagram is subject to potentially large uncertainties, since any small error in the determination of the white dwarf mass translates into a huge relative error on its total age."
" It should also be pointed out that there is avery colmmon tendency to associate bright white dwarfs with vouug stars which is inaccurate since they cau be either bright massive white dwarfs and. imdoeed. in this case their total age is simall or bright low mass white dwarts with a low mass progenitor. which las a large main sequence lifetime. and therefore the reverse is true Pinto et al 1901, Iseru et al 1999)."," It should also be pointed out that there is a very common tendency to associate bright white dwarfs with young stars, which is inaccurate since they can be either bright massive white dwarfs — and, indeed, in this case their total age is small — or bright low mass white dwarfs with a low mass progenitor, which has a large main sequence lifetime, and therefore the reverse is true }az--Pinto et al 1994, Isern et al 1999)."
" As it cau be seen in figure 5. in all the colormagnitude jagranus the isochrones show a pronounced turn-off at their cimuuer cud which. as the age of the isochrone oeicroases, it is located at inercasinely larger maguitudes."," As it can be seen in figure 5, in all the color–magnitude diagrams the isochrones show a pronounced turn-off at their dimmer end which, as the age of the isochrone increases, it is located at increasingly larger magnitudes."
 The presence of this turnoff is due to the coutribution of the most massive white dwarfs. while the upper portion of the isochrones closely resembles the cooling track of a ~ 0.51 AZ; object.," The presence of this turn–off is due to the contribution of the most massive white dwarfs, while the upper portion of the isochrones closely resembles the cooling track of a $\sim$ 0.54 $M_{\sun}$ object."
 The shape of the turn-off is modulated. in the JWN colours (for ages lavecr than f=78 Cyr). by the intrinsic turn to the blue of the inclivicual cooling tracks clearly sccu in figure |.," The shape of the turn-off is modulated, in the $J-K$ colours (for ages larger than $t=7-8$ Gyr), by the intrinsic turn to the blue of the individual cooling tracks clearly seen in figure 4."
 For the Wo7 and WoR colors the turn to the blue of the individual cooling sequences begins to contribute siguificantly to the isochronues for ages f£—13.LL C., For the $V-I$ and $V-R$ colors the turn to the blue of the individual cooling sequences begins to contribute significantly to the isochrones for ages $t=13-14$ Gyr.
 Iu this work we have computed the cooling sequences of very cool DA white dwarfs with C/O cores., In this work we have computed the cooling sequences of very cool DA white dwarfs with C/O cores.
 These cooling sequences include the most accurate plivsical description of the thermodvuamic quantities for both the core and the envelope and have been computed with a well tested. selfconsistent evolutionary code.," These cooling sequences include the most accurate physical description of the thermodynamic quantities for both the core and the envelope and have been computed with a well tested, self–consistent evolutionary code."
 We have computed as well the release of latent heat using the most up to date physical iuputs for the liquidsolid phase transition and we have found that for the most recent equations of state of the degenerate core the release of latent heat amounts to (Q.SEkp per particle., We have computed as well the release of latent heat using the most up to date physical inputs for the liquid–solid phase transition and we have found that for the most recent equations of state of the degenerate core the release of latent heat amounts to $\sim 0.8 k_{\rm B}T$ per particle.
" Additionally, the release of eravitational energy associated to pliase separation durus crvstallization has been also properly taken iuto account."," Additionally, the release of gravitational energy associated to phase separation during crystallization has been also properly taken into account."
 Color indices for several baudpasses have also been computed from the most recent svuthetie spectra., Color indices for several bandpasses have also been computed from the most recent synthetic spectra.
 Our major fiudiues can be stummarized as follows., Our major findings can be summarized as follows.
 Firstly. we have fouud that the most recent cooling sequences of Tauscu (1999) considerably uuderestinmate the cooling age of C/O white dwarts with hydrogen dominatedLB. atmosplieres even when chemical fractionation is neglected.," Firstly, we have found that the most recent cooling sequences of Hansen (1999) considerably underestimate the cooling age of C/O white dwarfs with hydrogen dominated atmospheres even when chemical fractionation is neglected."
 We have traced back the possible source of discrepancy and we lave found that at a given core teniperature the 7;L velationship of ILuseu (1999) differs by roughly with respect to other evolutionary caleulatious computed with the same physical iuputs., We have traced back the possible source of discrepancy and we have found that at a given core temperature the $T_{\rm c}-L$ relationship of Hansen (1999) differs by roughly with respect to other evolutionary calculations computed with the same physical inputs.
 This is true not only when the comparison is done with the cooling sequences reported here but iso with the cooling sequences by other authors (uamely Althaus Beuvenuto 1998)., This is true not only when the comparison is done with the cooling sequences reported here but also with the cooling sequences by other authors (namely Althaus Benvenuto 1998).
 The ultimate reason for this discrepancy remains unidentified although we have explored several possible causes without much success., The ultimate reason for this discrepancy remains unidentified although we have explored several possible causes without much success.
 Secondly. we have also found that a conservative lower luit to the acctuulated time delay. introduced by the release of eravitational energy associated to pliase separation is roughly at log(L/EL:)=LS. in good agrecimeut with the results of Isern et al (2000).," Secondly, we have also found that a conservative lower limit to the accumulated time delay introduced by the release of gravitational energy associated to phase separation is roughly at $\log(L/L_{\sun})= -4.5$, in good agreement with the results of Isern et al (2000)."
 We have transformed our cooling sequences from the log(L/L: log(T.g) pluie into various colormaguitude diagrams. and found that intrinsically faint DA white dwarts have a pronounced turnoff iu the infrared colors. and therefore are bluer than previously thought. in good aerecinent with the results of Hausen (1998. 1999).," We have transformed our cooling sequences from the $\log(L/
L_{\sun})$ $\log(T_{\rm eff})$ plane into various color–magnitude diagrams, and found that intrinsically faint DA white dwarfs have a pronounced turn–off in the infrared colors, and therefore are bluer than previously thought, in good agreement with the results of Hansen (1998, 1999)."
 Cooling isochrones taking iuto account the evolutionary time of the white dwarf progenitors suitable for the analysis of white dwarfs m stellar clusters have been produced as well., Cooling isochrones taking into account the evolutionary time of the white dwarf progenitors — suitable for the analysis of white dwarfs in stellar clusters — have been produced as well.
 They show in all colors a turnoff at the faiter eud. due to the contribution of the more massive objects: this turioff is modulated by the iutrimsic turn to the blue of the individual cooling tracks.," They show in all colors a turn–off at the fainter end, due to the contribution of the more massive objects; this turn–off is modulated by the intrinsic turn to the blue of the individual cooling tracks."
 Finally we would like to stress the inportauce of having reliable models of the evolution of white dwarts aud. thus. given the substantial differcuces in the inferred ages found bv different authors. and the complexity of the evolutionary codes needed to compute realistic cooling sequences. niore iudependoent calculations are highly desixalje.," Finally we would like to stress the importance of having reliable models of the evolution of white dwarfs and, thus, given the substantial differences in the inferred ages found by different authors, and the complexity of the evolutionary codes needed to compute realistic cooling sequences, more independent calculations are highly desirable."
 Q.5e1m This work has beeu supported by the DGES erauts PD970983C0302 and PBOT0983C03 by the NSF eraut AST 9731138 and by the CIRIT.," 0.5cm This work has been supported by the DGES grants PB97–0983–C03–02 and PB97--0983--C03--03, by the NSF grant AST 97–31438 and by the CIRIT."
 We sincerely thank our referee. D. IEuisen. for very valuable conmunents aud criticisin which have considerably iiproved the original manuscript.," We sincerely thank our referee, B. Hansen, for very valuable comments and criticism which have considerably improved the original manuscript."
 We also want to thank P. Dergerou for kindly providing us with model atinospheres., We also want to thank P. Bergeron for kindly providing us with model atmospheres.
 One of us. EGB. also acknowledges the support received frou Sun MicroSystems under the Academic Equipment Craut AEG7821990325SP.," One of us, EGB, also acknowledges the support received from Sun MicroSystems under the Academic Equipment Grant AEG–7824–990325–SP."
times could be easily reassigned. but because of the high HRC background rate and a telemetry limit of 184 events/sec. onlyvalid events. which make up approximately ol the/otal events in these observations. are recorded in data received on the ground.,"times could be easily reassigned, but because of the high HRC background rate and a telemetry limit of 184 events/sec, only events, which make up approximately of the events in these observations, are recorded in data received on the ground."
 We have reassigned the time of every valid event to the preceding valid event during ground processing: the event Gimmes will (hen be correct about. INutu]Nor: ol Uae me., We have reassigned the time of every valid event to the preceding valid event during ground processing; the event times will then be correct about $N_{valid}/N_{total}$ of the time.
 One can place an upper limit of 6/ on the (nme errors associated will this process. however. simply by excluding events with a time-shift of more than 9/: the average timing error of such events will. of course. be less Chan 6/.," One can place an upper limit of $\delta t$ on the time errors associated with this process, however, simply by excluding events with a time-shift of more than $\delta t$; the average timing error of such events will, of course, be less than $\delta t$."
 However. if 0/ is too small. too many events get excluded to perform sensitive timing studies.," However, if $\delta t$ is too small, too many events get excluded to perform sensitive timing studies."
 We have adopted a o/ of 2 ms. retaining of the original counts in the corrected data set.," We have adopted a $\delta t$ of 2 ms, retaining of the original counts in the corrected data set."
 Count rates were derived [or each observation segment [rom event lists filtered to exclude times of high background and telemetry saturation. when the (elemeterecl valid event rate exceeded 184 count 1. and to exclude high pulse-height events that arise entirely from backeround.," Count rates were derived for each observation segment from event lists filtered to exclude times of high background and telemetry saturation, when the telemetered valid event rate exceeded 184 count $^{-1}$, and to exclude high pulse-height events that arise entirely from background."
 These rates corresponding to OU order only are listed in Table 1.., These rates corresponding to 0th order only are listed in Table \ref{t:obs}.
 The count rates for the different observation segments are statistically consistent. (hough the 2000 March observation (113) rate of 0.2195£0.0020 count + lies aand 1.96 above the rate for the combined 2001 October series rate of 0.2155220.0008.," The count rates for the different observation segments are statistically consistent, though the 2000 March observation (113) rate of $0.2195\pm 0.0020$ count $^{-1}$ lies and $1.9\sigma$ above the rate for the combined 2001 October series rate of $0.2155\pm
0.0008$."
 Such deviations will be obtained by chance about oof the time when the count rates do not vary: il is thus more likely attributable to quantum efficiency (QE) variations in the detector ou small scales., Such deviations will be obtained by chance about of the time when the count rates do not vary; it is thus more likely attributable to quantum efficiency (QE) variations in the detector on small scales.
 Indeed. Pons et ((2001) note two ROSAT URI observations obtained 3 vr apart that are consistent toLv.," Indeed, Pons et (2001) note two ROSAT HRI observations obtained 3 yr apart that are consistent to."
.. The fluxes corresponding to ILAI and PSPC count rates are higher than that obtained from our Chandra data by aand30%.. respectively.," The fluxes corresponding to HRI and PSPC count rates are higher than that obtained from our Chandra data by and, respectively."
 As remarked by Durwitz οἱ ((2001). these differences are likely attributable to absolute calibration uncertainties.," As remarked by Burwitz et (2001), these differences are likely attributable to absolute calibration uncertainties."
 We have also examined the ASCA SIS observation described Pons et ((2001) ancl find fluxes 30-40 llower than obtained by Chandra for the 20-30 rrange. but in agreement with Chandra shortward of 20A.," We have also examined the ASCA SIS observation described Pons et (2001) and find fluxes 30-40 lower than obtained by Chandra for the 20-30 range, but in agreement with Chandra shortward of 20."
 The EUVE DS count rate of Pons et ((2001) is also consistent with the Chandra observation within allowed uncertainties., The EUVE DS count rate of Pons et (2001) is also consistent with the Chandra observation within allowed uncertainties.
 Our search for pulsations used Chree different techniques. none of which found any evidence lor significant. variability.," Our search for pulsations used three different techniques, none of which found any evidence for significant variability."
 In contrast to Ransom et ((2002). we did not include a deceleration term: this will be discussed below.," In contrast to Ransom et (2002), we did not include a deceleration term; this will be discussed below."
"the polar cap, the drift motion of plasma particles, and their realistic energy distribution function are taken into account.","the polar cap, the drift motion of plasma particles, and their realistic energy distribution function are taken into account."
 It gives us the first opportunity to provide the quantitative comparison of the theoretical predictions with observational data., It gives us the first opportunity to provide the quantitative comparison of the theoretical predictions with observational data.
" Using numerical integration we can now model the mean profiles of radio pulsars and, hence, evaluate the physical parameters of the plasma flowing in the pulsar magnetosphere."," Using numerical integration we can now model the mean profiles of radio pulsars and, hence, evaluate the physical parameters of the plasma flowing in the pulsar magnetosphere."
 It is necessary to stress that the detailed discussion of the morphological properties of mean profiles is beyond the scope of our considerationpapers., It is necessary to stress that the detailed discussion of the morphological properties of mean profiles is beyond the scope of our consideration.
 The goal of this paper is in quantitative analysis of the propagation effects on the polarization characteristics of radio pulsars., The goal of this paper is in quantitative analysis of the propagation effects on the polarization characteristics of radio pulsars.
" In particular, we try to determine how the plasma parameters affect the S-shape of the position angle swing and the properties of the mean profile."," In particular, we try to determine how the plasma parameters affect the $S$ -shape of the position angle swing and the properties of the mean profile."
" At first, let us discuss the results obtained by numerical integration of equations (70))-(71)) for ""ordinary"" pulsar (its parameters are given in Table 2))."," At first, let us discuss the results obtained by numerical integration of equations \ref{t1}) \ref{t2}) ) for ”ordinary” pulsar (its parameters are given in Table \ref{table1}) )."
 Everywhere below the dashed curves on the intensity panel show the intensity profile without any absorption., Everywhere below the dashed curves on the intensity panel show the intensity profile without any absorption.
" As was already stressed, in this paper for simplicity we suppose that it repeats the particle number density profile shown in Fig. 6.."," As was already stressed, in this paper for simplicity we suppose that it repeats the particle number density profile shown in Fig. \ref{fig0}."
" If the dashed curve is not shown, then the absorption is fatal and only the original intensity (which is normalized to 100 in its maximum) is shown."," If the dashed curve is not shown, then the absorption is fatal and only the original intensity (which is normalized to 100 in its maximum) is shown."
 The dashed curves on panels show the prediction of the RVM-model (1))., The dashed curves on panels show the prediction of the RVM-model \ref{p.a.}) ).
" Finally, pulsar phase $ is measured in degrees everywhere below."," Finally, pulsar phase $\phi$ is measured in degrees everywhere below."
 On Fig., On Fig.
" 11 we show the intensity Το. (60)) (left panel) and the swing (right panel) for extraordinary X-mode as a function of the pulsar phase ó for ""non-rotating dipole"" without the wind component (model A); the drift effects are neglected as well.", \ref{figres1} we show the intensity $I_{\infty}$ \ref{intens}) ) (left panel) and the swing (right panel) for extraordinary X-mode as a function of the pulsar phase $\phi$ for ”non-rotating dipole” without the wind component (model A); the drift effects are neglected as well.
 The circular polarization degree does not exceed one percent here and that is why this curve is not presented in this picture., The circular polarization degree does not exceed one percent here and that is why this curve is not presented in this picture.
 It results from approximately constant Gp along the ray., It results from approximately constant $\beta_{B}$ along the ray.
" For this reason, as we see, the curve is nicelyfitting by the RVM model."," For this reason, as we see, the curve is nicelyfitting by the RVM model."
" Further, the upper solid line corresponds to fo— 0.25, and the lower one corresponds to fo—0.0025."," Further, the upper solid line corresponds to $f_0 = 0.25$ , and the lower one corresponds to $f_0 = 0.0025$."
 The lower intensity curve shows that for the very small core in the number density fo=0.0025 (ri/Ro= 0.05) the absorption is fatal and the emission cannot escape from the magnetosphere., The lower intensity curve shows that for the very small core in the number density $f_0 = 0.0025$ $r_{\perp}/R_0 = 0.05$ ) the absorption is fatal and the emission cannot escape from the magnetosphere.
" As was already stressed, this property can be easily explained."," As was already stressed, this property can be easily explained."
" Indeed, for fo«1 the rarefied region of the ""hollow cone"" is actually absent, and the rays pass the cyclotron resonance in the region of rather dense plasma."," Indeed, for $f_0 \ll 1$ the rarefied region of the ”hollow cone” is actually absent, and the rays pass the cyclotron resonance in the region of rather dense plasma."
" On the other hand, for fo~1 the number density in the region of the cyclotron resonance is low enough for rays to escape the magnetosphere without strong absorption."," On the other hand, for $f_0 \approx 1$ the number density in the region of the cyclotron resonance is low enough for rays to escape the magnetosphere without strong absorption."
 On Fig., On Fig.
" 12 we show the swing (lower panels), the intensity [ος (solid lines on the top panels), and the Stokes parameter V (dotted lines) as a function of the pulsar phase @ for extraordinary wave for magnetic field model C and for various multiplicity parameter A."," \ref{figres2} we show the swing (lower panels), the intensity $I_{\infty}$ (solid lines on the top panels), and the Stokes parameter $V$ (dotted lines) as a function of the pulsar phase $\phi$ for extraordinary wave for magnetic field model C and for various multiplicity parameter $\lambda$."
 It is obvious that the absorption increases with increasing A., It is obvious that the absorption increases with increasing $\lambda$.
" In most cases, the trailing part of the mean profile is absorbed (see Dyks et al."," In most cases, the trailing part of the mean profile is absorbed (see Dyks et al."
 2010 as well)., 2010 as well).
" Besides, as one can see, the curves differ significantly from the RVM one (dashed lines) as )growing."," Besides, as one can see, the curves differ significantly from the RVM one (dashed lines) as growing."
Thispropertycanbeeasilyunderstoodaswell., This property can be easily understood as well.
"Indeed, astheescapera (7)) increases as \?/°, for large enough A the polarization properties of the outgoing waves are to be formed in the vicinity of the light cylinder, ie., in the region with quasi-homogeneous magnetic field (see Fig. 5))."," Indeed, as the escape radius $r_{\rm esc}$ \ref{resc}) ) increases as $\lambda^{2/5}$, for large enough $\lambda$ the polarization properties of the outgoing waves are to be formed in the vicinity of the light cylinder, i.e., in the region with quasi-homogeneous magnetic field (see Fig. \ref{figmagnfield}) )."
" Further, as was already mentioned, the drift effect causes the curve to be shifted to the trailing part of the mean profile."," Further, as was already mentioned, the drift effect causes the curve to be shifted to the trailing part of the mean profile."
 It is necessary to stress that the opposite shift is to take place if we neglect the drift effect on the dielectric tensor (Andrianov Beskin 2010; Wang et al., It is necessary to stress that the opposite shift is to take place if we neglect the drift effect on the dielectric tensor (Andrianov Beskin 2010; Wang et al.
 2010)., 2010).
" One can note that for highvalues of multiplicity parameter (i.e., for full absorption of the trailing part of the mean pulse) the observer will detect approximately constant Finally, as one can see from Eqn. (78)),"," One can note that for highvalues of multiplicity parameter (i.e., for full absorption of the trailing part of the mean pulse) the observer will detect approximately constant Finally, as one can see from Eqn. \ref{V}) ),"
" the maximum of circular polarization is also shifted to the trailing side, as larger deviations from the S-shape produce larger circular polarization."," the maximum of circular polarization is also shifted to the trailing side, as larger deviations from the $S$ -shape produce larger circular polarization."
 It is not visible on this picture because the trailing side of the beam is absorbed., It is not visible on this picture because the trailing side of the beam is absorbed.
 On Fig., On Fig.
" 13 we show the same dependences for various Lorentz-factors of outgoing plasma yo=10, 50, and 300."," \ref{figres3} we show the same dependences for various Lorentz-factors of outgoing plasma $\gamma_0 = 10$, $50$ , and $300$ ."
" As the escape radius resc (7)) decreases as y~°/°, the largest shift of the p.a. curve takes place for small yo=10."," As the escape radius $r_{\rm esc}$ \ref{resc}) ) decreases as $\gamma^{-6/5}$, the largest shift of the $p.a.$ curve takes place for small $\gamma_0 = 10$."
" Finally, on Fig."," Finally, on Fig."
" 14. one can see the same dependences for various wave frequencies v=0.03, 0.2, and 0.5 GHz."," \ref{figres4} one can see the same dependences for various wave frequencies $\nu= 0.03$, $0.2$, and $0.5$ GHz."
" As Toscxν ?7/5, the largest shift of the curve takes place for small frequencies."," As $r_{\rm esc} \propto \nu^{-2/5}$ , the largest shift of the curve takes place for small frequencies."
" One can note that the full investigation of frequency dependence of mean profiles of radio pulsars is more complicate and must include, e.g., the detailed analysis of frequency dependence of the emission radius rem (BGI)."," One can note that the full investigation of frequency dependence of mean profiles of radio pulsars is more complicate and must include, e.g., the detailed analysis of frequency dependence of the emission radius $r_{\rm em}$ (BGI)."
 'This is beyond the scope of the article., This is beyond the scope of the article.
" As was demonstrated above, in general the sign of the circular polarization remains constant for a given mode."," As was demonstrated above, in general the sign of the circular polarization remains constant for a given mode."
 But under certain conditions the change ofthe V sign in the same mode may occur., But under certain conditions the change ofthe $V$ sign in the same mode may occur.
 It can take placewhen we cross the directivity pattern inthe veryvicinity ofthe magneticaxis., It can take placewhen we cross the directivity pattern inthe veryvicinity ofthe magneticaxis.
 Suchan example for the X-mode is shown on Fig., Suchan example for the X-mode is shown on Fig.
" for P= 1s, ὁ= 49°,a= 48.5°, A= 10°, yo= 50, and Tem=100 R."," \ref{core} for $\mbox{P = 1s}$ , $\zeta = 49^{\circ}$ ,$\alpha = 48.5^{\circ}$ , $\lambda = 10^3$ , $\gamma_{0} = 50$ , and $r_{\rm em} = 100 \, R$ ."
SOULCeS.,sources.
 Although thus successful in energies above 0.6 keV. the LAINB+2\UNL model again falls below the observed spectra. when extrapolated tosvard lower energies.," Although thus successful in energies above 0.6 keV, the LMXB+2MKL model again falls below the observed spectra, when extrapolated toward lower energies."
 This can be attributed to vel another diffuse component of a still lower temperature leaking into the spectra. as identified later in $ 3.3 using the ciffuse-emission spectra.," This can be attributed to yet another diffuse component of a still lower temperature leaking into the spectra, as identified later in $\S$ 3.3 using the diffuse-emission spectra."
 In order to quantilv the diffuse X-ray emission using the ΧΑΑ Νεο data. we jointly analvzed the residual MOS and PN spectra of the central 6’ region. obtained by removing (he point sources as described in § 2.," In order to quantify the diffuse X-ray emission using the ${\it XMM}$ ${\it Newton}$ data, we jointly analyzed the residual MOS and PN spectra of the central $\arcmin$ region, obtained by removing the point sources as described in $\S$ 2."
 The spectra. shown in Figure 4.. are considerably softer (han those for bright point sources (Figure 3)). and bear a clear hump over 0.71.0 keV due to the Fe-L complex indicative of the dominance of soft thermal emission.," The spectra, shown in Figure \ref{fig:0.6-7keV}, are considerably softer than those for bright point sources (Figure \ref{fig:lmxb}) ), and bear a clear hump over 0.7–1.0 keV due to the Fe-L complex indicative of the dominance of soft thermal emission."
 Moreover. the spectra extend significantly toward higher energies than would be expected Lor the softer thermal emission.," Moreover, the spectra extend significantly toward higher energies than would be expected for the softer thermal emission."
 This is mainly considered to have two origins., This is mainly considered to have two origins.
 One is the ~15% spill-over of photons Irom the detected bright point sources out of the accumulation regions: (his is estimated to contribute ~2.6x105 erg st., One is the $\sim 15\%$ spill-over of photons from the detected bright point sources out of the accumulation regions; this is estimated to contribute $\sim 2.6 \times 10^{38}$ erg $^{-1}$.
 The other is contribution from discrete sources below our detection limit. which is estimated to be <5xLO erg ! based on the Log AN Log 5$ relation of Nongetal. (2002)..," The other is contribution from discrete sources below our detection limit, which is estimated to be $\lesssim 5 \times 10^{37}$ erg $^{-1}$ based on the Log ${\it N}$ –Log ${\it S}$ relation of \citet{M31_LF_Kong}. ."
" Since the former dominates. we model the hard component in the ""dilfuse spectra (Figure 4)) bv the LAINB model in contrast to SEAOL who used a PL modeling. which is ruled out in our analvses (§ 3.1)."," Since the former dominates, we model the hard component in the “diffuse” spectra (Figure \ref{fig:0.6-7keV}) ) by the LMXB model in contrast to SEA01 who used a PL modeling, which is ruled out in our analyses $\S$ 3.1)."
 Because (he absorption column densitv of the summed. point-source spectra is consistent wilh the Galactic value (Table 2)). we here and herealter apply the Galactic Nyy to the LAINB component. as well as lo the (hin-thermal plasma components (58 3.1).," Because the absorption column density of the summed point-source spectra is consistent with the Galactic value (Table \ref{tab:lmxb}) ), we here and hereafter apply the Galactic $N_{\rm H}$ to the LMXB component, as well as to the thin-thermal plasma components $\S$ 3.1)."
 The abundances of elements are allowed to vary freely. but their relative values are constrained to follow the solar ratios.," The abundances of elements are allowed to vary freely, but their relative values are constrained to follow the solar ratios."
 We lentatively limited the fit energv band to 0.67 keV. to compare with the ASC results (Paper 1).," We tentatively limited the fit energy band to 0.6–7 keV, to compare with the ${\it ASCA}$ results (Paper 1)."
 The Si-Ix energy band (1.681.80 keV) of both MOS detectors was excluded. because the background. contribution and uncertainty are rather large there al. 2002)..," The Si-K energy band (1.68–1.80 keV) of both MOS detectors was excluded, because the background contribution and uncertainty are rather large there \citep{CAL_Ferrando}."
 Figure daa shows the simplest model. ie. one LAINB plus one-temperature MIXL model (LAINB+IAINEL model). fitted simultaneously to the MOS and PN spectra over (he 0.67 keV range.," Figure \ref{fig:0.6-7keV}a a shows the simplest model, i.e., one LMXB plus one-temperature MKL model (LMXB+1MKL model), fitted simultaneously to the MOS and PN spectra over the 0.6–7 keV range."
 The obtained parameters are summarized in Table 3.., The obtained parameters are summarized in Table \ref{tab:fit}.
 Thus. an MIXL component with the temperature 0.30.4 keV roughly accounts for the softer part of the spectrum. but Chere remains excess around 0.9 keV. which makes the fit statistically unacceptable.," Thus, an MKL component with the temperature 0.3–0.4 keV roughly accounts for the softer part of the spectrum, but there remains excess around 0.9 keV, which makes the fit statistically unacceptable."
 If. we use the same model as SEAOL. namely one PL plus one MIXL model (PL+IAMINL). the PL photon index ( 1.8) and the MIXL temperature (~0.4 keV) become close to their results (Figure 4bb and Table 3)). but the excess remains as well.," If we use the same model as SEA01, namely one PL plus one MKL model (PL+1MKL), the PL photon index $\sim 1.8$ ) and the MKL temperature $\sim 0.4$ keV) become close to their results (Figure \ref{fig:0.6-7keV}b b and Table \ref{tab:fit})), but the excess remains as well."
 Since, Since
10 observed. average count rate becomes three times that --n quiescence.,the observed average count rate becomes three times that in quiescence.
" Since the material is now connecting to field ines much further in (74 £24,4)) the magnetic colatitude of 1e accretion region climbs from ο330.", Since the material is now connecting to field lines much further in 4 $R$ ) the magnetic colatitude of the accretion region climbs from to.
.. Aceretion also arrives from a greater. range of azimuth. extending the accretion are around. the pole owards a complete ring.," Accretion also arrives from a greater range of azimuth, extending the accretion arc around the pole towards a complete ring."
 Parts of the extended. ring are rerclore always disappearing ancl appearing over the limb of the white cwarf and this. coupled with a gradient in accretion rate alone the ring. produces a sinusoidal spin pulse.," Parts of the extended ring are therefore always disappearing and appearing over the limb of the white dwarf and this, coupled with a gradient in accretion rate along the ring, produces a sinusoidal spin pulse."
 Davs 3 4: The accretion rate and the observed count rate decline gradually., Days 3 4: The accretion rate and the observed count rate decline gradually.
 Dav 5: The aceretion rate drops further and. the magnetosphere expands., Day 5: The accretion rate drops further and the magnetosphere expands.
 “Phe absorption decreases anc so the observed count rate increases. despite the lower accretion rate.," The absorption decreases and so the observed count rate increases, despite the lower accretion rate."
 The lower pole begins to appear again and the pulse amplitude drops. with Hux from the lower pole filling in the minima.," The lower pole begins to appear again and the pulse amplitude drops, with flux from the lower pole filling in the minima."
 The accretion ring at the upper pole shrinks and again lies entrielv on the visible hemisphere at some spin phases: thus the spin pulse again becomes Lat topped., The accretion ring at the upper pole shrinks and again lies entriely on the visible hemisphere at some spin phases; thus the spin pulse again becomes flat topped.
 Dav 6: XY Ari returns to quiescence., Day 6: XY Ari returns to quiescence.
 To back up the results so far we present some simple modelling of the accretion regions and thus of the spin pulse profiles., To back up the results so far we present some simple modelling of the accretion regions and thus of the spin pulse profiles.
 “Phe model follows that by Ixim BBeouermann (1995). tracing magnetic field lines back from the disce disruption radius. Rina... to their location on the white chvarl.," The model follows that by Kim Beuermann (1995), tracing magnetic field lines back from the disc disruption radius, $R$, to their location on the white dwarf."
 The geometry is completely specified by Laue. ó and 7. and by assuming a cipolar field and a negligible shock height.," The geometry is completely specified by $R$, $\delta$ and $i$, and by assuming a dipolar field and a negligible shock height."
" As the accretion rate is likely to be >greatest. alonee ""downhill fieldlines. we have introduced. a function. D cose. where à is the angle between the tangent to the field and the disc at the capture point."," As the accretion rate is likely to be greatest along `downhill' fieldlines, we have introduced a function, $A \pm B \cos\alpha$ , where $\alpha$ is the angle between the tangent to the field and the disc at the capture point."
 This weights the intensity of the X-ray emission around the rings. using opposite signs of B for the upper and lower pole respectively.," This weights the intensity of the X-ray emission around the rings, using opposite signs of $B$ for the upper and lower pole respectively."
 Depending on the choice of ;4 and £ the emission can occur from. all points on the ring (cf2 0). a aare (D= 0). or somewhere inbetween.," Depending on the choice of $A$ and $B$ the emission can occur from all points on the ring $A=0$ ), a arc $B=0$ ), or somewhere inbetween."
 The code sums the ας from the accretion regions not obscured by the white dwarf. ignoring absorption effects.," The code sums the flux from the accretion regions not obscured by the white dwarf, ignoring absorption effects."
 Fig., Fig.
 9 shows simulations with the parameters adjusted to mateh the NY Ari pulse profiles (Pig., 9 shows simulations with the parameters adjusted to match the XY Ari pulse profiles (Fig.
 3)., 3).
 The bottom curve uses parameters suitable for quiescence (/ =S27. 0 =157. Bus == 91Ut5.4)) and assumes that both poles are visible and are equally bright.," The bottom curve uses parameters suitable for quiescence $i$ =, $\delta$ =, $R$ = ) and assumes that both poles are visible and are equally bright."
 ΙΕ such a model is entirely. symmetric no net modulation is seen., If such a model is entirely symmetric no net modulation is seen.
 The Iow-amplitude modulation seen in quiescence requires svmmetry breaking. for instance a non-zero shock height of order 0.01. {να oor a dipole placed olf center by 00.05 Rea.," The low-amplitude modulation seen in quiescence requires symmetry breaking, for instance a non-zero shock height of order 0.01 $R$ or a dipole placed off center by 0.05 $R$."
 Phus for the bottom curve the dipole was ollset by 0.05 Rec iin a direction perpendicular to the plane formed by the spin axis and the (centred) dipole axis., Thus for the bottom curve the dipole was offset by 0.05 $R$ in a direction perpendicular to the plane formed by the spin axis and the (centred) dipole axis.
 This reproduces the small peak at phase 0.7 (Fig., This reproduces the small peak at phase 0.7 (Fig.
 3)., 3).
 The top profile in Fig., The top profile in Fig.
 9 uses the reduced. Zhi ool {νι aancl assumes that only the upper pole is visible. as deduced [or the peak of outburst.," 9 uses the reduced $R$ of 4 $R$ and assumes that only the upper pole is visible, as deduced for the peak of outburst."
 In order to obtain the 990 per cent modulation depth the are length has to be., In order to obtain the 90 per cent modulation depth the arc length has to be.
. The middle profile uses Zhi S= 7 Bau this reduces € compared to the outburst. peak. and so the upper pole lies entirely on the visible face of the white dwarl at some spin phases. producing a Lat-toppeἱ profile.," The middle profile uses $R$ = 7 $R$; this reduces $\epsilon$ compared to the outburst peak, and so the upper pole lies entirely on the visible face of the white dwarf at some spin phases, producing a flat-topped profile."
 We have assumed that the lower pole is partially visible. and given it a 50 per cent weighting compared. to the upper pole.," We have assumed that the lower pole is partially visible, and given it a 50 per cent weighting compared to the upper pole."
 This fills in the mimima. producing a pulse profile comparable to that during the decline from outburst (the slight asvmmetry in the curve arises because the centered dipole was used for all three curves).," This fills in the mimima, producing a pulse profile comparable to that during the decline from outburst (the slight asymmetry in the curve arises because the off-centered dipole was used for all three curves)."
 Figs., Figs.
 9 33 show that the model reproduces the gross features of the data., 9 3 show that the model reproduces the gross features of the data.
 Phere are some discrepancies. for instance the Ilat-topped section of the pulse observed. during the decline," There are some discrepancies, for instance the flat-topped section of the pulse observed during the decline"
"telescope) could give a similar FoM for (wo, wa) compared with a conservative space-based slitless survey, it will not give competitive FoM for (.Xo.67,X1.33, X2.0) (see Table 2).","telescope) could give a similar FoM for $w_0,w_a$ ) compared with a conservative space-based slitless survey, it will not give competitive FoM for $X_{0.67},X_{1.33},X_{2.0}$ ) (see Table 2)."
 It is important to recognize that both space and ground galaxy redshift surveys are required to obtain definitive measurement of dark energy using galaxy clustering., It is important to recognize that both space and ground galaxy redshift surveys are required to obtain definitive measurement of dark energy using galaxy clustering.
" Ongoing ground-based surveys, BOSS andWiggleZ?,, will enable us to test the methodology for extracting dark energy constraints from galaxy clustering data, and improve our understanding of systematic effects."," Ongoing ground-based surveys, BOSS and, will enable us to test the methodology for extracting dark energy constraints from galaxy clustering data, and improve our understanding of systematic effects."
" Proposed ground-based surveys, suchas BigBOSS (Schlegeletal.2009) andHETDEX®,, will be complementary to space-based surveys in using different tracer populations and redshift coverage."," Proposed ground-based surveys, suchas BigBOSS \citep{Schlegel09} and, will be complementary to space-based surveys in using different tracer populations and redshift coverage."
" There are other tracers of cosmic large scale structure that can be observed from the ground, and are also highly complementary in probing dark energy to the space-based surveys discussed in this paper."," There are other tracers of cosmic large scale structure that can be observed from the ground, and are also highly complementary in probing dark energy to the space-based surveys discussed in this paper."
" For example, ground-based Lya forest data can be used to study clustering of matter at z=2 to 4 (Croftetal. 2002),, and"," For example, ground-based $\alpha$ forest data can be used to study clustering of matter at $z= 2$ to 4 \citep{Croft02}, , and"
pointed out that for ος + C resonances far below the height of the Coulomb barrier. the entrance channel width Py is much smaller than (the total resonance width.,"pointed out that for $^{12}$ C + $^{12}$ C resonances far below the height of the Coulomb barrier, the entrance channel width $\Gamma_{1}$ is much smaller than the total resonance width."
 The latter (svhich is of order ~ 100 keV for the observed resonances above 2.4 MeV) is also noticeably smaller than the resonance energy., The latter (which is of order $\sim$ 100 keV for the observed resonances above 2.4 MeV) is also noticeably smaller than the resonance energy.
 In this case we have (Fowler. Caughlan. Zimmerman 1967) Therefore the partial width Py in equation (3.5) is to be evaluated al the resonance enerev £L—E..," In this case we have (Fowler, Caughlan, Zimmerman 1967) Therefore the partial width $\Gamma_{1}$ in equation (3.5) is to be evaluated at the resonance energy $E=E_{r}$."
 Ileve we notice that the resonance energy is shifted by the plasma effects., Here we notice that the resonance energy is shifted by the plasma effects.
 We take ££. to be the shifted resonance energy., We take $E_{r}$ to be the shifted resonance energy.
" The shilted resonance enerev £, is related (o the resonance energy in the vacuum E? by the relationship where the expression lor C is given by equation (3.3).", The shifted resonance energy $E_{r}$ is related to the resonance energy in the vacuum $E_{r}^{0}$ by the relationship where the expression for $C$ is given by equation (3.3).
" The barrier penetration factor P,CE) for the pure Coulomb potential Z?e?/r of the identical nuclei is known to be Therefore. the enhancement factor a of the resonant thermonuclear reaction rates which arises because of the plasma screening ellects is"," The barrier penetration factor $P_{0}(E)$ for the pure Coulomb potential $Z^{2}e^{2}/r$ of the identical nuclei is known to be Therefore, the enhancement factor $\alpha$ of the resonant thermonuclear reaction rates which arises because of the plasma screening effects is"
interactions are not fully understood and quantified.,interactions are not fully understood and quantified.
 This is especially true in the presence of the gas compoucut. as the above cited works are based on a purely collisionless 1uodoeling.," This is especially true in the presence of the gas component, as the above cited works are based on a purely collisionless modeling."
 Iu this paper we analyze the effect of the gaseous component ou the evolution of galactic bars and take a nunierical approach., In this paper we analyze the effect of the gaseous component on the evolution of galactic bars and take a numerical approach.
 In particular. we vary the eas coutent in the disk aud adwvauce models with various munerical resolution. but keep the stellar disk aud DM halo paraimaters uuchauged. using the Standard Model of Paper EL. We are especially interested iu how the gas presence and the associated nunerical resolution affect the basic parameters of a stellar bar. as well as the aueular 1iomentum transfer process in a barred disk aud in a disk-halo svsteni.," In particular, we vary the gas content in the disk and advance models with various numerical resolution, but keep the stellar disk and DM halo paramaters unchanged, using the Standard Model of Paper I. We are especially interested in how the gas presence and the associated numerical resolution affect the basic parameters of a stellar bar, as well as the angular momentum transfer process in a barred disk and in a disk-halo system."
 For this purpose. we perform a detailed comparison of models with gas to the standard collisionless⋅⋅ model published⋅ i ENPaper I. Woo nei then hybrid3 7 N-body ENand Sinooth: ParticlDEAN TIvdrodvuainics (SPI) FTARLL code (e. Heller Shlosiman 1991: Heller. Shlosniu Athanassoula 2007: Romano-Diaz et al.," For this purpose, we perform a detailed comparison of models with gas to the standard collisionless model published in Paper I. We use the hybrid $N$ -body and Smooth Particle Hydrodynamics (SPH) FTM-4.4 code (e.g., Heller Shlosman 1994; Heller, Shlosman Athanassoula 2007; Romano-Diaz et al."
 2009) to evolve the stellar and gaseous disks embedded in the DAL halos., 2009) to evolve the stellar and gaseous disks embedded in the DM halos.
 The eravitational forces ave calculated using the FalcON routine (Delnen 2002) which scales as O(N)., The gravitational forces are calculated using the FalcON routine (Dehnen 2002) which scales as $O(N)$.
" The adopted units are the suue as ia Paper E the units of mines and distanceB are taken as E104)NI. and E10 kpe respectively,", The adopted units are the same as in Paper I: the units of mass and distance are taken as $10^{11}~{\rm M_\odot}$ and 10 kpc respectively.
 This makes the unit of tine equal to LTosLOT yu. when ο lo and the velocity unit 208aus+.," This makes the unit of time equal to $4.7\times 10^7$ yr, when $G = 1$ , and the velocity unit $208~~{\rm km~s^{-1}}$."
 The evavitational softening is eq=0.016 for stars and DM particles.," The gravitational softening is $\epsilon_{\rm grav} = 
0.016$ for stars and DM particles."
 For the gas particles we use a dynamical softening., For the gas particles we use a dynamical softening.
 The eravitational softening is set to the smoothing length wuless the smoothing length falls below the fixed limiting valne έναν , The gravitational softening is set to the smoothing length unless the smoothing length falls below the fixed limiting value $\epsilon_{\rm dyn}$.
The models consist of a stellar disk with WV.=2& 105. a eas disk with Na0100 and DM halo with Vpyy=10° collisionless particles.," The models consist of a stellar disk with $N_{\rm *} = 2\times 10^5$ , a gas disk with $N_{\rm gas} = 4\times 10^4$ and DM halo with $N_{\rm DM} = 10^6$ collisionless particles."
 Models were evolved for about a IDuibble time. At=—270 in the €adopted units.DE which translatesUitte: fo ↕⊲≻⊤," Models were evolved for about a Hubble time, $\Delta t=270$ in the adopted units, which translates to 12.7 Gyr."
≼∶∏⋅ During∙⋅. the model evolution.. woe have routinely. observed the formation of very compact acciunulatious of gas cudparticles iu the ceutral disk mUregion.," During the model evolution, we have routinely observed the formation of very compact accumulations of gas particles in the central disk region."
 Iu the models with a substautial eas component we have observed the ornation of secondary gaseous bars which decoupled roni) the large-scale bars and contracted subsequently ο spatial scales where iusufficieut resolution resulted iu Hattened “blobs” a few softening lengths im radius. iu aerectuent with simulations of Eugluuüer Shlosiman (2001).," In the models with a substantial gas component we have observed the formation of secondary gaseous bars which decoupled from the large-scale bars and contracted subsequently to spatial scales where insufficient resolution resulted in flattened `blobs' a few softening lengths in radius, in agreement with simulations of Englmaier Shlosman (2004)."
 Besides the gradual capture of additional gas articles. the morphological‘ structure of the blob has evolved very little. aud was bevoud the resolution. of. our uodels.," Besides the gradual capture of additional gas particles, the morphological structure of the blob has evolved very little and was beyond the resolution of our models."
 For the sake of shortening the computational ue. we have replaced the gas particles trapped in he center bv stellar ones.," For the sake of shortening the computational time, we have replaced the gas particles trapped in the center by stellar ones."
 This operation was repeated whenever the central gas accumulation was substautiallv slowing down the overall evolution., This operation was repeated whenever the central gas accumulation was substantially slowing down the overall evolution.
 We kept the softening of the iudividual particles when they were converted from oue type to the other., We kept the softening of the individual particles when they were converted from one type to the other.
 We have runi a large uuuber of tests to verify that this action did not affect the evolution of the bar., We have run a large number of tests to verify that this action did not affect the evolution of the bar.
 This was achieved by running parallel models with and without the gas particle replacciuent., This was achieved by running parallel models with and without the gas particle replacement.
 The initial conditions of the stellar and DM particles were created with the procedures described in Paper I. using the density profiles from Heruquist (1993).," The initial conditions of the stellar and DM particles were created with the procedures described in Paper I, using the density profiles from Hernquist (1993)."
 The mass volume density distribution in the disk is eiven iu cexliudiical coordinates by where Ma is the disk mass. fi is a radial scale length and ty is a vertical scaleheight.," The mass volume density distribution in the disk is given in cylindrical coordinates by where $M_\mathrm{d}$ is the disk mass, $h$ is a radial scale length and $z_0$ is a vertical scaleheight."
" The deusitv of the spherical halo is given by wherePM AA, is. the mass-- of the halo. rc isD a Cassia""EE ↸⊳∏↑∪↕−↥↥⋜∥∐∏↴∖⋜⋯≼↧↸↕↴∖↑↕∐∖↸⊳∪⊓∖↥⋜∥∐∏↴∖∙∩↕↴∖↑∐"," The density of the spherical halo is given by where $M_\mathrm{h}$ is the mass of the halo, $r_\mathrm{c}$ is a Gaussian cutoff radius and $\gamma$ is the core radius."
"↸∖⋅⋅⋅↴ ⋅↴ ⋅⋅⋅↴ ⋅↴ normalization coustaut defined by with,. y=fre", $\alpha$ is the normalization constant defined by with $q=\gamma/r_\mathrm{c}$.
" The article. s. dispersion. velocities and asvuunuetric drift)velocities, corrections were calculated using 1ioinents of the collisionless Doltzinaun— equation."," The particle velocities, dispersion velocities and asymmetric drift corrections were calculated using moments of the collisionless Boltzmann equation."
 Since models thus constructed are not in exact virial cquilibrimu. the halo courponeut was relaxed for fo1 m the frozen disk potential.," Since models thus constructed are not in exact virial equilibrium, the halo component was relaxed for $t\sim 40$ in the frozen disk potential."
 DecanusEversay weOA areqvo ]interestedPORTO to quantifv+MA. theD effectMoe of1 the gas fraction aud spatial resolution on the bar evolution. we use the pure stellar model SD from Paper I as our heuchimark model (Table 1).," Because we are interested to quantify the effect of the gas fraction and spatial resolution on the bar evolution, we use the pure stellar model SD from Paper I as our benchmark model (Table 1)."
 A fixed fraction of stellar disk particles at f—0 were converted. to ideutical mass gas particles aud re-baluiced using the central attraction forces from the total nass distyilition., A fixed fraction of stellar disk particles at $t=0$ were converted to identical mass gas particles and re-balanced using the central attraction forces from the total mass distribution.
" The is considered to he isothermal with Z4,=10 K. and initiallygas moves on circular orbits."," The gas is considered to be isothermal with $T_{\rm gas}=10^4$ K, and initially moves on circular orbits."
 uM created i set of molIs COVEme a two (uunuensonatouM paranieter Spaces Dy VALENS που Gals mass fraction.f... in the disk. aud by chaneine the ⋅⋅⋅ ↕∐⋯⊓∐∶↴∙⊾↖⇁⋜↧↕⋯∖∪↕↑↕∐∖∶↴∙⊾↥⋅⋜↧↖⇁↕↑⋜↧⊓∪∐⋜↧↕↴∖↴∪⇈↸∖∐∐↕∶↴⋁∙←↽≖⇂∖↽⊔∙ ⋅ ⋅∙ 1 iο," We have created a set of models covering a two dimensional parameter space: by varying the gas mass fraction, in the disk, and by changing the limiting value of the gravitational softening, $\epsilon_{\rm dyn}$, in the gas."
" js(Y:S. ""iiο sM Qi: he de(Y:s meSS ;WelsUM e mconsa", The sum of $+$ gas mass was kept constant in the models.
 m cu DT i Wet l VAM Jucons ⋅⋅∣⊓ and ⊽∙⋅ ando 0.016. 0Οδ dO.1," We used the values $f_\mathrm{g} = 0\%$ , and , and $ = 0.016$ , 0.05 and 0.1."
 P hesofteninguscdforthestellaraid:. DM compon, The softening used for the stellar and DM components is fixed at 0.016.
ci aud values used in: each inodel., Table \ref{param} shows the combination of and values used in each model.
 Disks with a high coutent of eas tend to form dense chuups of gas which interact withthe stellar courponenut and raisethe velocity dispersion in the disk., Disks with a high content of gas tend to form dense clumps of gas which interact withthe stellar component and raisethe velocity dispersion in the disk.
 Tf this rise iu the disk: ‘temperature’ ⋅⋅ijs substautial.⋅ thebar instability," If this rise in the disk `temperature' is substantial, thebar instability"
 Tf this rise iu the disk: ‘temperature’ ⋅⋅ijs substautial.⋅ thebar instability⋅," If this rise in the disk `temperature' is substantial, thebar instability"
 Tf this rise iu the disk: ‘temperature’ ⋅⋅ijs substautial.⋅ thebar instability⋅κ," If this rise in the disk `temperature' is substantial, thebar instability"
 Tf this rise iu the disk: ‘temperature’ ⋅⋅ijs substautial.⋅ thebar instability⋅κ.," If this rise in the disk `temperature' is substantial, thebar instability"
sequence of Type 3 inodels. the polarization at the edge eradually increases as a function of D.s/λος until D.s/Ds0.23.,"sequence of Type 3 models, the polarization at the edge gradually increases as a function of $\ratio$ until $\ratio \approx 0.23$."
 Past this poiut. the maxinuun polarization for the flament occurs at the edge. so the width of the depolarized region becomes equal to the width of the filament.," Past this point, the maximum polarization for the filament occurs at the edge, so the width of the depolarized region becomes equal to the width of the filament."
 The polarization patterus duc to helically maguetized filaments become more complicated for filamcuts that are inclined relative to the plane of the sky., The polarization patterns due to helically magnetized filaments become more complicated for filaments that are inclined relative to the plane of the sky.
 As an example. in Figure L.. we show maps of one filament tilted 107. 307. aud 607 relative to the plane of the sky.," As an example, in Figure \ref{fig:inc}, we show maps of one filament tilted $10^\circ$, $30^\circ$, and $60^\circ$ relative to the plane of the sky."
 The fikunent model shown in Fieve Lis the sale one as shown in the top two paucls of Figure 1.., The filament model shown in Figure \ref{fig:inc} is the same one as shown in the top two panels of Figure \ref{fig:types}.
 Each polarization vector always remains parallel or perpendicular to the svuumetry axis for the reasous discussed in Section 12. above. but the patterus become asviunuetrce relative to the central axis of the filament.," Each polarization vector always remains parallel or perpendicular to the symmetry axis for the reasons discussed in Section \ref{sec:patterns} above, but the patterns become asymmetric relative to the central axis of the filament."
 It is particularly notable that laments inclined relative to the line of sight usually result iu asvuuuetric polarization patterns that include a 90° flip im the oricutation of the polarization vectors ou oue or both sides of the filament., It is particularly notable that filaments inclined relative to the line of sight usually result in asymmetric polarization patterns that include a $90^\circ$ flip in the orientation of the polarization vectors on one or both sides of the filament.
 The main purpose of this paper is to calculate polarization maps for the models of filamentary uolecular clouds that we presented in FPI aud FP2., The main purpose of this paper is to calculate polarization maps for the models of filamentary molecular clouds that we presented in FP1 and FP2.
 When the filament is orieuted parallel to the plane of he sky. the polarization patterus resulting from our models cau be qualitatively classified iuto three main vpes of behaviour.," When the filament is oriented parallel to the plane of the sky, the polarization patterns resulting from our models can be qualitatively classified into three main types of behaviour."
 The polarization vectors are parallel to the svuuuctry axis of the filament (Type 1) when the toroidal field is dominant. aud perpendicular (Type 2) for models whose poloidal field dominates.," The polarization vectors are parallel to the symmetry axis of the filament (Type 1) when the toroidal field is dominant, and perpendicular (Type 2) for models whose poloidal field dominates."
 The third type of pattern is more complicated. with polarization vectors that flip from being aligned parallel o the filament axis to perpendicular at some radius or radii (n projection).," The third type of pattern is more complicated, with polarization vectors that flip from being aligned parallel to the filament axis to perpendicular at some radius or radii (in projection)."
" Generally. the first type of xittern occurs when P.5/D,,s~0.1. while the second type occurs DB.s/D,,s~ 0:33."," Generally, the first type of pattern occurs when $\ratio \appleq 0.1$, while the second type occurs $\ratio\appgeq 0.33$ ."
" The third type of xilarization pattern occurs at intermediate values of δες Our most important result is that helical fields result in ""polarization holes.” iu which the emission is depolarized at some positious in the interior of he filament."," The third type of polarization pattern occurs at intermediate values of $\ratio$ Our most important result is that helical fields result in “polarization holes,” in which the emission is depolarized at some positions in the interior of the filament."
" Our Type Limodels qualitatively agree with the polarization structure of the ""integral-shiaped? filament iu Orion A (Matthews Wilson 2000).", Our Type 1 models qualitatively agree with the polarization structure of the “integral-shaped” filament in Orion A (Matthews Wilson 2000).
 Alauv of our models result in polarization patterns that contain 90° flips iu the oricutation of the »)larization vectors., Many of our models result in polarization patterns that contain $90^\circ$ flips in the orientation of the polarization vectors.
 The flips are sviuuetric about the svnunetry axis when a filament with comparable »oloidal and toroidal magnetic field streueths is oriented parallel to the plane of the sla., The flips are symmetric about the symmetry axis when a filament with comparable poloidal and toroidal magnetic field strengths is oriented parallel to the plane of the sky.
 Fibuneuts that are inclined relative to the line of sight generally result iu αποταο polarization flips ou one or both sides of he central axis of the filament., Filaments that are inclined relative to the line of sight generally result in asymmetric polarization flips on one or both sides of the central axis of the filament.
 We lave considered only 110dels of nou-tragmented. cvlindrically svuuuctric flameuts in this paper.," We have considered only models of non-fragmented, cylindrically symmetric filaments in this paper."
 This is obviously an idealization of real Glamenutary clouds. many of which have suffered gravitational fracimentation aud formed strings of embedded cores (cf.," This is obviously an idealization of real filamentary clouds, many of which have suffered gravitational fragmentation and formed strings of embedded cores (cf."
 Schneider Ehucercen 1979. Dutrev et al.," Schneider Elmegreen 1979, Dutrey et al."
 1991: see also FP2 and ΕΤ). and may also deviate from perfect cvlndzrical svunuetry.," 1991; see also FP2 and FP3), and may also deviate from perfect cylindrical symmetry."
 We will address the polarization patterus of embedded cores in a future analysis., We will address the polarization patterns of embedded cores in a future analysis.
 It is difficult to comment ou the ecueral effects of asvuuuetry., It is difficult to comment on the general effects of asymmetry.
 However. it is worthwhile to cousider the simplest type of non-axisvnuunetrie perturbation. that corresponding to a kink mode in which the flament is beut into a transverse sinusoidal wave.," However, it is worthwhile to consider the simplest type of non-axisymmetric perturbation, that corresponding to a kink mode in which the filament is bent into a transverse sinusoidal wave."
 If the sinusoid lies iu the plane of the sky.the frout to back sviuuietiy," If the sinusoid lies in the plane of the sky,the front to back symmetry"
the range of accretion rates in which delayed mixed bursts occur becomes notably truncated as well.,the range of accretion rates in which delayed mixed bursts occur becomes notably truncated as well.
" More importantly, a new burst regime appears at accretion rates close to Mg, where the bursts have o«100 and are hence prompt."," More importantly, a new burst regime appears at accretion rates close to $\dot{M}_{\mathrm{Edd}}$, where the bursts have $\alpha < 100$ and are hence prompt."
" This regime of bursts near the Eddington limit was hypothesized by CNO06, but it was not present in either the NHO3 or CN06 models due to the exclusion of hot CNO cycle breakout reactions."," This regime of bursts near the Eddington limit was hypothesized by CN06, but it was not present in either the NH03 or CN06 models due to the exclusion of hot CNO cycle breakout reactions."
 The usual delayed mixed burst regime and this new prompt mixed burst regime are separated by a short range of accretion rates in which bursts do not occur., The usual delayed mixed burst regime and this new prompt mixed burst regime are separated by a short range of accretion rates in which bursts do not occur.
" These results obtained with larger values of agree rather well with previous theoretical models (Fujimotofi,etal.1981;AyasliBildsten1998;PaczynskiFiskeretal.2003;Heger 2005),, most of which predicted bursts to occur up to roughly the Eddington limit."," These results obtained with larger values of $\frp$ agree rather well with previous theoretical models \citep{FHM81,AJ82,P83,T85,TWL96,B98,FHLT03,HCW05}, most of which predicted bursts to occur up to roughly the Eddington limit."
 Figure | thus explains the origin of the differences between those models and the models of NH03 and 6Ν06., Figure \ref{alphagraph} thus explains the origin of the differences between those models and the models of NH03 and CN06.
" If the breakout reaction rate is large, say ftp20.1, the results agree with most published models (which include breakout reactions in full and if the rate is small, fiyS0.1, the results are similar to strength),those obtained by NH03 and CN06 (who effectively set fi,— 0)."," If the breakout reaction rate is large, say $\frp \gtrsim 0.1$, the results agree with most published models (which include breakout reactions in full strength), and if the rate is small, $\frp \lesssim 0.1$, the results are similar to those obtained by NH03 and CN06 (who effectively set $\frp = 0$ )."
 We reiterate that the latter results agree much better with observations., We reiterate that the latter results agree much better with observations.
" For laceSS0.1, the reached during steady-state nuclear burning in the temperaturesaccreted layer are too low for significant leakage out of the hot CNO cycle via breakout reactions, and so these reactions have a negligible effect on the onset of type I X-ray bursts, regardless of the true IO(a,y)?Ne cross section."," For $\lacc \lesssim 0.1$, the temperatures reached during steady-state nuclear burning in the accreted layer are too low for significant leakage out of the hot CNO cycle via breakout reactions, and so these reactions have a negligible effect on the onset of type I X-ray bursts, regardless of the true $^{15}$ $\alpha$ $\gamma$ $^{19}$ Ne cross section."
" However, breakout reactions should affect the nuclear flow during the burst itself, and hence they could influence the burst lightcurve and possibly also the onset of subsequent bursts (Fiskeretal.2006)."," However, breakout reactions should affect the nuclear flow during the burst itself, and hence they could influence the burst lightcurve and possibly also the onset of subsequent bursts \citep{FGWD06}."
". CNO06 showed that the delayed mixed bursts of NH03 occur when Xg, the column depth at which H is depleted via stable nuclear as measured from the stellar surface, is less than but close burningto Nye, the column depth at which He is depleted via stable nuclear burning."," CN06 showed that the delayed mixed bursts of NH03 occur when $\sh$, the column depth at which H is depleted via stable nuclear burning as measured from the stellar surface, is less than but close to $\she$, the column depth at which He is depleted via stable nuclear burning."
" The hot CNO cycle breakout reaction sequence (1) eliminates seed nuclei from the hot CNO cycle which slows H burning, thereby increasing Mg, and (ii) provides an additional pathway by which He may burn, thereby decreasing Xe."," The hot CNO cycle breakout reaction sequence (i) eliminates seed nuclei from the hot CNO cycle which slows H burning, thereby increasing $\sh$, and (ii) provides an additional pathway by which He may burn, thereby decreasing $\she$."
" Consequently, one expects the regime of delayed mixed bursts to occur at lower lace if breakout reactions are included."," Consequently, one expects the regime of delayed mixed bursts to occur at lower $\lacc$ if breakout reactions are included."
" Furthermore, CN06 showed that the interplay between H burning via the hot CNO cycle and He burning via triple-o reactions is integral to generating the oscillations that precede delayed mixed bursts."," Furthermore, CN06 showed that the interplay between H burning via the hot CNO cycle and He burning via $\alpha$ reactions is integral to generating the oscillations that precede delayed mixed bursts."
 Breakout reactions diminish this interplay by eliminating the hot CNO cycle seed nuclei and should therefore reduce the range of accretion rates over which delayed mixed bursts occur., Breakout reactions diminish this interplay by eliminating the hot CNO cycle seed nuclei and should therefore reduce the range of accretion rates over which delayed mixed bursts occur.
" Figure l illustrates both of these effects (compare fj,=1 with Sip=0)."," Figure \ref{alphagraph}
 illustrates both of these effects (compare $\frp=1$ with $\frp=0$ )."
" For accretion rates above those at which delayed mixed bursts occur, NHO3 found that nuclear burning is always stable and therefore bursts do not occur, in accord with observations (vanetal.1979,1988;Cornelisseetal.2003;RemillardParadijs 2006)."," For accretion rates above those at which delayed mixed bursts occur, NH03 found that nuclear burning is always stable and therefore bursts do not occur, in accord with observations \citep{vPCLJ79,vPPL88,Cetal03,RLCN06}."
". According to their model, for 0.3Slace1, steady-state H burning via the hot CNO cycle increases the effective radiative cooling rate and thereby suppresses a thin-shell thermal instability (CN06)."," According to their model, for $0.3 \lesssim \lacc \lesssim 1$, steady-state H burning via the hot CNO cycle increases the effective radiative cooling rate and thereby suppresses a He-triggered thin-shell thermal instability (CN06)."
" However,He-triggered their model did not include breakout reactions."," However, their model did not include breakout reactions."
" If the 'O(a,y)!°Ne breakout reaction rate is significant in this regime, this reaction will suppress hot CNO cycle H burning, and thus He burning via triple-a reactions will govern the total reaction rate, since the α reaction rate will be the slowest rate in the nuclear flow at these high /a¢¢."," If the $^{15}$ $\alpha$ $\gamma$ $^{19}$ Ne breakout reaction rate is significant in this regime, this reaction will suppress hot CNO cycle H burning, and thus He burning via $\alpha$ reactions will govern the total reaction rate, since the $\alpha$ reaction rate will be the slowest rate in the nuclear flow at these high $\lacc$."
" The notion that the triple-a reaction rate is the slowest rate in the nuclear flow is precisely what zone type I X-ray burst ignition models assume (Fujimotosten2000;Hegeretal. 2005),, and we suggest that this is the reason why the results of one-zone and time-dependent multi-zone models agree so well."," The notion that the $\alpha$ reaction rate is the slowest rate in the nuclear flow is precisely what one-zone type I X-ray burst ignition models assume \citep{FHM81,P83,B98,CB00,HCW05}, and we suggest that this is the reason why the results of one-zone and time-dependent multi-zone models agree so well."
" All of these models predict that nuclear burning is thermally stable for temperatures 5x105 K, which are reached only at accretion rates lace>1 (e.g.,Schatzetal.1999)."," All of these models predict that nuclear burning is thermally stable for temperatures $T \gtrsim 5 \times
10^{8}$ K, which are reached only at accretion rates $\lacc \gtrsim 1$ \citep[e.g.,][]{SBCW99}."
". Figure 1 illustrates that, for relatively large values of fij, there is a regime of prompt bursts at accretion rates greater than the accretion rates at which delayed mixed bursts occur, and the critical accretion rate above which these prompt mixed bursts cease is roughly Miggq, in very good agreement with these other models."," Figure \ref{alphagraph} illustrates that, for relatively large values of $\frp$, there is a regime of prompt bursts at accretion rates greater than the accretion rates at which delayed mixed bursts occur, and the critical accretion rate above which these prompt mixed bursts cease is roughly $\dot{M}_{\mathrm{Edd}}$, in very good agreement with these other models."
 Figure 1 illustrates also that nuclear burning is stable for the small of Jace between these two regimes., Figure \ref{alphagraph} illustrates also that nuclear burning is stable for the small range of $\lacc$ between these two regimes.
" These accretion rates are rangehigh enough to suppress delayed mixed bursts, but they are not high enough to cause a sufficient leakage out of hot CNO cycle H burning and trigger a thermal instability."," These accretion rates are high enough to suppress delayed mixed bursts, but they are not high enough to cause a sufficient leakage out of hot CNO cycle H burning and trigger a thermal instability."
" The model we in this work suggests that, if the true >O(a,7)!°Ne cross presentsection is greater than approximately 0.1 of the CF88 rate, type I X-ray bursts should occur in systems with accretion rates near the Eddington limit."," The model we present in this work suggests that, if the true $^{15}$ $\alpha$ $\gamma$ $^{19}$ Ne cross section is greater than approximately $0.1$ of the CF88 rate, type I X-ray bursts should occur in systems with accretion rates near the Eddington limit."
" Observations indicate that this is not the case, since low-mass X-ray binaries with lace20.3 generally do not exhibit bursts."," Observations indicate that this is not the case, since low-mass X-ray binaries with $\lacc \gtrsim 0.3$ generally do not exhibit bursts."
" This suggests that the true O(a,7)!®Ne reaction rate is less than that proposed by CF88.", This suggests that the true $^{15}$ $\alpha$ $\gamma$ $^{19}$ Ne reaction rate is less than that proposed by CF88.
" This conclusion is complementary to that of Fiskeretal.(2006),, who were the first to propose that the occurrence of type I X-ray bursts is sensitive to the strength of the PO(a,A)'?Ne reaction rate, and who found that the existence of bursts in with lace720.1 suggests a lower bound on this rate."," This conclusion is complementary to that of \citet{FGWD06}, who were the first to propose that the occurrence of type I X-ray bursts is sensitive to the strength of the $^{15}$ $\alpha$ $\gamma$ $^{19}$ Ne reaction rate, and who found that the existence of bursts in systems with $\lacc
\approx 0.1$ suggests a lower bound on this rate."
 This systemslower bound corresponds to fry70.05 in our notation., This lower bound corresponds to $\frp \approx 0.05$ in our notation.
" We also note that a relatively low !°O(a,7)!?Ne reaction rate is consistent with the existence of carbon-triggered superbursts, since a low"," We also note that a relatively low $^{15}$ $\alpha$ $\gamma$ $^{19}$ Ne reaction rate is consistent with the existence of carbon-triggered superbursts, since a low"
"W531 is one of the largest ccolplexes in the Galaxy. with intense star-forming regions that lave been observed froni radio fo lear-infrared wavelengths (o.ο, Ghosh et al.","W31 is one of the largest complexes in the Galaxy, with intense star-forming regions that have been observed from radio to near-infrared wavelengths (e.g. Ghosh et al."
 1989: Dhu ct al., 1989; Blum et al.
 2001: Ian Ίου 2002)., 2001; Kim Koo 2002).
 At low spalal resolution. W31 appears as three main exteed lTroglons: 0. 0.1 aud O.L (Shaver Coss 1970) (Fig.," At low spatial resolution, W31 appears as three main extended regions: $-$ 0.3, $-$ 0.1 and $-$ 0.4 (Shaver Goss 1970) (Fig."
 1)., 1).
 The radio nenila ) lies within W231 on the plane o [the sky (Fig., The radio nebula $-$ 0.3 lies within W31 on the plane of the sky (Fig.
 1). aud has dvawn considerable attention iu recen vears due to the iutrieuiug objects nearby.," 1), and has drawn considerable attention in recent years due to the intriguing objects nearby."
 At one tine. 0.3 was SUESed tobe a plerionic superhova remnant poweredby a rare soft gamuna-ray repeater. SCR 20 Usullsariui Frail 1993: Ixouvoliotou et al.," At one time, $-$ 0.3 was suggested to be a plerionic supernova remnant powered by a rare soft gamma-ray repeater, SGR $-$ 20 (Kulkarni Frail 1993; Kouveliotou et al."
 1998)., 1998).
 SCR. 2. dn urn. was though to be associated with au alHost οally rare luminous blue varialle (LBV) sar (van Werkwijk et al.," SGR $-$ 20, in turn, was thought to be associated with an almost equally rare luminous blue variable (LBV) star (van Kerkwijk et al."
 1995) which lies at the time-variable (iu oth Hux and morphology) cox! of this nebula (Vasisht. Frail. I&ulkunui 1995).," 1995) which lies at the time-variable (in both flux and morphology) core of this nebula (Vasisht, Frail, Kulkarni 1995)."
 Towever. the revised Iuter-Planetary Network. (IPN) localization of SCR 20 provicC8 a »osition 1iconsistent (see Fie.," However, the revised Inter-Planetary Network (IPN) localization of SGR $-$ 20 provides a position inconsistent (see Fig."
11) with that of the ΤΟΝ star and radio core of (IIurlev et al., 1) with that of the LBV star and radio core of $-$ 0.3 (Hurley et al.
" 1999). hough the LDV position is consistent with the radio core within the uwertainties (~ 2""),"," 1999), though the LBV position is consistent with the radio core within the uncertainties $\sim 2 \arcsec$ )."
" Receut and infrared observations confirm that the SCR lies ~12"" away from the LBV aud radio core (Eikeuberry et al.", Recent and infrared observations confirm that the SGR lies $\sim 12 \arcsec$ away from the LBV and radio core (Eikenberry et al.
 2001: Ixaplan ct al., 2001; Kaplan et al.
 2002)., 2002).
 Furthermore. Cacnusler et al. (," Furthermore, Gaensler et al. ("
2001) argue that 4.3 is nof a supernova reninaut at all.|mit is rather powered by the tremendous wind of the LBV star at its core.,"2001) argue that $-$ 0.3 is not a supernova remnant at all, but is rather powered by the tremendous wind of the LBV star at its core."
 Iuared observations of the field of SCR 20 revealti the LBV star is not alouc. but appears to be part of a cluster of eumibedded. hot. hninous stars (Fuchsetal. 1999)). and the IPN position for SCR 20 Is consistent with ieership in the star cluster (Eikeuberry et al.," Infrared observations of the field of SGR $-$ 20 reveal that the LBV star is not alone, but appears to be part of a cluster of embedded, hot, luminous stars \cite{fuc99}) ), and the IPN position for SGR $-$ 20 is consistent with membership in the star cluster (Eikenberry et al."
 2001)., 2001).
 Cuven this somewhat confusing history. we take a moment to sTunnuaurize our curren understanding of CL0.0-0.3: 1.," Given this somewhat confusing history, we take a moment to summarize our current understanding of G10.0-0.3: 1."
 CHO.0.3 Is a radio nebula (NOT stpernowva reniit) wit1 enission powered by the LBV star spatially coin‘ident with its core., G10.0-0.3 is a radio nebula (NOT supernova remnant) with emission powered by the LBV star spatially coincident with its core.
 2., 2.
 The LBV Say is par of a cluster of huiLS Ssars enibeded iu a amolecular cloud., The LBV star is part of a cluster of luminous stars embedded in a molecular cloud.
 3., 3.
 yous likely anot10r nenaY of this cluster of stars. and is spatially distiict from t16 LBV star.," is likely another member of this cluster of stars, and is spatially distinct from the LBV star."
 We plot in Fig., We plot in Fig.
 1 the various oljects (ancl t111 relative location on he plaιο of the skv) flat we will discuss i1 this paper., 1 the various objects (and their relative location on the plane of the sky) that we will discuss in this paper.
" Because of the proximity of these unusualiusight oljeC""5. he distance to them can provide significant inte(their oluwsical properties. eiviug this measurenient parictlar nuportauce."," Because of the proximity of these unusual objects, the distance to them can provide significant insight into their physical properties, giving this measurement particular importance."
 Corbel et al. (, Corbel et al. (
16JOT) proposed distaIce estimate based on observations of molecular clouds aleng he line of sight.,1997) proposed a distance estimate based on observations of molecular clouds along the line of sight.
" They used CO spectroscopy to esnuate he hwdrogeu coluun density towards 0,D.? ancl youn that an absorption column «ensitv as a fuucion of distance."," They used $CO$ spectroscopy to estimate the hydrogen column density towards $-$ 0.3, and from that an absorption column density as a function of distance."
 Taking the measured X-rax absorption fvvards SCR 20 and aji estimate of the optical extiuction o the LBV star. thev cowcluded that (L3. aand W31 he 115x|LL kpe freLL the Sun.," Taking the measured X-ray absorption towards SGR $-$ 20 and an estimate of the optical extinction to the LBV star, they concluded that $-$ 0.3, and W31 lie $14.5 \pm 1.4$ kpc from the Sun."
 Ilowever. this situation seeLLC complicated V Lowor infrared stellar spectroscopy bv Bhuu et al. ," However, this situation seemed complicated by newer infrared stellar spectroscopy by Blum et al. ("
2001) aud racBo/nulluineter observatious bv Iam Ίνου (2002).,2001) and radio/millimeter observations by Kim Koo (2002).
 Bhim et al. (, Blum et al. (
2001) present iuyared spectra of nembers of a star cluster iu he ireeion 0.3. also within W31 ou the plane of he sky.,"2001) present infrared spectra of members of a star cluster in the region $-$ 0.3, also within W31 on the plane of the sky."
 Based ou the spectra. they derive spectralΜπτοςitv classes and extinctions for the stars. which. combined with oeinfrared photometry. place them and 3 naildieuouslv at a distance of dot3. Lispe.," Based on the spectra, they derive spectral/luminosity classes and extinctions for the stars, which, combined with infrared photometry, place them and $-$ 0.3 unambiguously at a distance of $d \simeq 3.4$ kpc."
 Towever. he extinction towards 0.3 is 1uuch siualler (Ανzm15 lage) tiui the extinction towards aud he star chister close to (Eikeuberry et al," However, the extinction towards $-$ 0.3 is much smaller $\Delta
A_V \approx 15$ mag) than the extinction towards and the star cluster close to (Eikenberry et al."
 2001) aud the correlated ταν absorpion towards, 2001) and the correlated X-ray absorption towards
"those given by Rybka (1984), separately for the stars in aand for all stars inHevelius.","those given by Rybka (1984), separately for the stars in and for all stars in."
". In most cases the identifications are identical, but there are differences."," In most cases the identifications are identical, but there are differences."
 We have identified 24 stars (among which 1 repeated entry) that Rybka could not identify., We have identified 24 stars (among which 1 repeated entry) that Rybka could not identify.
 In 6 cases where two stars are plausible counterparts we choose the stars that Rybka did not choose., In 6 cases where two stars are plausible counterparts we choose the stars that Rybka did not choose.
 In 48 cases (among which 1 repeated entry) we do not agree with the identification in Rybka (1984); this includes 9 stars which we cannot identify., In 48 cases (among which 1 repeated entry) we do not agree with the identification in Rybka (1984); this includes 9 stars which we cannot identify.
 This number doesnot include the 21 cases where our emendation to Rybka leads to a different Hipparcos identification (see refs:rybkae))., This number does include the 21 cases where our emendation to Rybka leads to a different Hipparcos identification (see \\ref{s:rybkae}) ).
" In some cases the identification given by Rybka (1984) has such a large positional offset, or is so faint, that we consider our rejection secure; in other cases we choose another closer and/or brighter star as a more plausible counterpart."," In some cases the identification given by Rybka (1984) has such a large positional offset, or is so faint, that we consider our rejection secure; in other cases we choose another closer and/or brighter star as a more plausible counterpart."
 Details may be found in refs:notes.., Details may be found in \\ref{s:notes}.
" reft:rybka shows that there are 16 stars, oone percent of the"," \\ref{t:rybka} shows that there are 16 stars, one percent of the"
Taurus population to be coeval.,Taurus population to be coeval.
 It is thus legitimate to compare our brown cwarf disk masses with those of T Tauri stus in star forming regions wilh similar ages., It is thus legitimate to compare our brown dwarf disk masses with those of T Tauri stars in star forming regions with similar ages.
 The most appropriate wav lo do this is to compare the of disk mass to object mass., The most appropriate way to do this is to compare the of disk mass to object mass.
 We estimate masses for our brown dwarf targets by converting (heir spectral tvpes to elfective temperatures using the scale by Luhimanetal.(2008a) and comparing these temperatures with the most recent evolutionary (racks by Daralfeetal.(2003) assuming an age of MMvr., We estimate masses for our brown dwarf targets by converting their spectral types to effective temperatures using the scale by \citet{lsm03} and comparing these temperatures with the most recent evolutionary tracks by \citet{bcb03} assuming an age of Myr.
 The derived object masses are listed in Table L.., The derived object masses are listed in Table \ref{targets}.
 Although the uncertainties in (he conversion [rom spectral tvpes (ο masses are considerable. mainly due the evolutionary tracks at voung ages (seeDaraffeetal.2002).. all our targets are likely to have masses between 0.01 and AZ...," Although the uncertainties in the conversion from spectral types to masses are considerable, mainly due the evolutionary tracks at young ages \citep[see][]{bca02}, all our targets are likely to have masses between 0.01 and $\,M_{\odot}$."
 We note that the model uncertainties lead to svstematic errors. (hus (he relative masses in our sample are more reliable.," We note that the model uncertainties lead to systematic errors, thus the relative masses in our sample are more reliable."
" The same procedure was applied to derive masses for the voung brown cdwarfs and very low mass stars observed by. (2003).. which are included in the following analvsis. in cases that we did not observe in ΟΥ Survey,"," The same procedure was applied to derive masses for the young brown dwarfs and very low mass stars observed by \citet{kap03}, which are included in the following analysis, in cases that we did not observe in our survey."
 As comparison samples for higher mass stars. we used the results from Osterloh& (1995).. Nuernberger.Chini.&Zinnecker (1997).. Nuernbergerοἱal. (1993).. and Natta.Grinin.&Mannings(2000).," As comparison samples for higher mass stars, we used the results from \citet{ob95}, \citet{ncz97}, \citet{nby98}, and \citet{ngm00}."
. All these papers contain lists with object masses aud ages as well as disk masses determined either rom SED modeling or mnm measurements for objects mainly belonging to star forming regions in Taurus. p OOph. and Lupus.," All these papers contain lists with object masses and ages as well as disk masses determined either from SED modeling or mm measurements for objects mainly belonging to star forming regions in Taurus, $\rho$ Oph, and Lupus."
 Again. all object masses should be considered as rough estimates. but for our purposes even uncertainties of ave tolerable.," Again, all object masses should be considered as rough estimates, but for our purposes even uncertainties of are tolerable."
 To avoid being biased by an age spread (see above). we used only the subsample of stars with ages between 1 and MMyr. where disk masses are known not to show a significant trend with age.," To avoid being biased by an age spread (see above), we used only the subsample of stars with ages between 1 and Myr, where disk masses are known not to show a significant trend with age."
 All comparison stars can thus be considered to be coeval with our brown cdwarls., All comparison stars can thus be considered to be coeval with our brown dwarfs.
 The final sample of comparison sources comprises of 52 objects with masses between 0.03 and 2.7M... among them 25 upper and 2 lower limits.," The final sample of comparison sources comprises of 52 objects with masses between 0.03 and $\,M_{\odot}$, among them 25 upper and 2 lower limits."
 Their ratios of disk ancl object mass are plotted in Fig., Their ratios of disk and object mass are plotted in Fig.
 3. along with our values lor Taurus brown cdwarls., \ref{f3} along with our values for Taurus brown dwarfs.
 Fig., Fig.
 3. does not reveal a significant change of disk to object mass ratio with object mass., \ref{f3} does not reveal a significant change of disk to object mass ratio with object mass.
 Specificallv. there is no obvious difference in (he verv low mass regine. Le. among the lowest mass stars aud brown dwaarls. where ejection might become important (Goodwin.Whitworth.&Ward-Thompson 2004).," Specifically, there is no obvious difference in the very low mass regime, i.e., among the lowest mass stars and brown dwarfs, where ejection might become important \citep{gww04}."
. The average mass ratio (for the detections) is in the stellar and in the substellar regime. thus comparable given the large scatter and uncertainties.," The average mass ratio (for the detections) is in the stellar and in the substellar regime, thus comparable given the large scatter and uncertainties."
 There is no hint of an underabundaney of massive disks among brown cwarls: 413-1 (9 out of 22) detected disks around stars have ratios > 2%... whereas (here are (wo," There is no hint of an underabundancy of 'massive' disks among brown dwarfs: $41\pm 14$ (9 out of 22) detected disks around stars have ratios $>2$ , whereas there are two"
The observed properlies of dense cores BBenson Alvers 1989) form the basis of the standard model of isolated star formation.,The observed properties of dense cores Benson Myers 1989) form the basis of the standard model of isolated star formation.
" In this model. (he ""starless"" dense core represents the earliest identifiable stage of the star Formation process."," In this model, the “starless” dense core represents the earliest identifiable stage of the star formation process."
 The physical conditions in this early stage have a profound impact on Cie evolution of protostars towards (he main sequence., The physical conditions in this early stage have a profound impact on the evolution of protostars towards the main sequence.
 The iniüal density structure. particularly in the innermost regions. affects the collapse dvnanmies and the time dependence of the mass accretion rate and therefore many of the observable properties of protostars. including huminosity.," The initial density structure, particularly in the innermost regions, affects the collapse dynamics and the time dependence of the mass accretion rate and therefore many of the observable properties of protostars, including luminosity."
 A quantitative understanding of the collapse dynamics has been hindered by uncertain knowledge of appropriate initial conditions (André... Ward-Thompson Barsony 2000).," A quantitative understanding of the collapse dynamics has been hindered by uncertain knowledge of appropriate initial conditions (André,, Ward-Thompson Barsony 2000)."
 In the popular theory of “inside-out” collapse. a spherical starless core loses turbulent aud niagnelic support and relaxes to a balance between gravity. and thermal pressure. an r7 density. distribution is established and the core collapses from the inside-out wilh a constant mass accretion rate (Shu 1971).," In the popular theory of “inside-out” collapse, a spherical starless core loses turbulent and magnetic support and relaxes to a balance between gravity and thermal pressure, an $r^{-2}$ density distribution is established and the core collapses from the inside-out with a constant mass accretion rate (Shu 1977)."
 However. if collapse begins before the density distribution fully relaxes. then a central region of relatively constant density remains and the mass accretion rate is an order of magnitude larger at early Games (Foster Chevalier 1993).," However, if collapse begins before the density distribution fully relaxes, then a central region of relatively constant density remains and the mass accretion rate is an order of magnitude larger at early times (Foster Chevalier 1993)."
" This phenomenon has been identified with the voungest ""Class 0 protostars. which exhibit especially powerful outflows (Ilenriksen. Andre Bontemps 1997: Andre. Ward-Thompson Barsony 2000)."," This phenomenon has been identified with the youngest “Class 0” protostars, which exhibit especially powerful outflows (Henriksen, Andre Bontemps 1997; Andre, Ward-Thompson Barsony 2000)."
 Detter observations of starless cores are needed to determine the initial conditions., Better observations of starless cores are needed to determine the initial conditions.
of the Kochaneketal.(2001). GLE. combined with the strong bias towards luminous galaxies. the integrated cross-section of the galaxy population converges rapidly at faint magnitudes.,"of the \citet{2001ApJ...560..566K} GLF, combined with the strong bias towards luminous galaxies, the integrated cross–section of the galaxy population converges rapidly at faint magnitudes."
 Adopting a faint limit of A.|5 for the integration results in a prediction that 760 per cent of the absorber cross—section should result from galaxies brighter than our observational limit of AJ*|1.5. in good agreement with the fraction of our 30 absorbers where we see such a galaxy.," Adopting a faint limit of $M^* + 5$ for the integration results in a prediction that $\simeq$ per cent of the absorber cross–section should result from galaxies brighter than our observational limit of $M^* +1.5$, in good agreement with the fraction of our 30 absorbers where we see such a galaxy."
 Given the power law luminosity—-dependent scaling provides an effective and simple parameterisation of our results. we adopt the relation in order to make quantitative estimates of the filling-factor of the absorption and to compare the properties of the absorbers with other classes of quasar absorption Systems.," Given the power law luminosity–dependent scaling provides an effective and simple parameterisation of our results, we adopt the relation in order to make quantitative estimates of the filling-factor of the absorption and to compare the properties of the absorbers with other classes of quasar absorption systems."
 The A-band imaging results of the systems provide new constraints on the relationship between absorbers and their associated galaxies., The $K$ -band imaging results of the systems provide new constraints on the relationship between absorbers and their associated galaxies.
 In. Wild.Hewett.&Pettini(2007.Section4.1) we used observations of DLAs. together with the relative number densities of DLAs and absorbers. to propose that absorption can arise at projected radii of ~ I0Kkkpe from the centre of a galaxy.," In \citet[][Section 4.1]{2007MNRAS.374..292W} we used observations of DLAs, together with the relative number densities of DLAs and absorbers, to propose that absorption can arise at projected radii of $\sim10$ kpc from the centre of a galaxy."
 This simple calculation assumes a circular geometry for the absorbers and a unit filling factor (ie. the gas is not patchy)., This simple calculation assumes a circular geometry for the absorbers and a unit filling factor (i.e. the gas is not patchy).
 Clearly. this model contrasts with the much larger value of KKkpe derived from the A -band imaging.," Clearly, this model contrasts with the much larger value of kpc derived from the $K$ -band imaging."
 Our imaging results reveal that the dominant contributor to the cross-section is the inner part of more extended DLA absorbers. centred on relatively luminous galaxies.," Our imaging results reveal that the dominant contributor to the cross-section is the inner part of more extended DLA absorbers, centred on relatively luminous galaxies."
 Rather. the absorbers themselves trace much larger structures associated primarily with luminous galaxies where the filling factor is low.," Rather, the absorbers themselves trace much larger structures associated primarily with luminous galaxies where the filling factor is low."
 Taking the value of dP?/dz=0.025 from Wild.Hewett.Pettini(2007). for absorbers at 21 and integrating the luminosity-dependent absorber cross-section (Section 5.4.2)) using the Kochaneketa.(2001) GLF down to A.|5 produces a value of a=40)—Kkpc. equivalent to a radius. /?7=11.ΕΚΚΡοΟ.," Taking the value of $dP/dz=0.025$ from \citet{2007MNRAS.374..292W} for absorbers at $z$$\simeq$ 1 and integrating the luminosity-dependent absorber cross–section (Section \ref{sec:ldcc}) ) using the \citet{2001ApJ...560..566K} GLF down to $M^*+5$ produces a value of $\sigma^*=400$ $^2$, equivalent to a radius, $R^*=11.4$ kpc."
 Our observaions produce an observed mean galaxy-absorber separation of kkpe for galaxies with AJ=A. leading to a maximum radius out to which an absorber may be of kkpc.," Our observations produce an observed mean galaxy-absorber separation of kpc for galaxies with $M=M^*$, leading to a maximum radius out to which an absorber may be of kpc."
 The inferred filling-factor is thus only €11.4/36)7—0.1., The inferred filling-factor is thus only $^2$ =0.1.
 The large mean projected separations observed also provide an explanation for thesmall contribution of the absorbers to the star formation rate density at >~| determined by Wild.Hewett.&Pettini(2007). from AA3727.3730 emission observed in the SDSS spectra.," The large mean projected separations observed also provide an explanation for thesmall contribution of the absorbers to the star formation rate density at $z\sim1$ determined by \citet{2007MNRAS.374..292W} from $\lambda\lambda$ 3727,3730 emission observed in the SDSS spectra."
 The fraction of the A'-band luminosity of the associated galaxies that would fall within the 3700 diameter of the SDSS spectroscopic fibres is only 305 pper cent., The fraction of the $K$ -band luminosity of the associated galaxies that would fall within the 0 diameter of the SDSS spectroscopic fibres is only $\pm5$ per cent.
 Thus. the star formation associated with the central luminous parts of the associated galaxies would have been missed and. assuming the observed galaxies are directly responsible for the absorption. the results of Wild.Hewett.&Pettini(2007) do not constrain directly the contribution of absorber galaxies to the star formation rate of the Universe at 2~1.," Thus, the star formation associated with the central luminous parts of the associated galaxies would have been missed and, assuming the observed galaxies are directly responsible for the absorption, the results of \citet{2007MNRAS.374..292W} do not constrain directly the contribution of absorber galaxies to the star formation rate of the Universe at $z\sim1$."
 In this section we compare our imaging results to those for the more familiar classes of strong absorbers and DLAs., In this section we compare our imaging results to those for the more familiar classes of strong absorbers and DLAs.
 The equivalent widths of the absorbers are very large: the mean rest-frame ΕΕνοτορ of the sample imaged in this paper isAA.. with a minimum value ofΑΑ.. and 18 of the 30 systems have equivalent widths greater than (Table 6).," The equivalent widths of the absorbers are very large; the mean rest-frame $W_{\lambda2796}$ of the sample imaged in this paper is, with a minimum value of, and 18 of the 30 systems have equivalent widths greater than (Table \ref{tab:caiiprop}) )."
 Where possible. we therefore focus on results for similarly strong absorbers.," Where possible, we therefore focus on results for similarly strong absorbers."
" The only published galaxy luminosity function for absorbers is from the survey of 52 absorbers by Steidel.Dickinson.&Persson (1994).. who find that. after correcting for an observed dependence of gas cross-section on galaxy luminosity (67x L1), the Av-band luminosity function of absorbers in the redshift range 0.2<+1.0 is consistent with the :=0 luminosity function of Mobasher.Sharples.&Ellis (1993)."," The only published galaxy luminosity function for absorbers is from the survey of 52 absorbers by \citet{1994ApJ...437L..75S}, who find that, after correcting for an observed dependence of gas cross-section on galaxy luminosity $\sigma\propto L^{0.4}$ ), the $K$ -band luminosity function of absorbers in the redshift range $0.2\le z \le1.0$ is consistent with the $z$ =0 luminosity function of \citet{1993MNRAS.263..560M}. ."
. This implies a weaker galaxy luminosity-dependence of the cross-section than for absorbers., This implies a weaker galaxy luminosity-dependence of the cross-section than for absorbers.
 Several caveats apply to the comparison with Steidel et al., Several caveats apply to the comparison with Steidel et al.
's result.,'s result.
 Firstly. the equivalent width distribution of the absorber sample differs considerable from ours. with their systems selected to have M»zo570.3AA..," Firstly, the equivalent width distribution of the absorber sample differs considerable from ours, with their systems selected to have $W_{\lambda2796} >$."
 Secondly. a recent reassessment of the identitication procedure for host galaxies in this survey by Churchill.Steidel. (2005)... has revealed a small number of potential systematicbiases which remain to be investigated.," Secondly, a recent reassessment of the identification procedure for host galaxies in this survey by \citet{2005ASPC..331..387C}, , has revealed a small number of potential systematicbiases which remain to be investigated."
 Our second comparison is with the r-band imaging of the environments of 15 very strong. ΕΕντου22.7 AA..," Our second comparison is with the $r$ -band imaging of the environments of 15 very strong, $W_{\lambda2796} \ge 2.7$ ,"
under study.,under study.
" This lack of evolution contrasts with the strong evolution of major mergers, for which the relative fractions are ~5%//40%//55% at z=0.8 (similar to the minor ones), and ~10%//60%//30% at z=0.5."," This lack of evolution contrasts with the strong evolution of major mergers, for which the relative fractions are $\sim$ at $z = 0.8$ (similar to the minor ones), and $\sim$ at $z = 0.5$."
" From z~0.8 to z~0.5, the fraction of wet major mergers decreases by a factor of two, while dry and mixed mergers increase their importance."," From $z\sim0.8$ to $z\sim0.5$, the fraction of wet major mergers decreases by a factor of two, while dry and mixed mergers increase their importance."
" Our major merger trends are in agreement with using an expanded data set, as well as previous works, e.g.,??."," Our major merger trends are in agreement with using an expanded data set, as well as previous works, e.g.,."
". These results show that the relative fraction of dry and mixed major mergers become more important with cosmic time for Lg=Li galaxies in our redshift range due to the lack of blue primaries with major companions at low redshift, rather than from an increase in the major merger fractions of red galaxies as also pointed out by?."," These results show that the relative fraction of dry and mixed major mergers become more important with cosmic time for $L_{B} \gtrsim L_{B}^{*}$ galaxies in our redshift range due to the lack of blue primaries with major companions at low redshift, rather than from an increase in the major merger fractions of red galaxies as also pointed out by."
". Previous work finds that the major merger fraction from close pairs depends on mass, with more massive galaxies having higher merger fractions(??)."," Previous work finds that the major merger fraction from close pairs depends on mass, with more massive galaxies having higher merger fractions."
". If blue principal galaxies at z=0.8 were more massive by a factor of 3 than at z=0.5 because of our B—band luminosity selection, this would explain the observed trend in f>!""°."," If blue principal galaxies at $z = 0.8$ were more massive by a factor of 3 than at $z = 0.5$ because of our $B-$ band luminosity selection, this would explain the observed trend in $f_{\rm m}^{\rm blue}$."
" Using stellar masses determined in?,, we do not find a significant change (less than 0.1 dex) in the median mass of red, log(Mysea/Mo)~10.8, and blue, log(Myxbtue/Mo)~10.3, principal galaxies."," Using stellar masses determined in, we do not find a significant change (less than 0.1 dex) in the median mass of red, $\log\,(\overline{M_{\star, {\rm red}}}/M_{\odot}) \sim 10.8$, and blue, $\log\,(\overline{M_{\star, {\rm blue}}}/M_{\odot}) \sim 10.3$, principal galaxies."
 This supports that the observed trends reflect a real evolution in the merger properties of blue galaxies., This supports that the observed trends reflect a real evolution in the merger properties of blue galaxies.
" In addition, our results imply that more massive (red) galaxies have higher merger fractions than lower mass (blue) galaxies, in agreement with and?."," In addition, our results imply that more massive (red) galaxies have higher merger fractions than lower mass (blue) galaxies, in agreement with and."
". The study of the major and minor merger fraction in mass selected galaxies is beyond the scope of the present paper, and we will address this issue in a future work."," The study of the major and minor merger fraction in mass selected galaxies is beyond the scope of the present paper, and we will address this issue in a future work."
 Our goal in this section is to estimate the minor merger (1/10<pw 1/4) rate of bright galaxies in the range 0.2<z<0.95., Our goal in this section is to estimate the minor merger $1/10 \leq \mu < 1/4$ ) rate of bright galaxies in the range $0.2 \leq z < 0.95$ .
" In the following we name the the number of mergers per Gyr per galaxy, noted R."," In the following we name the the number of mergers per Gyr per galaxy, noted $R$."
" Because the parameters involved in the translation of the merger fraction to the merger rate are better constrained for major mergers, we estimate them first and then expand to the minor merger rate."," Because the parameters involved in the translation of the merger fraction to the merger rate are better constrained for major mergers, we estimate them first and then expand to the minor merger rate."
" Following?,, we define the major merger rate as where the factor Cj takes into account the lost companions in the inner 5/7! kpc and the factor Οι is the fraction of the observed close pairs that finally merge in a typical timescale Tym."," Following, we define the major merger rate as where the factor $C_{\rm p}$ takes into account the lost companions in the inner $5h^{-1}$ kpc and the factor $C_{\rm m}$ is the fraction of the observed close pairs that finally merge in a typical timescale $T_{\rm MM}$."
" We take Cp=r5""—5h""!kpc).", We take $C_{\rm p} = r_{\rm p}^{\rm max}/(r_{\rm p}^{\rm max}-5h^{-1}\ {\rm kpc})$.
" The typical merger timescale depends on /($""*γρ and can be estimated by cosmological and N-body simulations.", The typical merger timescale depends on $r_{\rm p}^{\rm max}$ and can be estimated by cosmological and $N$ -body simulations.
" We compute the major merger timescales from the cosmological simulations of ?,, based on the Millennium simulation(?)."," We compute the major merger timescales from the cosmological simulations of , based on the Millennium simulation."
". These major merger timescales, denoted TEM refer to major mergers (u>1/4 in stellar mass), and depend mainly οπής” and on the stellar mass of the principal galaxy, with a weak dependence on redshift in our range of interest (see?,, for details)."," These major merger timescales, denoted $T_{\rm MM}^{K08}$, refer to major mergers $\mu > 1/4$ in stellar mass), and depend mainly on $r_{\rm p}^{\rm max}$ and on the stellar mass of the principal galaxy, with a weak dependence on redshift in our range of interest (see, for details)."
" Taking log(M,/Mo)10.7 as the average stellar mass of our principal galaxies with a close companion, we obtain the values in Table 7 for r5*=30, 50 and 100 4! kpc, and Av""**=500 km s! "," Taking $\log\, (M_{\star}/M_{\odot}) = 10.7$ as the average stellar mass of our principal galaxies with a close companion, we obtain the values in Table \ref{tmm} for $r_{\rm p}^{\rm max} = 30$, 50 and 100 $h^{-1}$ kpc, and $\Delta v^{\rm max} = 500$ km $^{-1}$ ."
In every case we assume an uncertainty of 0.2 dex in the mass of the principal galaxies to estimate the error in TEM., In every case we assume an uncertainty of 0.2 dex in the mass of the principal galaxies to estimate the error in $T_{\rm MM}^{K08}$.
" These timescales already include the factor Cy,???),, so we take Cy,=1 in the following."," These timescales already include the factor $C_{\rm m}$, so we take $C_{\rm m} = 1$ in the following."
" These timescales are for central - satellite mergers, and satellite - satellite pairs could have different timescales."," These timescales are for central - satellite mergers, and satellite - satellite pairs could have different timescales."
" However, only 1 of the 103 close pairs under study is satellite - satellite, so the use of principal - satellite timescales is justified."," However, only 1 of the 103 close pairs under study is satellite - satellite, so the use of principal - satellite timescales is justified."
" We also remark that the velocity condition Av™*=500 km s! selects close bound systemseven when they are located in denseenvironments, but in these environments the probability"," We also remark that the velocity condition $\Delta v^{\rm max} = 500$ km $^{-1}$ selects close bound systemseven when they are located in denseenvironments, but in these environments the probability"
The rate of growth of the bubble radius near the critical radius is given by HH. where fis the cdvnamical prefactor 1992a).,"The rate of growth of the bubble radius near the critical radius is given by $\left( \frac{dr}{dt}\right) _{macro}=k\left( r-R_{c}\right) $ , where $k$ is the dynamical prefactor ."
. The prefactor A has been evaluated. by solving the equations of relativistic fIuid dynamics in all regions. in(2001).," The prefactor $k$ has been evaluated, by solving the equations of relativistic fluid dynamics in all regions, in."
". The result (also taking into account heat conduction) is where c, is a constant (the velocity of the sound in the mediun around the saddle configuration) and A,,. 75, and £, are the thermal conductivity and shear and bulk viscosity coellicients of the neutron matter. respectively."," The result (also taking into account heat conduction) is where $c_{s}$ is a constant (the velocity of the sound in the medium around the saddle configuration) and $\lambda _{n}$, $\eta _{n}$ and $\xi _{n}$ are the thermal conductivity and shear and bulk viscosity coefficients of the neutron matter, respectively."
 The first term in the above equation is the same as obtained by(1996).. corresponding to the case of non-viscous malter.," The first term in the above equation is the same as obtained by, corresponding to the case of non-viscous matter."
 The second term is similar to the results obtained by and(1994).. except will a minor dillerence . i.e. instead of 4 there is a factor c? in the numerator.," The second term is similar to the results obtained by and, except with a minor difference , i.e. instead of $4$ there is a factor $c_{s}^{2}$ in the numerator."
 With the use of Eq. (24)), With the use of Eq. \ref{pre}) )
 for the prefactor it follows that the rate of Formation of quark bubbles in neutron matter is given by To obtain the number € of the net strange quark matter bubbles. the rate must be multiplied bv the time interval available for prompt nucleation. A’ and bv the volume Vj. where the nucleation can take place in (he dense core1992).," for the prefactor it follows that the rate of formation of quark bubbles in neutron matter is given by To obtain the number $\xi $ of the net strange quark matter bubbles, the rate $j$ must be multiplied by the time interval available for prompt nucleation, $\Delta t$ and by the volume $V_{0}$, where the nucleation can take place in the dense core."
. We impose that at least one quark bubble appears (which would suffice to convert thewholeneutron star)., We impose that at least one quark bubble appears (which would suffice to convert thewholeneutron star).
 Therefore. with (he useof Eq. (25)).," Therefore, with the useof Eq. \ref{j}) ),"
 one obtains the following, one obtains the following
any photo-z algorithm.,any $z$ algorithm.
 Of course. if this method is to work. then the subsample with spectral information be able to provide an accurate estimate of p(z[e).," Of course, if this method is to work, then the subsample with spectral information be able to provide an accurate estimate of $p(z|{\bm c})$."
 The convolution method. of the previous subsection provides a simple way of illustrating how one should use the output from photo-z codes that actually provide a properly calibrated probability distribution. Z(z|e) for cach set. of colors €. to estimate dNels.," The convolution method of the previous subsection provides a simple way of illustrating how one should use the output from $z$ codes that actually provide a properly calibrated probability distribution ${\cal L}(z|{\bm c})$ for each set of colors ${\bm c}$, to estimate ${\rm d}N/{\rm d}z$."
" I also shows in what sense the codes should be ""unbiased.", It also shows in what sense the codes should be `unbiased'.
 In particular. equation (5)) suggests that one can estimate cLN(z)/dz by summing over all the objects in the dataset. weighting cach by its £(z[e).," In particular, equation \ref{Npz|c}) ) suggests that one can estimate ${\rm d}N(z)/{\rm d}z$ by summing over all the objects in the dataset, weighting each by its ${\cal L}(z|{\bm c})$ ."
 This is because Equation (6)) shows that if Z(z|e) does not have the same shape as p(z|e). then use of Z(z]e) will lead to a bias: this is the pernicious bias which must. be reduced whether or not (z|e) equals the spectroscopic redshift is. in some sense. irrelevant. (," This is because Equation \ref{NLz|c}) ) shows that if ${\cal L}(z|{\bm c})$ does not have the same shape as $p(z|{\bm c})$, then use of ${\cal L}(z|{\bm c})$ will lead to a bias; this is the pernicious bias which must be reduced – whether or not $\langle z|{\bm c}\rangle$ equals the spectroscopic redshift is, in some sense, irrelevant. ("
In the case of a one-to-one mapping between e and C. (z|e? is the same as the quantity (z[O which we discussed in the previous subsections.),"In the case of a one-to-one mapping between ${\bm c}$ and $\zeta$, $\langle z|{\bm c}\rangle$ is the same as the quantity $\langle z|\zeta\rangle$ which we discussed in the previous subsections.)"
 Satisfving Z(z|e)=p(z|e) is nontrivial., Satisfying ${\cal L}(z|{\bm c}) = p(z|{\bm c})$ is nontrivial.
 This is perhaps most easily seen by supposing that the template or training set consists of two galaxy types (earlv- and Iate-tvpes. sav). for which the same observed colors are associated with two different redshifts.," This is perhaps most easily seen by supposing that the template or training set consists of two galaxy types (early- and late-types, say), for which the same observed colors are associated with two different redshifts."
 In this case. if the photo-z: algorithms are working well. then Z(z|e) will be bimodal for at least some ο.," In this case, if the $z$ algorithms are working well, then ${\cal L}(z|{\bm c})$ will be bimodal for at least some ${\bm c}$."
 Llowever. if the sample of interest. only contains LRGs. then p(z|e) may actually be unimocal.," However, if the sample of interest only contains LRGs, then $p(z|{\bm c})$ may actually be unimodal."
 As a result. C(z|e)xp(z|e) unless proper priors on the templates are used. or care has been taken to insure that the training set is representative of the sample of interest.," As a result, ${\cal L}(z|{\bm c}) \ne p(z|{\bm c})$ unless proper priors on the templates are used, or care has been taken to insure that the training set is representative of the sample of interest."
 We can perform a similar analysis of the luminosity function., We can perform a similar analysis of the luminosity function.
 In this case. the key is to recognize that. in a magnitude limited survey. the quantity which is most directly.allected by the photometric redshift’ error is not the luminosity function. ολ) itsel but. the luminosity," In this case, the key is to recognize that, in a magnitude limited survey, the quantity which is most directlyaffected by the photometric redshift error is not the luminosity function $\phi(M)$ itself, but the luminosity"
"with T7, =lol"" JaslongasthemagneticfieldislessthanlOnoftheequipartilion hegaspressurebeingequallolhemagnelicpr 310!",with T_1 = as long as the magnetic field is less than of the equipartition (the gas pressure being equal to the magnetic pressure) field.
 weindthaltlhem, The outer boundary of the flow and the opening angle of the outflow are hardly changed if this condition is met.
agnelicfieldhastobesmallerthanattoflheequiparlition[ieldloproducelheoul [low.," For T_1 = 3, we find that the magnetic field has to be smaller than of the equipartition field to produce the outflow."
For flow. Iheproductionoftheoul flow/spossibleontywhenthemagneticfieldislessthanl," For the highest temperature T_1 = flow, the production of the outflow is possible only when the magnetic field is less than of the equipartition field."
Woflheequipartition |, So we conclude that the Compton preheated outflow is possible in the magnetic CDAFs as long as the magnetic field is less than from to of the equipartition field.
 Tot aceretion flows like ADAFs have a number of physical characteristics that compliment the classic low-temperature disk flows., Hot accretion flows like ADAFs have a number of physical characteristics that compliment the classic low-temperature disk flows.
 In previous work. we have explored (he consequences of the high temperature and the tvo-dimensional density structure of these flows. and have found that ADAFs may be able to produce raclialively driven outflows.," In previous work, we have explored the consequences of the high temperature and the two-dimensional density structure of these flows, and have found that ADAFs may be able to produce radiatively driven outflows."
 We subsequently have noted that self-similar CDAFs have even more suitable properties for producing outflows as compared to ADAFs: steeper poloidal density gradients and higher radiation efficiencies., We subsequently have noted that self-similar CDAFs have even more suitable properties for producing outflows as compared to ADAFs: steeper poloidal density gradients and higher radiation efficiencies.
 In (his paper. we have studied the conditions for self-similar two-dimensional CDAFs (NIA: QG) to develop radiativelv heated polar outflows.," In this paper, we have studied the conditions for self-similar two-dimensional CDAFs (NIA; QG) to develop radiatively heated polar outflows."
 1., 1.
 We have found that CDAFs produce enough luminosity and photon energy to drive polar outflows via Compton heating lor a reasonable range of mass accretion rate. or. equally. luminosity. as long as the magnetic field is less than from to of the equipartition field.," We have found that CDAFs produce enough luminosity and photon energy to drive polar outflows via Compton heating for a reasonable range of mass accretion rate, or, equally, luminosity, as long as the magnetic field is less than from to of the equipartition field."
 When the electron temperature saturates around. 1011IN at the inner region. polar outflows are possible lor 8x10.*<L/LyS4x10? for radiation efficiency of 10.7. where Ly is the Eddington humninositv.," When the electron temperature saturates around $10^{11}\K$ at the inner region, polar outflows are possible for $8\times10^{-7} \lesssim L/L_E
\lesssim 4\times10^{-5}$ for radiation efficiency of $10^{-2}$, where $L_E$ is the Eddington luminosity."
" The huninositv range for whieh outflow exists is narrower for lower electron temperature flows and disappears completely [or electron temperature <6x10""IX 2."," The luminosity range for which outflow exists is narrower for lower electron temperature flows and disappears completely for electron temperature $\lesssim 6\times
10^9\K$ 2."
 In most cases. oulllows are well collimated along the rotation axis. wilh an opening angle tvpicallv in the range <107.," In most cases, outflows are well collimated along the rotation axis, with an opening angle typically in the range $\lesssim 10^\circ$."
 3., 3.
 If we. instead of taking elliciency as constant. assume (hat it is proportional to the nass accretion rate 1. the solutions are qualitatively the same but are more self similar. 1.e.. opening angle and outer boundary. Gu Schwarzschild units) depend less strongly on i.," If we, instead of taking efficiency as constant, assume that it is proportional to the mass accretion rate $\dot m$, the solutions are qualitatively the same but are more self similar, i.e., opening angle and outer boundary (in Schwarzschild units) depend less strongly on $\dot m$."
 4., 4.
 Outflow is more probable for small viscosity parameter)., Outflow is more probable for small viscosity parameter.
 The treatment in this work is not completely satisfactory in the sense that the dvnamics and the temperature profiles of CDAFs are not sell-consistently solved., The treatment in this work is not completely satisfactory in the sense that the dynamics and the temperature profiles of CDAFs are not self-consistently solved.
 However. it was not," However, it was not"
center of inass of pulsar B. aud therefore applies a torque sufficient to produce the observed spin augular momenutum.,"center of mass of pulsar B, and therefore applies a torque sufficient to produce the observed spin angular momentum."
 The kick aud offset. vectors iust lie iu the plane perpendicular to AS (see Equation 3))., The kick and offset vectors must lie in the plane perpendicular to $\Delta \vec{S}$ (see Equation \ref{eq:spin-change}) ).
) The kick velocity. vy. the offset. vector relative to the center of mass. r. aud the chauge in B's spin vector are related by where Ap=A/pvyy is the change in linear moimeutuim induced by a chauge tn velocity of vy iu au object with mass Mj.," The kick velocity, $\vec{v}_K$, the offset vector relative to the center of mass, $\vec{r}$, and the change in B's spin vector are related by where $\Delta \vec{p} = M_B\vec{v}_K$ is the change in linear momentum induced by a change in velocity of $\vec{v}_K$ in an object with mass $M_B$."
 The offset leneth r and kick velocity magnitude ey must then satisfy tlie inequality where AS is the inaguitude of the chauge of B's spin: the equality holds ouly when the kick aud offset are perpendicular to each other., The offset length $r$ and kick velocity magnitude $v_K$ must then satisfy the inequality where $\Delta \vec{S}$ is the magnitude of the change of B's spin; the equality holds only when the kick and offset are perpendicular to each other.
 The relative orientation of the current spin provides a coustraint on the kick direction iu this scenario., The relative orientation of the current spin provides a constraint on the kick direction in this scenario.
" Let the colatitude of AS relative to tlie pre-SN orbital plane be 04: because the auele between S, and Sga is greater than 90 degrees. no matter the maeuitucde of Sij we must have Ox>116 degrees."," Let the colatitude of $\Delta
\vec{S}$ relative to the pre-SN orbital plane be $\theta_{\Delta}$; because the angle between $\vec{S}_0$ and $\vec{S}_{SN}$ is greater than 90 degrees, no matter the magnitude of $\vec{S}_0$ we must have $\theta_{\Delta} \geq 116$ degrees."
 Let the plane perpendicular to AS make an angle  with respect to the pre-SN orbital plane., Let the plane perpendicular to $\Delta \vec{S}$ make an angle $\psi$ with respect to the pre-SN orbital plane.
 Then we have ο=180—04<cU61 degrees., Then we have $\psi = 180 - \theta_{\Delta} \leq 64$ degrees.
 The kick. vj. lies in this plane aud therefore must have colatitude Oy that satisfies This geometry is illustrated in Figure 2..," The kick, $\vec{v}_K$, lies in this plane and therefore must have colatitude $\theta_K$ that satisfies This geometry is illustrated in Figure \ref{fig:spin-geometry}."
 The coustraint in Exquatiou 5 is consistent. with the constraint on kick colatitude in Figure 7 of (Wongetal.2010).. but tighter.," The constraint in Equation \ref{eq:theta-constraint} is consistent with the constraint on kick colatitude in Figure 7 of \citep{Wong2010}, but tighter."
 The coustraiut on the magnitude of the spin change. Equation(2).. together with Equation(1).. imply a lower limit ou the offset distance [If we assume that the core of pulsar B's progenitor was in syuchronous. rigic-bocdy rotation just before the supernova then seο.κLOYeen’s! and the limit ou the offset distance rises by a [actor of 1.5: Wongetal.(2010) used the measured semi-inajor axis. eccentricity and proper motion of the JOT37-3039 system (see Table 1)) to coustrain the kick imparted to the system by the secoud SN.," The constraint on the magnitude of the spin change, Equation, together with Equation, imply a lower limit on the offset distance If we assume that the core of pulsar B's progenitor was in synchronous, rigid-body rotation just before the supernova then $S^{SR}_0 \simeq 2 \times 10^{45}\, \textnormal{g}\, \textnormal{cm}^2\,
\textnormal{s}^{-1}$, and the limit on the offset distance rises by a factor of $1.8$: \citet{Wong2010} used the measured semi-major axis, eccentricity and proper motion of the J0737-3039 system (see Table \ref{tab:system-parameters}) ) to constrain the kick imparted to the system by the second SN."
"ας coming from the intracluster gas heats the clouds located in the cooling Low at a distance r from the cluster center. and so thermal balance between heating and. cooling defines an equilibrium temperature of the sub-clouds at a distance ro—[es inside the cooling How region (i.e. fu<recu). The following column densities are adopted. for a typical small cloud. (with ng,=10 οι and the orto-para ratio equal to Neo=108em?andNg,2107 ?. which corresponds to a CO abundance: geo~510 7.","flux coming from the intracluster gas heats the clouds located in the cooling flow at a distance $r$ from the cluster center, and so thermal balance between heating and cooling defines an equilibrium temperature of the sub-clouds at a distance $r=R_{eq}$ inside the cooling flow region (i.e. $R_{eq}<r_{cool}$ The following column densities are adopted for a typical small cloud (with $n_{H_2}=10^6$ $^{-3}$ and the orto-para ratio equal to $N_{CO}=10^{14}\, {\rm cm}^{-2} \, \, {\rm and} \, \, 
N_{H_2}=2 \times 10^{18} \, {\rm cm}^{-2}$ , which corresponds to a $CO$ abundance: $\eta_{CO}\sim 5 \times 10^{-5}$ ."
" For ££« instead. we assume the primordial ratio ""HDo4sLO”. In Figure 1 we plotted the equilibrium temperature of clumps at the equilibrium distance {ή for different values of the attenuation factor 7.", For $HD$ instead we assume the primordial ratio $\eta_{HD} \sim 7 \times 10^{-5}$ In Figure 1 we plotted the equilibrium temperature of clumps at the equilibrium distance $R_{eq}$ for different values of the attenuation factor $\tau$.
 We see that. low equilibrium temperatures are achieved at distances smaller than the cooling radius Veal., We see that low equilibrium temperatures are achieved at distances smaller than the cooling radius $r_{cool}$.
 In Figure 2 we plotted the equilibrium distance as a function of the equilibrium temperature for cillerent values of geoand 7=2.5is kept fixed., In Figure 2 we plotted the equilibrium distance as a function of the equilibrium temperature for different values of $\eta_{CO}$and $\tau=2.5$is kept fixed.
 Indeed. the CO ," Indeed, the $CO$ "
Consider the PRE to be a square of angular size PA/D.,Consider the PRF to be a square of angular size $P\lambda/D$.
 Figure 3.) plots the value of Asa as a function of P.," Figure \ref{ASNfig}
 plots the value of $A_{SN}$ as a function of $P$ ."
 In the limit P«I of fine sampling. Asgx reduces to that of the Airy pattern. (for pupil obscuration e= 0.25).," In the limit $P\ll1$ of fine sampling, $A_{SN}$ reduces to that of the Airy pattern, $A_{\rm Airy}=3.35(\lambda/D)^2$ (for pupil obscuration $\epsilon=0.25$ )."
 The sky noise is thus equivalent to that in a circular aperture of radius 1.053A/D., The sky noise is thus equivalent to that in a circular aperture of radius $1.03\lambda/D$ .
 In the limit Po>1. dsx is simply the pixel area Apis=P?(A/Dy*.," In the limit $P\gg1$, $A_{SN}$ is simply the pixel area $A_{\rm pix}=P^2(\lambda/D)^2$ ."
 It should be noted. however. that the common heuristic approximation Ayazmda;+Apis will Asgx. aud hence the required exposure time. by up to between these limits.," It should be noted, however, that the common heuristic approximation $A_{SN}\approx A_{\rm
Airy} + A_{\rm pix}$ will $A_{SN}$, and hence the required exposure time, by up to between these limits."
 lu particular. note that Nyquist-sampline pixels (22= 0.5) degrade the Airy Asa level by13%.. while pixels at the Airy FWHAIL of P?=1.22 degrade the speed by about a factor 1.5. asstunine interlacing to the Nyquist level.," In particular, note that Nyquist-sampling pixels $P=0.5$ ) degrade the Airy $A_{SN}$ level by, while pixels at the Airy FWHM of $P=1.22$ degrade the speed by about a factor 1.5, assuming interlacing to the Nyquist level."
 The degradation of point-source S/N is more severe if the detector has significant cdilfusion of charge before collection iuto pixels., The degradation of point-source $S/N$ is more severe if the detector has significant diffusion of charge before collection into pixels.
 The dashed line in Figure 3. shows Asa when there is Ciaussiau charge diffusion with σ of one-half pixel., The dashed line in Figure \ref{ASNfig} shows $A_{SN}$ when there is Gaussian charge diffusion with $\sigma$ of one-half pixel.
 Iu this case even the Nyquist-sizecl pixels increase Aga by30%... and the sv for P=1 is z10 pixels. three times worse than the pure Airy value.," In this case even the Nyquist-sized pixels increase $A_{SN}$ by, and the $A_{SN}$ for $P=1$ is $\approx10$ pixels, three times worse than the pure Airy value."
 On the other haud. the overall speed of a photometry project can beau increasing Πίο of pixel sizeif thenumber of pixels is coustrained.," On the other hand, the overall speed of a photometry project can bean increasing function of pixel sizeif thenumber of pixels is constrained."
 If the science goals require surveyiug a fixed. large," If the science goals require surveying a fixed, large"
Mik 421 is one of the brightest extragalactic X-ray sources in the sky.,Mrk 421 is one of the brightest extragalactic X-ray sources in the sky.
 It belongs to the eroup of DL Lac objects. a sub-class of radio-loud active galactic nuclei (AGN).," It belongs to the group of BL Lac objects, a sub-class of radio-loud active galactic nuclei (AGN)."
 The BL Lac, The BL Lac
If (he rotational Irequency is much smaller (han (he Ixepler frequency. the deviations from spherical sviumetry are negligible and the moment of inertia can be calculated applving the slow-rotation approximation discussed. briefly in Section 3.,"If the rotational frequency is much smaller than the Kepler frequency, the deviations from spherical symmetry are negligible and the moment of inertia can be calculated applying the slow-rotation approximation discussed briefly in Section 3."
 For this case Schutz(2005) showed that the moment of inertia can be very well approximated by Eq. (21))., For this case \citet{Lattimer:2005} showed that the moment of inertia can be very well approximated by Eq. \ref{eq.14}) ).
 In Fig., In Fig.
 4 we display the moment of inertia as a function of stellar mass lor slowly rotating neulron stars as computed with the empirical relation (21))., \ref{fig.4} we display the moment of inertia as a function of stellar mass for slowly rotating neutron stars as computed with the empirical relation \ref{eq.14}) ).
 As shown in Fig. 3..," As shown in Fig. \ref{fig.3},"
 above (he neutron star racius remains approximately constant before reaching (he maximum lass supported by a given EOS., above $\sim 1.0M_{\sun}$ the neutron star radius remains approximately constant before reaching the maximum mass supported by a given EOS.
 The moment of inertia (2~ALR?) (hus increases almost linearly with stellar mass for all models., The moment of inertia $I\sim MR^2$ ) thus increases almost linearly with stellar mass for all models.
 Right before (he maximum mass is achieved. the neutron star radius starts to decrease (Fig. 3)).," Right before the maximum mass is achieved, the neutron star radius starts to decrease (Fig. \ref{fig.3}) ),"
 which causes the sharp drop in the moment of inerlia observed in Fie. 4., which causes the sharp drop in the moment of inertia observed in Fig. \ref{fig.4}.
 Since { is proportional to the mass and the square of the racius. il is more sensitive to the density dependence of (he nuclear sviumetry energy. which determines (he neutron star radius.," Since $I$ is proportional to the mass and the square of the radius, it is more sensitive to the density dependence of the nuclear symmetry energy, which determines the neutron star radius."
 Here we recall that the ο=—1 EOS has much stiller symmetry energy (will respect to the one of the c=0 EOS). which results in neutron star models with larger radii and. in turn. momenta of inertia.," Here we recall that the $x=-1$ EOS has much stiffer symmetry energy (with respect to the one of the $x=0$ EOS), which results in neutron star models with larger radii and, in turn, momenta of inertia."
 For instance. for a “canonical” neutron star (.A/=LAA. ). the difference in the moment of inertia is more than 30% with the.=0 and the .c=—1 EOSs.," For instance, for a “canonical” neutron star $M=1.4M_{\sun}$ ), the difference in the moment of inertia is more than $30\%$ with the $x=0$ and the $x=-1$ EOSs."
 In Fig., In Fig.
" 5. we take another view of the moment of inertia where J is scaled by AL""? as a function of the stellar mass (after (Lattimer&Schutz 2005))).", \ref{fig.5} we take another view of the moment of inertia where $I$ is scaled by $M^{3/2}$ as a function of the stellar mass (after \citep{Lattimer:2005}) ).
 The discovery of the extremely relativistic binary. pulsar PSR J0737-3039A.D. provides an unprecedented opportunitv to test. General Relativity and physics of pulsars (Bureayal.," The discovery of the extremely relativistic binary pulsar PSR J0737-3039A,B provides an unprecedented opportunity to test General Relativity and physics of pulsars \citep{Burgay2003}."
 2003).. Lattimer&Schutz(2005). estimated that the moment of inertia of the A component of the svstem should be measurable with an accuracy of about1056., \citet{Lattimer:2005} estimated that the moment of inertia of the A component of the system should be measurable with an accuracy of about.
.. Given that the masses of both stars are already. accurately determined by observations. a measurement of (he moment of inertia of even one neutron star could have enormous importance for the neutron star physics (Lattimer&Schutz2005).. (," Given that the masses of both stars are already accurately determined by observations, a measurement of the moment of inertia of even one neutron star could have enormous importance for the neutron star physics \citep{Lattimer:2005}. ("
The significance of such a measurement is illustrated in Fig. 5..,The significance of such a measurement is illustrated in Fig. \ref{fig.5}.
 As pointed by Lattimer&Schutz(2005).. it is clear that very few EOSs would survive these constraints.)," As pointed by \citet{Lattimer:2005}, it is clear that very few EOSs would survive these constraints.)"
" Thus. theoretical predictions of the moment of inertia are verv timely,"," Thus, theoretical predictions of the moment of inertia are very timely."
 Caleulations of the moment of inertia of pulsar A (Ay=1.338M... vy= 44.0512) have been reported by Morrisonetal.(2004) and Dejgeretal.(2005).," Calculations of the moment of inertia of pulsar A $M_A=1.338M_{\sun}$, $\nu_A=44.05Hz$ ) have been reported by \citet{Morrison:2004} and \citet{BBH2005}."
. In Table 3. we show the moment of inertia and (other selected quantities) of PSR JOT37-3039A computed with the RAS code using the EOSs emploved in this study., In Table \ref{tab.3} we show the moment of inertia and (other selected quantities) of PSR J0737-3039A computed with the $RNS$ code using the EOSs employed in this study.
" Our results with the APR EOS are in very good agreement with those by Morrisonetal.(2004) (PPP—1,24x10Pg em?) and Bejgeretal.(2005). (LHP=1.23x107g em).", Our results with the APR EOS are in very good agreement with those by \citet{Morrison:2004} $I^{APR}=1.24\times 10^{45}g$ $cm^2$ ) and \citet{BBH2005} $I^{APR}=1.23\times 10^{45}g$ $cm^2$ ).
 In the last column of Table 3. we also include results computed with the empirical relation (Eq. (21)))., In the last column of Table \ref{tab.3} we also include results computed with the empirical relation (Eq. \ref{eq.14}) )).
 From a comparison with the results from (he exact numerical caleulation we conclude that Eq. (21)), From a comparison with the results from the exact numerical calculation we conclude that Eq. \ref{eq.14}) )
 is an excellent approximation for the moment of inertia of slowly-rotating neutron stars. (, is an excellent approximation for the moment of inertia of slowly-rotating neutron stars. (
The average uncertainty of Eq. (21)),The average uncertainty of Eq. \ref{eq.14}) )
 is ~ 2%. except lor the DBIIF--DonnD EOS for," is $\sim 2\%$ , except for the DBHF+BonnB EOS for"
"miergers are very Common in massive galaxies. with nearly all galaxies at AZ,>Lott hhaving a minor merger occurring within the past 0. Cyr at every redshift.","mergers are very common in massive galaxies, with nearly all galaxies at $M_{\star} > 10^{11}$ having a minor merger occurring within the past 0.4 Gyr at every redshift."
 At lower stellar masses. there is an increasingly [large dillerence in the minor merecr lraction evolution. with a high ratio of fuiawiner2709 at 2=3. which declines rapidly: with redshift by an order of magnitude at >=0.," At lower stellar masses, there is an increasingly large difference in the minor merger fraction evolution, with a high ratio of $f_{\rm m,minor} > 0.3$ at $z = 3$, which declines rapidly with redshift by an order of magnitude at $z=0$."
 This means that in the Millennium simulation Less massive galaxies undergo many more minor mergers at high redshift than in the recent. past., This means that in the Millennium simulation less massive galaxies undergo many more minor mergers at high redshift than in the recent past.
 One could argue that i£ the mass resolution of numerical simulations could be increased indefinitely. the minor merger fraction would be unity for all galaxies.," One could argue that if the mass resolution of numerical simulations could be increased indefinitely, the minor merger fraction would be unity for all galaxies."
 This is certainly true if one considers as minor merecrs all accretion events. independently of the mass ratio between the central galaxy and the infalling cloud.," This is certainly true if one considers as minor mergers all accretion events, independently of the mass ratio between the central galaxy and the infalling cloud."
 As we show in Fig. 2..," As we show in Fig. \ref{mass},"
 between SO and 90 per cent of mergers in the Millennium. simulation have mass ratios larger than 0.01., between 80 and 90 per cent of mergers in the Millennium simulation have mass ratios larger than 0.01.
 This mass ratio is a small number and includes accretion of small satellites that clo not alfect the morphology. of the central galaxy ancl would not be found as members of galaxy pairs., This mass ratio is a small number and includes accretion of small satellites that do not affect the morphology of the central galaxy and would not be found as members of galaxy pairs.
" A fraction of these merging satellites have stellar masses below the resolution of the simulation. especially when galaxies with AL«101"" aare considered."," A fraction of these merging satellites have stellar masses below the resolution of the simulation, especially when galaxies with $M_{\star} < 10^{10}$ are considered."
" However. above M,1027AL... most events we consider as minor mergers have a non-negligible mass ratio."," However, above $M_{\star} > 10^{10}$, most events we consider as minor mergers have a non-negligible mass ratio."
 Iniposing a minimum value for the mass ratio for mergers to be considered in our analysis. as for example a minimum mass ratio of 0.05. would mocifyv the minor merger results by about a factor of 2 for the most massive galaxies. but would not change the qualitative behaviour of the minor merger fractions and rates.," Imposing a minimum value for the mass ratio for mergers to be considered in our analysis, as for example a minimum mass ratio of 0.05, would modify the minor merger results by about a factor of 2 for the most massive galaxies, but would not change the qualitative behaviour of the minor merger fractions and rates."
 The redshift’ evolution of the minor merger fraction as a function of stellar mass shown in Fig., The redshift evolution of the minor merger fraction as a function of stellar mass shown in Fig.
 3 is an indication that resolution ellects alone do not shape the distribution of the minor merger fractions as a function of stellar mass., \ref{fig2} is an indication that resolution effects alone do not shape the distribution of the minor merger fractions as a function of stellar mass.
 According to Fig. 2..," According to Fig. \ref{mass},"
 major mergers represent between 5 and LO per cent of all merger events in the Millennium., major mergers represent between 5 and 10 per cent of all merger events in the Millennium.
 Fig., Fig.
 4 shows predictions of the merger rate as a function of stellar mass in units of 5 +., \ref{fig0} shows predictions of the merger rate as a function of stellar mass in units of $^{-3}$ $^{-1}$.
 The. evolution of the merger rate with redshift is shown for both major mergers (left panel) and minor mergers (right. panel)., The evolution of the merger rate with redshift is shown for both major mergers (left panel) and minor mergers (right panel).
 The cillerent lines indicate results at different redshifts., The different lines indicate results at different redshifts.
 The merger rate tends to be higher for less massive galaxies and lower for the most massive galaxies., The merger rate tends to be higher for less massive galaxies and lower for the most massive galaxies.
 Phe shape of the merger rate as a function of stellar mass mostly reflects the shape of the stellar mass function mai(z). which enters the definition of the merger rate according to Eq. 4..," The shape of the merger rate as a function of stellar mass mostly reflects the shape of the stellar mass function $n_{\rm gm} (z)$, which enters the definition of the merger rate according to Eq. \ref{mergerrate}."
" In fact. while the number density. of galaxies. varies by up to five orders of magnitude for 10 <Ad,1017 systems. the merger fractions vary by at most a factor of 100. as shown in Fig. 3.."," In fact, while the number density of galaxies varies by up to five orders of magnitude for $10^9$ $<M_{\star} < 10^{12}$ systems, the merger fractions vary by at most a factor of 100, as shown in Fig. \ref{fig2}."
" The major and the minor merger rates are fairly constant for ealaxies with Al,<107AL... at 1 and decrease with time for z<I."," The major and the minor merger rates are fairly constant for galaxies with $M_{\star} < 10^{11}$, at $z > 1$ and decrease with time for $z < 1$."
" The merger rates of galaxies with AZ,>1m sstronegly decrease with redshift.", The merger rates of galaxies with $M_{\star} > 10^{11}$ strongly decrease with redshift.
 This partially rellects the increase in the number density of galaxies with time., This partially reflects the increase in the number density of galaxies with time.
 In this Subsection. we examine in detail how the observed merger fractions compare with those predicted. by. the Millennium simulation.," In this Subsection, we examine in detail how the observed merger fractions compare with those predicted by the Millennium simulation."
 The merger fraction. as defined in Subsection 2.1.. is the fraction. of galaxies undergoing a merger within a given stellar mass and redshift range.," The merger fraction, as defined in Subsection \ref{sec21}, is the fraction of galaxies undergoing a merger within a given stellar mass and redshift range."
 Fie., Fig.
" 5 shows a comparison between the observed. and simulated merger fractions for galaxies with AZ,>10:3 ((upper panels). AM,107 (middle panels) and 10"" iAL.<107 (lower panels) as a function of redshift."," \ref{figcomp} shows a comparison between the observed and simulated merger fractions for galaxies with $M_{\star} > 10^{11}$ (upper panels), $M_{\star} > 10^{10}$ (middle panels) and $10^{9}$ $<M_{\star}< 10^{10}$ (lower panels) as a function of redshift."
 To demonstrate the effect. of the time-scale for merging. the," To demonstrate the effect of the time-scale for merging, the"
Without implementing these improvements. it 15 difficult to say whether the cores of the modeled helmet streamers would agree better with observation. or if there is a flaw. or a missing element. in the structural assumptions used to create the model.,"Without implementing these improvements, it is difficult to say whether the cores of the modeled helmet streamers would agree better with observation, or if there is a flaw, or a missing element, in the structural assumptions used to create the model."
 Creating a model of the corona based on high-density sheets. with extended longitudes above the observed position of filaments. has given us a good overall agreement of the outline of large helmet streamers. in other words. the legs of helmet streamer are well simulated.," Creating a model of the corona based on high-density sheets, with extended longitudes above the observed position of filaments, has given us a good overall agreement of the outline of large helmet streamers, in other words, the legs of helmet streamer are well simulated."
 This is in support of the multiple current sheet model of helmet streamers (Eddy1973:Crookeretal. 1993).," This is in support of the multiple current sheet model of helmet streamers \citep{eddy1973,crooker1993}."
. Indeed. without resorting to an empirical model. the direct observation of enhanced brightness in the streamer legs. correlated with the position of prominences/active regions in the low corona is supportive of this view.," Indeed, without resorting to an empirical model, the direct observation of enhanced brightness in the streamer legs, correlated with the position of prominences/active regions in the low corona is supportive of this view."
 Figure 7 illustrates schematically two possible views of the magnetic topology within the large helmet streamer., Figure \ref{schematic} illustrates schematically two possible views of the magnetic topology within the large helmet streamer.
 The left schematic is based on a large system of closed loops which form the streamer core., The left schematic is based on a large system of closed loops which form the streamer core.
 In the right schematic. systems of closed loops are smaller. and localized directly above prominences or active regions.," In the right schematic, systems of closed loops are smaller, and localized directly above prominences or active regions."
 The helmet shape of the streamer is caused by the merging of open field over a large region. and the overall topology of the streamer is one of multiple current sheets.," The helmet shape of the streamer is caused by the merging of open field over a large region, and the overall topology of the streamer is one of multiple current sheets."
 The structure of our model is based on the latter view. and it has considerable success in replicating both the overall shape of streamers. and the sheets of enhanced brightness (the streamer legs) directly above the position of prominences or active regions.," The structure of our model is based on the latter view, and it has considerable success in replicating both the overall shape of streamers, and the sheets of enhanced brightness (the streamer legs) directly above the position of prominences or active regions."
 On the other hand. the PMO synoptic map of figure 3. shows one filament situated just North of the center of the streamer base.," On the other hand, the PMO synoptic map of figure \ref{sourcestrmr} shows one filament situated just North of the center of the streamer base."
 This prominence can be seen in the He I CCHIP image of figure 6.. underlying a small. dark. arched cavity extending above its position to a height of iin the EIT aand the lowest height of the MKIV coronameter’s field of view.," This prominence can be seen in the He I CHIP image of figure \ref{latprof}, underlying a small, dark, arched cavity extending above its position to a height of in the EIT and the lowest height of the MKIV coronameter's field of view."
 This ts very different from the enhanced brightness above the neighboring prominence. under the northernmost streamer leg.," This is very different from the enhanced brightness above the neighboring prominence, under the northernmost streamer leg."
 Is it possible then that some prominences have high-density sheets extending into the corona to form streamers. or the legs of large helmet streamers. while other prominences are surrounded by a larger system of closed loops?," Is it possible then that some prominences have high-density sheets extending into the corona to form streamers, or the legs of large helmet streamers, while other prominences are surrounded by a larger system of closed loops?"
 It is known that prominences can vary considerably in structure (Heinzel&Anzer2006.andreferences within).., It is known that prominences can vary considerably in structure \citep[and references within]{heinzel2006}.
 If the closed field or arcade associated with a prominence is restricted to low, If the closed field or arcade associated with a prominence is restricted to low
1998).. along 10 lines of sight in the disk of the Milkv Way.,", along 10 lines of sight in the disk of the Milky Way."
 The data reveal an edge to the disk al R=13.9£z0.6 kpe., The data reveal an edge to the disk at $R=13.9\pm0.6$ kpc.
 Why there should be a sharp edge to the stus. while the eas profile is much more extended and does not show such an edge (Ixalberla&Ixerp2009).. remains unclear. but a critical test would be the measurement of the gas and star density in situ (IXennicutt.1989)..," Why there should be a sharp edge to the stars, while the gas profile is much more extended and does not show such an edge \citep{2009ARA&A..47...27K}, remains unclear, but a critical test would be the measurement of the gas and star density in situ \citep{1989ApJ...344..685K}."
 There are (wo new deep survevs of the Milkv Wax disk in the near-infrared: the UIXIDSS Galactic Plane Survey (GPS). that is mapping the Northern disk (Lucasetal.2008).. and the VISTA Variables in the (VVV) Survey. that is mapping 520 square degrees in (he southern disk aud bulge of our Galaxy (Minnitietal.2010)..," There are two new deep surveys of the Milky Way disk in the near-infrared: the UKIDSS Galactic Plane Survey (GPS), that is mapping the Northern disk \citep{2008MNRAS.391..136L}, and the VISTA Variables in the (VVV) Survey, that is mapping 520 square degrees in the Southern disk and bulge of our Galaxy \citep{2010NewA...15..433M}."
 While previous observations in the Galactic plane have been limited by crowding. source dimnuess. and interstellar extinction. these new public surveys allow us to pierce through the whole disk. reaching red giants in the horizontal branch (clump giants) located at the other side of our galaxy [ον (he first lime in different δαν pass-bands.," While previous observations in the Galactic plane have been limited by crowding, source dimness, and interstellar extinction, these new public surveys allow us to pierce through the whole disk, reaching red giants in the horizontal branch (clump giants) located at the other side of our galaxy for the first time in different near-IR pass-bands."
" The VVV survey observes in the Z (0.87 microns). Y (1.02 microns). J (1.25 microns). // (1.64 microns). and A, (2.14 microns) bands. while the UIXIDSS-GDP5 is observing in the J (1.25 microns). Jf (1.64 microns). and A. (2.14 microns) bands; with the photometry uniformly calibrated in the 2421ASS system 2006 ).."," The VVV survey observes in the $Z$ (0.87 microns), $Y$ (1.02 microns), $J$ (1.25 microns), $H$ (1.64 microns), and $K_{\rm s}$ (2.14 microns) bands, while the UKIDSS-GPS is observing in the $J$ (1.25 microns), $H$ (1.64 microns), and $K$ (2.14 microns) bands; with the photometry uniformly calibrated in the 2MASS system \citep{2006AJ....131.1163S}."
 Chump giants are excellent tools to study (he structure of our galaxy. because of two main reasons. discussed below: (1) they are good distance indicators. aud (1) thev are good tracers of the mass in the disks of galaxies.," Clump giants are excellent tools to study the structure of our galaxy because of two main reasons, discussed below: (i) they are good distance indicators, and (ii) they are good tracers of the mass in the disks of galaxies."
 The UIXKIDSS-GDPS and VVV Survevs give the largest homogeneous census of clump giants out to well bevond the extent of the Milky Way stellar disk., The UKIDSS-GPS and VVV Surveys give the largest homogeneous census of clump giants out to well beyond the extent of the Milky Way stellar disk.
" The limiting DoPhot (Schechteretal.1993) magnitude of single epoch VVV images processed by the Cambridge Astronomical Survey Unit (CASU) pipeline v1.0 is A,=17.5 in the disk fields. and the DoPhot photometry on the combined epochs allow to reach approximately J=20 mag. and A,=18.5 mag (Fie. 1))."," The limiting DoPhot \citep{1993PASP..105.1342S} magnitude of single epoch VVV images processed by the Cambridge Astronomical Survey Unit (CASU) pipeline v1.0 is $K_{\rm s}=17.5$ in the disk fields, and the DoPhot photometry on the combined epochs allow to reach approximately $J=20$ mag, and $K_{\rm s}=18.5$ mag (Fig. \ref{cmd}) ),"
" slightly fainter (han the UIIDSS-GPS. and much fainter than the near-IR survevs 2MÀSS and DENIS. for which the limit is A,=14.3 mag (Skrutskieetal.2006).."," slightly fainter than the UKIDSS-GPS, and much fainter than the near-IR surveys 2MASS and DENIS, for which the limit is $K_{\rm s}=14.3$ mag \citep{2006AJ....131.1163S}."
 The distance probed along the line of sight depends on the reddening of the fields., The distance probed along the line of sight depends on the reddening of the fields.
 For example. in disk fields with low extinction (Ay<3 mag) the UKIDSS-GPS and VVV surveys would see clump giants bevond 50 kpe.," For example, in disk fields with low extinction $A_{\rm
 V}<3$ mag) the UKIDSS-GPS and VVV surveys would see clump giants beyond 50 kpc."
 Therelore we can search for the edge of the disk of our galaxy. using well calibrated standard candles., Therefore we can search for the edge of the disk of our galaxy using well calibrated standard candles.
 The chunp giants with known parallaxes measured) by the Iipparcos satellite are the best collectively calibrated standarcl candles (Paczvnski&Stanek 1993).., The clump giants with known parallaxes measured by the Hipparcos satellite are the best collectively calibrated standard candles \citep{1998ApJ...494L.219P}. .
Ομ(10σαι)=10 kms +.,$v_{\rm wind}(10^{16} {\rm cm})=10$ km $^{-1}$.
 In all cases we set the initial teniperature of the CSM/ISM and the ejecta to 7=101 Kk. For the magnetic field. we look at the case of a constant magnetic field. strength.with values of 3. 10. and 20 μέ. We restrict ourselves for now to parallel shocks. where the magnetic field is aligned with the shock normal. for reasons mentioned in 4..," In all cases we set the initial temperature of the CSM/ISM and the ejecta to $T=10^4$ K. For the magnetic field we look at the case of a constant magnetic field strength,with values of $3$, $10$, and $20$ $\umu$ G. We restrict ourselves for now to parallel shocks, where the magnetic field is aligned with the shock normal, for reasons mentioned in \ref{sec:method}."
 In this case the magnetic field. is the same on both sides of the shock and 1e cilfusion coellicient depends only on the particle energy. in the case of Bohm diffusion., In this case the magnetic field is the same on both sides of the shock and the diffusion coefficient depends only on the particle energy in the case of Bohm diffusion.
 While the code has an MIID. solver ancl we can follow the magnetic Ποιά dynamically. this only becomes interesting in two dimensions. something we celer until a later work.," While the code has an MHD solver and we can follow the magnetic field dynamically, this only becomes interesting in two dimensions, something we defer until a later work."
 We solve the Euler equations ancl parametrise the magnetic Field strength for use in the SDEs., We solve the Euler equations and parametrise the magnetic field strength for use in the SDEs.
 We find (sec 6.5)) that particles diffuse far enough downstream to reach the reverse shock such that they are re-accelerated. provided that the magnetic field in the ejecta is strong enough to confine the particles.," We find (see \ref{sec:rev}) ) that particles diffuse far enough downstream to reach the reverse shock such that they are re-accelerated, provided that the magnetic field in the ejecta is strong enough to confine the particles."
 We therefore investigate two cillerent. cases for the magnetic field in the ejecta: we either. parametrise the magnetic field. strength such that it decays with radius as δὲ7 as to represent the expanding fossil Ποια of the progenitor. frozen. into the plasma while conserving magnetic Dux. or we keep the magnetic field fixed at the same value as in the ISM/CSM.," We therefore investigate two different cases for the magnetic field in the ejecta: we either parametrise the magnetic field strength such that it decays with radius as $R^{-2}$ as to represent the expanding fossil field of the progenitor, frozen into the plasma while conserving magnetic flux, or we keep the magnetic field fixed at the same value as in the ISM/CSM."
 In the first case. the magnetic field strength for reasonable progenitor fields is very weak. making the gvroradius (and thus the diffusion length-scale) of the particles very large when Bohm cillusion is assumed.," In the first case, the magnetic field strength for reasonable progenitor fields is very weak, making the gyroradius (and thus the diffusion length-scale) of the particles very large when Bohm diffusion is assumed."
 In the latter case. the diffusion in the ejecta proceeds at a similar rate as in the ISAL/CSAL. so particles that encounter the reverse shock can be re-accelerated. as we will illustrate in Sec. 6.5..," In the latter case, the diffusion in the ejecta proceeds at a similar rate as in the ISM/CSM, so particles that encounter the reverse shock can be re-accelerated, as we will illustrate in Sec. \ref{sec:rev}."
 At the forward. shock of SNRs. and. possibly. also at the reverse shock (??).. particles are believed to be accelerated to relativistic energies. up to 101 eV [for protons.," At the forward shock of SNRs, and possibly also at the reverse shock \citep{2008HelderVink, 2010ZirakashviliAharonian}, , particles are believed to be accelerated to relativistic energies, up to $10^{15}$ eV for protons."
 The Favoured. mechanism is cillusive shock acceleration (DSA). where every time the particle crosses the shock interlace it experiences a Lorentz boost. after which the momentum direction is randomised in the local frame.," The favoured mechanism is diffusive shock acceleration (DSA), where every time the particle crosses the shock interface it experiences a Lorentz boost, after which the momentum direction is randomised in the local frame."
 Since the plasma rest frames on either side of the shock dilfer in speed. this results in à svstematic increase of the particle momentum. every time the particle eveles from upstream to downstream and back (?).," Since the plasma rest frames on either side of the shock differ in speed, this results in a systematic increase of the particle momentum every time the particle cycles from upstream to downstream and back \citep{1978Bella}."
. 2222272 derive the spectrum. of. shock-accelerated particles. based on the energy gain per evcle versus the probability of being adyvected: away and. lost. from. the acceleration. process.," \citet{1977Krymskii, 1977Axfordetal, 1978Bella,1978Bellb, 1978BlandfordOstriker, 1983Drury} derive the spectrum of shock-accelerated particles, based on the energy gain per cycle versus the probability of being advected away and lost from the acceleration process."
" The escape probability is given by Dus=ANfre. where V, is the shock velocity. r the shock compression ratio and e is the velocity of the cosmic rays. equal to the speed. of light."," The escape probability is given by $P_{\rm esc}=4 V_{\rm s}/rc$, where $V_{\rm s}$ is the shock velocity, $r$ the shock compression ratio and $c$ is the velocity of the cosmic rays, equal to the speed of light."
 The mean time Af for a relativistic particle to complete a evcle probing the upstream. and the downstream plasma and the corresponding boost in momentum p is given by (77): where α is the cdilfusion cocllicient ancl we will use the subscript 1 for upstream. and 2 for downstream properties.," The mean time $\Delta t$ for a relativistic particle to complete a cycle probing the upstream and the downstream plasma and the corresponding boost in momentum $p$ is given by \citep{1983LagageCesarsky,1983Drury}: where $\kappa$ is the diffusion coefficient and we will use the subscript $1$ for upstream-, and $2$ for downstream properties."
 The test particle limit results in a power-law spectrum in momentum. with the number density of particles per unit momentum following the distribution with the slope.," The test particle limit results in a power-law spectrum in momentum, with the number density of particles per unit momentum following the distribution with the slope."
 Here £44=1./we ds the return probability in à shock crossing cvcle., Here $P_{\rm ret}=1-P_{\rm esc}$ is the return probability in a shock crossing cycle.
 The acceleration rate of relativistic protons or electrons of given encrey fo=pemyersoL GeV is: giving a tvpical time-scale for acceleration time of lt will be useful to consider the case of Bohm diffusion for comparison with the simulations. and to look at. the combined ellect of acceleration. and. synchrotron losses. which is important for electrons.," The acceleration rate of relativistic protons or electrons of given energy $E = pc \gg m_{\rm p}c^2\sim 1$ GeV is: giving a typical time-scale for acceleration time of It will be useful to consider the case of Bohm diffusion for comparison with the simulations, and to look at the combined effect of acceleration and synchrotron losses, which is important for electrons."
 The change in energy at the shock is then given by: [or a parallel shock. where &j=8»DpcE/35ZeD is the Bohm ciffusion coellicient. which depends on the evroradius of the particles. D the magnetic field strength. Z the charge number. e the elementary. electric charge and. m the mass of the partic," The change in energy at the shock is then given by: for a parallel shock, where $\kappa_1=\kappa_2=D_{\rm B} = c E/3ZeB$ is the Bohm diffusion coefficient, which depends on the gyroradius of the particles, $B$ the magnetic field strength, $Z$ the charge number, $e$ the elementary electric charge and $m$ the mass of the particle."
 We have introduced. ai term that accounts for svnchrotron losses with The mass dependence makes this term negligible for xotons., We have introduced a term that accounts for synchrotron losses with The mass dependence makes this term negligible for protons.
 Llieh-enerey electrons on the other hand lose a substantial fraction of their energy by svuchrotron racliation., High-energy electrons on the other hand lose a substantial fraction of their energy by synchrotron radiation.
 With a Thonison cross section of er=6.65.n102 cnm? or electrons this gives for A;=1.29.10 cm s , With a Thomson cross section of $\sigma_T=6.65\times10^{-25}$ $^2$ for electrons this gives for $\lambda_s=1.29\times10^{-9}$ cm s $^{-1}$ .
The same loss-term can be used to account for inverse Compton losses. where the energy that is lost in upscattering ohotons from the microwave background: corresponds. to svnchrotron losses in an equivalent magnetic field Boar= Ter (72)..," The same loss-term can be used to account for inverse Compton losses, where the energy that is lost in upscattering photons from the microwave background corresponds to synchrotron losses in an equivalent magnetic field $B_{\rm CMB}=3.27\ \umu$ G \citep{1998Reynolds}. ."
 When the actual magnetic field.is stronger han this value. svnchrotron losses will dominate over inverse Compton losses unless there is another source of photons hat boosts the inverse Compton scattering rate.," When the actual magnetic fieldis stronger than this value, synchrotron losses will dominate over inverse Compton losses unless there is another source of photons that boosts the inverse Compton scattering rate."
 Fornow we, Fornow we
"(Preston (1967))) y= Dio) = B, BL)(e)) js ad à maximum. and uis the limb-darkening coefficient.","\cite{1967ApJ...150..547P}) ) _l= ) = B_p ) is at a maximum, and $u$ is the limb-darkening coefficient."
 For ο and D stars. 4=0.35 (1997))) can be used.," For O and B stars, $u=0.35$ \cite{1997A&A...317..723T}) ) can be used."
 The phase-averaged ratios D;j/D. where B;=/Bilojdo. for all of the possible 7 aud i do not exceed 0.3.," The phase-averaged ratios $\ov{B_l}/B_p$, where $
\ds \ov{B}_l = \!\int_{0}^{2\pi}\!\! {B_l}(\phi)\, d\phi \, ,
$ for all of the possible $\beta$ and $i$ do not exceed 0.3."
 D; depends significantly on the inclination 7. which is determined bv ihe random orientation of the stellar rotation axis. and on the angle 3. which also varies over a wide range.," $\ov{B}_l$ depends significantly on the inclination $i$, which is determined by the random orientation of the stellar rotation axis, and on the angle $\beta$, which also varies over a wide range."
 Thus. the mean field strength 2; cannot be used for statistical studies.," Thus, the mean field strength $\ov{B}_l$ cannot be used for statistical studies."
 In addition. (he times at which the field measurements were made are sometimes distributed verv irregularly in stellar rotation phase.," In addition, the times at which the field measurements were made are sometimes distributed very irregularly in stellar rotation phase."
 As one of the possible characteristics of the mean magnetic field that depends weakly on the random values of / and 32. we will consider the rool-meansquare (mis) field widely used in the literature (see. e.g.. Dohlenderetal. (1993))).," As one of the possible characteristics of the mean magnetic field that depends weakly on the random values of $i$ and $\beta$ , we will consider the root-meansquare (rms) field widely used in the literature (see, e.g., \cite{1993A&A...269..355B}) ), =."
 llere. (he summation is over all » field measurements.," Here, the summation is over all $n$ field measurements."
" In addition to this quantity. we will introduce another field characteristic. =(3) where Z4, and Biya, ave (he minimum (given the sign) and maximum values of 23; determined from allfield measurements."," In addition to this quantity, we will introduce another field characteristic, =, where $B_{\mbox{\sz min}}$ and $B_{\mbox{\sz max}}$ are the minimum (given the sign) and maximum values of ${B_l}$ determined from allfield measurements."
"with Using the generalized. Laplace coefficients defined as (Warel 1989: IXorvcansky Pollack 93. hereafter 92]: where: we can rewrite V, uncer the form: ⋜⋯∠⇂↿↓↕∢⊾⋖⋅⇀∖↓≻↓⋅∢⋅⊳∖⊳∖⊲↓∪⊔∪⊔↿⇂↥⋖⊾↓⋰↓⋏∙≟↓∐↓⋯⊔∠⇂≻↕∠⇂∢⋅⇂↥⋜↧≻↿∪∣⋈⋅∠⇂⊲↓∖⋰⊓⇂⋖⋅∠⇂∣⋡∙∖⇁⋜↧⇂∎⋜","with and Using the generalized Laplace coefficients defined as (Ward 1989; Korycansky Pollack 93, hereafter KP93): where: we can rewrite $\Psi'_m$ under the form: and the expression on the right–hand–side has to be divided by a factor two if $m=0$."
↧≼∙↥∢≱↓⋅↿∖∖⊽∢≱↕⇂⋅↗∣∣∶∪⊳↾↓∖↓↕⋖⋅⊳∖⊔∣⋡≱∖≼∙↓⋰↓↓≻↿↼↓≻↼⊲↓⊔∠⇂⊲⊔∼⋜⋯⋅⊳∖↿↓⋯∣⇂↓↕⋖⊾ ⊏↥⇂⇂⋜↧↓↥↿↕↿, The subscript 'p' indicates that the quantities have to be evaluated at $r=r_p$.
↕∢⊾≻↓⋯∖⇁⋖⋅∣∪∣⋡∢⋅⋖⊾∖⇁⋜↧↓⋯⋯⊾∠⇂⋜∐∣⋮∶∣⋮↗∣⋡ In this paper we consider small perturbations. so that the basic equations can be linearized.," In this paper we consider small perturbations, so that the basic equations can be linearized."
 We denote Eulerian perturbations with a prime., We denote Eulerian perturbations with a prime.
 Since the equilibrium is axisvmmetric and steady. we can expand cach of the perturbed quantities in. Fourier series with respect to the variable ó ancl solve separately for cach value of m.," Since the equilibrium is axisymmetric and steady, we can expand each of the perturbed quantities in Fourier series with respect to the variable $\phi$ and solve separately for each value of $m$."
" The general problem may then be reduced to calculating the response of the disc to the real part of a complex. potential of the form Wi,exp(mswh]. where the amplitude V,nn is real and the frequeney w=mO,. and summing up over 2."," The general problem may then be reduced to calculating the response of the disc to the real part of a complex potential of the form $\Psi'_m(r) {\rm exp} \left[ i \left( m\varphi - \omega t \right)
\right]$, where the amplitude $\Psi'_m$ is real and the frequency $\omega \equiv m \Omega_p$, and summing up over $m$."
 We make all Eulerian fluid state variable perturbations complex by writing: where .X is any state variable., We make all Eulerian fluid state variable perturbations complex by writing: where $X$ is any state variable.
 The physical perturbations will be recovered by taking the real part of these complex quantities., The physical perturbations will be recovered by taking the real part of these complex quantities.
" We denote € the Lagrangian displacement. and. write its m th Fourier component as £,,(r)expηwf]."," We denote $\xib$ the Lagrangian displacement, and write its $m$ –th Fourier component as $\xib_m (r) {\rm exp} \left[ i
\left( m\varphi - \omega t \right) \right]$."
 Since here we ave only interested in the perturbations induced by the planet in the plane of the disc. we take £«=0.," Since here we are only interested in the perturbations induced by the planet in the plane of the disc, we take $\xi_z=0$."
 Lincarization of the induction equation (3)). or the equivalent. statement of the conservation of magnetic Dux. leads to an expression for the perturbed magnetic field in the form (Chandrasekhar Fermi 1953): The corresponding perturbation of the Lorentz force per unit volume is: Use of equations (17)) and (18)) with €=0 gives: where we have defined the dimensionless quantities:," Linearization of the induction equation \ref{induction}) ), or the equivalent statement of the conservation of magnetic flux, leads to an expression for the perturbed magnetic field in the form (Chandrasekhar Fermi 1953): The corresponding perturbation of the Lorentz force per unit volume is: Use of equations \ref{Bp}) ) and \ref{FMp}) ) with $\xi_z=0$ gives: where we have defined the dimensionless quantities:"
These simulations show a very interesting result in terms of the formation of the white dwarf discs.,These simulations show a very interesting result in terms of the formation of the white dwarf discs.
" There is a lack of bodies scattered onto star-grazing orbits by a single planet, at 30AU."," There is a lack of bodies scattered onto star-grazing orbits by a single planet, at 30AU."
" In fact in these simulations no test particle has a semi-major axis less than 0.5a,;.", In fact in these simulations no test particle has a semi-major axis less than $a_{pl}$.
" Even including the eccentricity, no test particles has a pericentre of less than 0.3a,;."," Even including the eccentricity, no test particles has a pericentre of less than $a_{pl}$."
" In order for an asteroid to be tidally disrupted and form the observed discs, it must be scattered onto an orbit with pericentre less than the tidal radius (on the order of Ro)."," In order for an asteroid to be tidally disrupted and form the observed discs, it must be scattered onto an orbit with pericentre less than the tidal radius (on the order of $R_{\odot}$ )."
" It may be that such a small percentage of test particles are scattered further in towards the star, but that our simulations do not find them because we have not included a sufficiently large number of test particles."," It may be that such a small percentage of test particles are scattered further in towards the star, but that our simulations do not find them because we have not included a sufficiently large number of test particles."
" Given this caveat, these simulations show that a single planet is incapable of producing the observed discs and that an inner planetary system is necessary."," Given this caveat, these simulations show that a single planet is incapable of producing the observed discs and that an inner planetary system is necessary."
 An analytic (or semi-analytic) model is useful to understand the physical mechanisms causing instability and allows us to get scaling laws that describe the behaviour over a wide range of parameter space., An analytic (or semi-analytic) model is useful to understand the physical mechanisms causing instability and allows us to get scaling laws that describe the behaviour over a wide range of parameter space.
 Analytically it is expected that test particles orbits close to the planet will become chaotic due to the overlap of mean motion resonances., Analytically it is expected that test particles orbits close to the planet will become chaotic due to the overlap of mean motion resonances.
 This condition defines the chaotic zone within which this occurs as (?):: where C—1.3., This condition defines the chaotic zone within which this occurs as \citep{Wisdom1980}: where $C=1.3$.
 Test particles on chaotic orbits will either be ejected or scattered into the inner planetary system., Test particles on chaotic orbits will either be ejected or scattered into the inner planetary system.
" Thus, the fraction of the disc mass that will be removed, Manatytic/Mbeit can be calculated, for a given surface density profile, assuming that all test particles with initial semimajor axes less than Achaotic=Apt+ÓQchaos are removed, for the index in the surface density profile, α—1 and K=Meets/tapi(22/5—1), for a disc with outer edge, aas=(23 Pay."," Thus, the fraction of the disc mass that will be removed, $M_{analytic}/M_{belt}$ can be calculated, for a given surface density profile, assuming that all test particles with initial semimajor axes less than $a_{chaotic} = a_{pl}+ \delta a_{chaos}$ are removed, for the index in the surface density profile, $\alpha =1$ and $K=M_{belt}/\pi a_{pl}(2^{2/3}-1)$, for a disc with outer edge, $a_{max}=(\frac{2}{1})^{2/3}a_{pl}$ ."
" This can be compared to our N-body simulations, where we show that most test particles with semi-major axes less than Gechaotic are removed, but that some test particles with a>GQchaotic are also removed."," This can be compared to our N-body simulations, where we show that most test particles with semi-major axes less than $a_{chaotic}$ are removed, but that some test particles with $a>a_{chaotic}$ are also removed."
 Fig., Fig.
" 3 shows initial semi-major axis and eccentricities of all test particles in the baseline simulation, with those test particles that are scattered by a 1Mwnep planet highlighted."," \ref{fig:initial} shows initial semi-major axis and eccentricities of all test particles in the baseline simulation, with those test particles that are scattered by a $1M_{Nep}$ planet highlighted."
 The higher the initial eccentricity of the test particles the higher the number of test particles outside of the chaotic zone that are removed., The higher the initial eccentricity of the test particles the higher the number of test particles outside of the chaotic zone that are removed.
 This also applies to inclination., This also applies to inclination.
" No test particles were ejected in the 10” yrs of this simulation as this timescale is too long for a Neptune mass planet to increase a test particles semi-major axis to greater than Agalactic=15, OOOAU."," No test particles were ejected in the $10^7$ yrs of this simulation as this timescale is too long for a Neptune mass planet to increase a test particles semi-major axis to greater than $a_{galactic}=15,000$ AU."
 The formulation for the chaotic zone (Eq. 3)), The formulation for the chaotic zone (Eq. \ref{eq:chaos}) )
 was developed for bodies on circular orbits., was developed for bodies on circular orbits.
" Although ? showed that the same formalism applies for eccentric planets, when all test particles have the forced (or the planet’s) eccentricity, the behaviour for test particles with high free eccentricities (and inclinations) is different."," Although \cite{Quillen06} showed that the same formalism applies for eccentric planets, when all test particles have the forced (or the planet's) eccentricity, the behaviour for test particles with high free eccentricities (and inclinations) is different."
" There are a few sets of simulations that show that the chaotic zone is larger for eccentric or inclined bodies (e.g., ?)), but there is no analytic prescription."," There are a few sets of simulations that show that the chaotic zone is larger for eccentric or inclined bodies (e.g., \cite{Veras04}) ), but there is no analytic prescription."
 Although the formalism for the chaotic zone (Eq. 3)), Although the formalism for the chaotic zone (Eq. \ref{eq:chaos}) )
" was developed specifically in terms of semi-major axis, Fig."," was developed specifically in terms of semi-major axis, Fig."
 3 shows that most of the structure in eccentricity can be described by the scattering of test particles with pericentres closer to the planet than the chaotic zone i.e. q<apid-Óacnaos., \ref{fig:initial} shows that most of the structure in eccentricity can be described by the scattering of test particles with pericentres closer to the planet than the chaotic zone i.e. $q< a_{pl}+\delta a_{chaos}$.
" This may be because such particles have strong close encounters with the planet, or because of the increased resonance width of eccentric orbits."," This may be because such particles have strong close encounters with the planet, or because of the increased resonance width of eccentric orbits."
" However, it also seems that the full picture is somewhat more complicated, especially for highly eccentric orbits."," However, it also seems that the full picture is somewhat more complicated, especially for highly eccentric orbits."
When the uncertainties in S. M. and H can be approximated as v/S. V/M. and \/H. respectively. then the 2 conditions above imply the same restrictions on the relative values of S and M and S and H as do 3 and 4.,"When the uncertainties in $S,$ $M,$ and $H$ can be approximated as $\sqrt{S},$ $\sqrt{M},$ and $\sqrt{H},$ respectively, then the $2$ conditions above imply the same restrictions on the relative values of $S$ and $M$ and $S$ and $H$ as do 3 and 4."
" We denote systems that satisfy the 30 conditions ""SSS-36."""," We denote systems that satisfy the $3\sigma$ conditions $3\sigma$ ."""
 The HR and 3e conditions select sources based on the dominance of flux in the S band., The HR and $3\sigma$ conditions select sources based on the dominance of flux in the S band.
 If. however. there is a large column of gas between the source and the detector. the preferential absorption of soft photons erodes the signal in S relative to that in the other bands. so that even genuine SSSs may satisfy neither condition.," If, however, there is a large column of gas between the source and the detector, the preferential absorption of soft photons erodes the signal in $S$ relative to that in the other bands, so that even genuine SSSs may satisfy neither condition."
 Especially if the flux is small. the HR and 36 conditions will fail to identify many SSSs.," Especially if the flux is small, the HR and $3\sigma$ conditions will fail to identify many SSSs."
 We therefore introduce weaker conditions that can identify additional SSSs. but which may also identify some harder sources.," We therefore introduce weaker conditions that can identify additional SSSs, but which may also identify some harder sources."
" Because the SSS status of sources identified via weaker conditions is less secure. we refer to all such sources as ""quasisoft"" sources (QSSs)."," Because the SSS status of sources identified via weaker conditions is less secure, we refer to all such sources as “quasisoft"" sources (QSSs)."
 Any QSSs with high enough count rates to allow spectral fits. and for which the fits shows a low effective temperature. can then be reclassified as SSSs.," Any QSSs with high enough count rates to allow spectral fits, and for which the fits shows a low effective temperature, can then be reclassified as SSSs."
 To proceed. we turn to another hallmark of SSSs: the lack of high-energy emission from many of the “classical” SSSs.," To proceed, we turn to another hallmark of SSSs: the lack of high-energy emission from many of the “classical"" SSSs."
 We use the condition. to identify 2 branches in the identification procedure.," We use the condition, to identify $2$ branches in the identification procedure."
 If the condition is satisfied. we know that the source cannot be hard. and devise further tests to determine if it 15 soft.," If the condition is satisfied, we know that the source cannot be hard, and devise further tests to determine if it is soft."
 We refer to the path defined by these further tests as7., We refer to the path defined by these further tests as.
 If the condition given in (10) is not satisfied. it is still possible for the source to be an SSS.," If the condition given in (10) is not satisfied, it is still possible for the source to be an SSS."
 implements those conditions that can select SSSs while rejecting hard sources., implements those conditions that can select SSSs while rejecting hard sources.
 If C(60/AC(6)<=1 and C(5)/AC(S)<2. then there is at most a weak detection in the bins corresponding to medium energies.," If $C(6)/\Delta\, C(6) \leq 1$ and $C(5)/\Delta\, C(5) \leq 2,$ then there is at most a weak detection in the bins corresponding to medium energies."
" Sources satisfying these conditions are called ""QSS-NOH'"": they exhibit little or no emission above 1.1 keV. If the NOH condition is not satisfied. then there is a detection above 1.1 keV. The condition M/A>3. creates another bifurcation which we can follow along 2 sub-branches. and1."," Sources satisfying these conditions are called “QSS-NOH”; they exhibit little or no emission above $1.1$ keV. If the NOH condition is not satisfied, then there is a detection above $1.1$ keV. The condition $M/\Delta\, M > 3,$ creates another bifurcation which we can follow along $2$ sub-branches, and."
2. If we have reached this point along the path. there is no hard emission. but there is significant emission in the M band.," If we have reached this point along the path, there is no hard emission, but there is significant emission in the $M$ band."
 We can call the source “supersoft” only if emission in the S band dominates and if. there is a steep decline in the number of counts from the low to high energy bins.," We can call the source “supersoft"" only if emission in the $S$ band dominates and if, there is a steep decline in the number of counts from the low to high energy bins."
" We therefore require that EV22EVI and C(5)>2.6C(6). Sources satisfying these conditions are designated “QSS-MNOH""."," We therefore require that ${{S}\over{\Delta\, S}} > 2\, {{M}\over{\Delta\, M}}$ and $C(5) > 2.6\, C(6).$ Sources satisfying these conditions are designated “QSS-MNOH""."
 Sources that do not satisfy the conditions for QSS-MNOH could be SSSs which are highly absorbed. or they could be sources with spectra that can best be viewed as flat in M relative to S —L.e.. little emission in the soft or hard bands. e.g.. a 200 eV blackbody.," Sources that do not satisfy the conditions for QSS-MNOH could be SSSs which are highly absorbed, or they could be sources with spectra that can best be viewed as flat in M relative to S –i.e., little emission in the soft or hard bands, e.g., a $200$ eV blackbody."
" We call such sources ""QSS-ENOH'""."," We call such sources “QSS-FNOH""."
 If we have reached this point along the path. there is no hard emission. but there is also little flux in the M band.," If we have reached this point along the path, there is no hard emission, but there is also little flux in the $M$ band."
 The source can be SSS only if there is emission in the S band., The source can be SSS only if there is emission in the $S$ band.
" If S/AS>3. we use the designation ""QSS-SNOH""."," If $S/\Delta\, S > 3,$ we use the designation “QSS-SNOH""."
 If the source fails the SNOH test. it is placed in the QSS-FNOH category.," If the source fails the SNOH test, it is placed in the QSS-FNOH category."
 In this case. H/AH>0.5. so there may be some hard emission.," In this case, $H/\Delta\, H > 0.5,$ so there may be some hard emission."
 If the hard emission is nevertheless a small component of the spectrum. the source could still be an SSS.," If the hard emission is nevertheless a small component of the spectrum, the source could still be an SSS."
 First we require that either (a) H/T<0.005. where T=S+M+H. or (b) both S/AS>3 and H/AH«1.," First we require that either (a) $H/T < 0.005,$ where $T= S + M + H,$ or (b) both $S/\Delta\, S > 3$ and $H/\Delta\, H < 1$."
" In addition. we require the same conditions on HR2 required for the ""HR conditions"": HR2«—0.8 and HR2«—0.8. At this point. however. we relax the conditions on HRI."," In addition, we require the same conditions on HR2 required for the “HR conditions"": $HR2 < -0.8$ and $HR2_{\Delta} < -0.8.$ At this point, however, we relax the conditions on HR1."
 That is. we allow for the possibility that absorption has eroded the flux in the S band relative to the flux in the M band.," That is, we allow for the possibility that absorption has eroded the flux in the S band relative to the flux in the M band."
" The new conditions on HRI are: HRI«—0.5 and HRI,«0. These conditions could. e.g.. be satisfied by an absorbed 100 eV SSS."," The new conditions on HR1 are: $HR1 < -0.5$ and $HR1_{\Delta} < 0.$ These conditions could, e.g., be satisfied by an absorbed $100$ eV SSS."
" Sources satisfying these conditions will be referred to as ""QSS-HR, .", Sources satisfying these conditions will be referred to as $_1$ ”.
" The 36, conditions represent a somewhat relaxed version of the 36 conditions."," The $3\, \sigma_1$ conditions represent a somewhat relaxed version of the $3\, \sigma$ conditions."
" We require a two-c detection in S. S/AS2. and that the hard flux be a small fraction of the total flux: H/T<0.005. As before. we require S/AS>3H/AH. But we relax the condition (8). replacing it with $/AS>M/AM. Systems satisfying this set of conditions are designated ""QSS-36,""."," We require a $\sigma$ detection in $S$, $S/\Delta\, S > 2,$ and that the hard flux be a small fraction of the total flux: $H/T < 0.005.$ As before, we require $S/\Delta\, S > 3\, H/\Delta\, H.$ But we relax the condition (8), replacing it with $S/\Delta\, S > M/\Delta\, M.$ Systems satisfying this set of conditions are designated $3 \sigma_1$ ""."
" Sources with 4/7<0.05. and A > 3are designated ""QSS-a”."," Sources with $H/T < 0.05,$ and $S\over \Delta\, S$ $> 3$ are designated $\sigma$ ”."
 The remaining sources are. for the purposes of this classification scheme. not SSSs.," The remaining sources are, for the purposes of this classification scheme, not SSSs."
One of the most incresting problems in eravitation fweory is the study of the relation tha exists between je critical phenonena aud the process of black hole formation.,One of the most interesting problems in gravitation theory is the study of the relation that exists between the critical phenomena and the process of black hole formation.
 The studies of non-lnemitv of the Einstein field equiλος near the hreshold of blacx hole formation reveal very rich phenomena [1]-- EMi|] which are quie sinular to critical phenomena i1 Statistical Mechanics aud. Quautuu Fick| Theory |--[5]..," The studies of non-linearity of the Einstein field equations near the threshold of black hole formation reveal very rich phenomena \cite{Chop93a}- \cite{Chop93c}, which are quite similar to critical phenomena in Statistical Mechanics and Quantum Field Theory \cite{Golden0}- \cite{Golden1}."
 Iu particuar. bv uuucerically stidvineg the ex:witational collapse of a massless scaar fiek in3 1- splerically svinmuctric spacctimes. Chopmils found that the mass of such forned black holes Hales a scaling forπ. Mpg=Cip)tppy). where ορ) is a constant aud depends ou he miial data. iid p paranieterizes a family of iniial ¢ata in such a way that when p>p black holes :ue forned. aud when p«p no black holes are formect.," In particular, by numerically studying the gravitational collapse of a massless scalar field in $3+1$ -dimensional spherically symmetric spacetimes, Choptuik found that the mass of such formed black holes takes a scaling form, $M_{BH} = C(p)\left(p -p^{*}\right)^{\gamma}$, where $C(p)$ is a constant and depends on the initial data, and $p$ parameterizes a family of initial data in such a way that when $p > p^{*}$ black holes are formed, and when $p < p^{*}$ no black holes are formed."
 Tt was shown hat. iu contrast to C(p). the expoueif 5 ds universal o all the faimilics of initial data stiicο," It was shown that, in contrast to $C(p)$, the exponent $\gamma$ is universal to all the families of initial data studied."
 Numerically it was determined as 5—0.37., Numerically it was determined as $\gamma \sim 0.37$.
 The soluion with p=p usually called the critical sok)1. 18 found also universal.," The solution with $p = p^{*}$, usually called the critical solution, is found also universal."
 Moreover. for the massless scalar field it is periodic. too.," Moreover, for the massless scalar field it is periodic, too."
 Universality of the Critical solution aud expoucut. as well as the power-law scaliue of the slack hole mass all have given rise to he nameCollapse.," Universality of the critical solution and exponent, as well as the power-law scaling of the black hole mass all have given rise to the name."
 Cjoptuik'« studies were soon geuceralized to otv matter fields [6.7].. and now the following seems clear: (a) There are two types of eriical colapse. deuding on whether the lack hole mass takes the scaling form (Mp) or not.," Choptuik's studies were soon generalized to other matter fields \cite{Gun00,Wang01}, and now the following seems clear: (a) There are two types of critical collapse, depending on whether the black hole mass takes the scaling form $M_{BH}$ ) or not."
 When it tases the scaling foru. the corresponding ¢ollapse is called Type ZZ collapse. aud when it does not it is calced Type| E collapse.," When it takes the scaling form, the corresponding collapse is called Type $II$ collapse, and when it does not it is called Type $I$ collapse."
" In the type Z7 collaye. all the critical solutions fod so far have either discrete sclfsimilariv (DSS) or homothetic self-iniarity (IISS). depending on the matter Ποια»,"," In the type $II$ collapse, all the critical solutions found so far have either discrete self-similarity (DSS) or homothetic self-similarity (HSS), depending on the matter fields."
 Iu the type £ collaoe. the critical solutions have neither DSS nor ISS.," In the type $I$ collapse, the critical solutions have neither DSS nor HSS."
 For certain matter fields. these two types of collapse ca1 οςexist. (," For certain matter fields, these two types of collapse can co-exist. ("
b) For Type £7 collapse. the corresponding exponent is uuiversal oulv witli respect to ceraln niater fields.,"b) For Type $II$ collapse, the corresponding exponent is universal only with respect to certain matter fields."
 Usually. different mater fields have different critical soutious anc. im the secucl. different exponents.," Usually, different matter fields have different critical solutions and, in the sequel, different exponents."
 But for a 8eiven matter feld the critical solution aud the exponent are universal., But for a given matter field the critical solution and the exponent are universal.
 So far. the studiss have been mainly restricted to splerically saunetrie case and it is not clear whether «x not the critica solution aud expoucut are universal with respect to different svuinetries of the spaceines [8.9.10].. (," So far, the studies have been mainly restricted to spherically symmetric case and it is not clear whether or not the critical solution and exponent are universal with respect to different symmetries of the spacetimes \cite{Cho03a,Cho03b,Wang03}. ("
0) A critical solution for both of the two types has one and only oue unstable mode.,c) A critical solution for both of the two types has one and only one unstable mode.
 This now is considered as oue of the main criteria for a solution to be exitical. (, This now is considered as one of the main criteria for a solution to be critical. (
0) The universality oftie exponent is closely related to the last property.,d) The universality of the exponent is closely related to the last property.
 In fact. usiug dimensional analysis |1 1|l-|11| one cal show that L[] where &Á denotes the unustable mode.," In fact, using dimensional analysis \cite{Even0}- \cite{Even3} one can show that $\gamma = \frac{1}{\left|k\right|}$, where $k$ denotes the unstable mode."
sky: only located at 6«O°.,sky: only located at $\delta<0^{\circ}$.
 Therefore. we were aiming at selecting additional targets in the data base of the HES.," Therefore, we were aiming at selecting additional targets in the data base of the HES."
" White dwarfs have been selected from wide angle surveys in the southern hemisphere before. and also in the HES (see below). """," White dwarfs have been selected from wide angle surveys in the southern hemisphere before, and also in the HES (see below). “"
"Classical"" UV excess surveys. like the Montreal-Cambridge-Tololo survey (MCT:??).. or the Edinburgh-Cape survey (EC:??) can efficiently select complete samples of hot stars. including WDs.","Classical” UV excess surveys, like the Montreal-Cambridge-Tololo survey \citep[MCT;
][]{Demersetal:1986,Lamontagneetal:2000}, or the Edinburgh-Cape survey \citep[EC; ][]{Stobieetal:1997,Kilkennyetal:1997} can efficiently select complete samples of hot stars, including WDs."
 However. completeness at the end is either sacrificed for efficiency. as in the MCT (see Fig. 2)).," However, completeness at the end is either sacrificed for efficiency, as in the MCT (see Fig. \ref{UBV_DA}) ),"
 where only objects witht?B<0.6 enter the sample of stars for which follow-up spectroscopy ts obtained (?).. or efficiency is sacrificed for completeness. as in the EC.," where only objects with $U-B<-0.6$ enter the sample of stars for which follow-up spectroscopy is obtained \citep{Lamontagneetal:2000}, or efficiency is sacrificed for completeness, as in the EC."
 It has been shown that the EC is (2). , It has been shown that the EC is \citep{Stobieetal:1997}. .
However. an intermediate selection. step based on photoelectric UBV photometry has to be used to eliminate the large fraction (~30 stars.," However, an intermediate selection step based on photoelectric $UBV$ photometry has to be used to eliminate the large fraction $\sim 30$ stars."
 In the HES. WDs enter the quasar candidate sample if they have UoB«0.18 (?)..," In the HES, WDs enter the quasar candidate sample if they have $U-B<-0.18$ \citep{hespaperIII}."
 However. HES quasar candidates are inspected manually at the computer screen. and in this process hot stars. and stars clearly exhibiting stellar absorption lines (like e.g. DA white dwarfs. having strong. broad lines over a wide temperature range: see Fig.," However, HES quasar candidates are inspected manually at the computer screen, and in this process hot stars, and stars clearly exhibiting stellar absorption lines (like e.g. DA white dwarfs, having strong, broad lines over a wide temperature range; see Fig."
 | in Appendix 3)) are rejected. and follow-up spectroscopy is not obtained for them in the course of the quasar survey.," \ref{WDmodels} in Appendix \ref{Sect:SpectralAtlas}) ) are rejected, and follow-up spectroscopy is not obtained for them in the course of the quasar survey."
 This results ina very efficient quasar selection: typically which follow-up spectroscopy is obtained quasars (2).., This results in a very efficient quasar selection: typically which follow-up spectroscopy is obtained quasars \citep{hespaperIII}. .
 The remaining (Gap<L220000 KK). helium-rich WDs (DBs.DZs:see2).. a couple of interesting peculiar objects. e.g.. magnetic DBs and magnetic DAs (??) have also been discovered in this way.," The remaining $T_{\mbox{\scriptsize eff}}\lesssim 20\,000$ K), helium-rich WDs \citep[DBs,
DZs; see][]{Friedrichetal:2000}, a couple of interesting peculiar objects, e.g., magnetic DBs \citep{magDB1} and magnetic DAs \citep{Reimersetal:1994,Reimersetal:1996} have also been discovered in this way."
" The selection of white dwarf candidates described in this paper aims at an selection, that is. the contamination of the sample with other objects (e.g.. hot subdwarfs) shall be às clean às possible. in order not to waste any observing time at the VLT."," The selection of white dwarf candidates described in this paper aims at an selection; that is, the contamination of the sample with other objects (e.g., hot subdwarfs) shall be as clean as possible, in order not to waste any observing time at the VLT."
 A description of the HES plate material. plate digitisation and data reduction can be found in ?)..," A description of the HES plate material, plate digitisation and data reduction can be found in \cite{hespaperIII}."
 In this section we describe some survey properties that are particularily important for stellar work m more detail than it was done in ?).. and we briefly repeat a general description of the HES. for better readability.," In this section we describe some survey properties that are particularily important for stellar work in more detail than it was done in \cite{hespaperIII}, and we briefly repeat a general description of the HES, for better readability."
 The HES is based on IIHa-J plates taken with the nm ESO Schmidt telescope and its prism., The HES is based on IIIa-J plates taken with the m ESO Schmidt telescope and its $^{\circ}$ prism.
" It covers the magnitude range 13.02B,=17.5 (2).."," It covers the magnitude range $13.0 \gtrsim B_J
\gtrsim 17.5$ \citep{hespaperIII}."
 The magnitude limits depend on plate quality., The magnitude limits depend on plate quality.
 Note that the value given for the faint limit is the completeness limit for quasar search. which we define as the magnitude corresponding to average photographic density in the B; band >56 above the diffuse plate background. where σ is the background noise.," Note that the value given for the faint limit is the completeness limit for quasar search, which we define as the magnitude corresponding to average photographic density in the $B_J$ band $>5\sigma$ above the diffuse plate background, where $\sigma$ is the background noise."
 The detection limit of the HES tsapproximately one magnitude deeper than the completeness limit., The detection limit of the HES isapproximately one magnitude deeper than the completeness limit.
 For stellar applications. the survey magnitude range depends on the object type searched for.," For stellar applications, the survey magnitude range depends on the object type searched for."
 E.g.. in oursearch," E.g., in oursearch"
the feature αἱ 5.65 ss.! which showed the largest correlated flux on the longest baseline. Manna Kea Saint Croix. and could be used for calibrating all the antennas of the array.,"the feature at 5.65 $^{-1}$ which showed the largest correlated flux on the longest baseline, Mauna Kea – Saint Croix, and could be used for calibrating all the antennas of the array."
 The calibration of the array was performed by means of fringe fitting to the reference feature., The calibration of the array was performed by means of fringe fitting to the reference feature.
 A correction to the position of the phase center adopted at the correlator was made using (he measured absolute position of the reference leature., A correction to the position of the phase center adopted at the correlator was made using the measured absolute position of the reference feature.
 The absolute position was measured by the Iringe rate method., The absolute position was measured by the fringe rate method.
 Since the position offset of different spectral features in W75N is known to be quite large compared to the size of the map. the mapping was carried oul in two steps.," Since the position offset of different spectral features in W75N is known to be quite large compared to the size of the map, the mapping was carried out in two steps."
 First. an approximate position of all spectral features relative {ο the reference feature was determined with the fringe rate method.," First, an approximate position of all spectral features relative to the reference feature was determined with the fringe rate method."
 Then maps of spectral features were constructed centered on these approximate positions., Then maps of spectral features were constructed centered on these approximate positions.
 From the maps relative positions aud angular dimensions were determined by fitting 6vo-dimensional Gaussians ancl deconvolving them with the beam., From the maps relative positions and angular dimensions were determined by fitting two-dimensional Gaussians and deconvolving them with the beam.
 Separate maps of spectral features were obtained in all SLokes parameters: I. Q. U and V. From the first three Stokes parameters polarization maps were constructed. which show total intensity I contours and linear polarization vectors.," Separate maps of spectral features were obtained in all Stokes parameters: I, Q, U and V. From the first three Stokes parameters polarization maps were constructed, which show total intensity I contours and linear polarization vectors."
 The 1667 MllIz data were considered as independent from the 1665 MIIz data and were calibrated separately. using the feature at the radial velocity of 9.8 + (Fig.2 lower) as a reference.," The 1667 MHz data were considered as independent from the 1665 MHz data and were calibrated separately, using the feature at the radial velocity of 9.8 $^{-1}$ (Fig.2 lower) as a reference."
 The absolute position of the 1667 MlIz relerence feature was determined by the Iringe rate method independently [rom the 1665 MIIZ measurements. with a lower accuracy than (he relative position measurement accuracy.," The absolute position of the 1667 MHz reference feature was determined by the fringe rate method independently from the 1665 MHz measurements, with a lower accuracy than the relative position measurement accuracy."
 The cross-correlated. circularly polarized spectra of the 1665 MIIz line are shown in Fig., The cross-correlated circularly polarized spectra of the 1665 MHz line are shown in Fig.
 3., 3.
 LW compared with spectra taken in 1975 by Daviesetal.(LOTT) one can see many changes.," If compared with spectra taken in 1975 by \citet{davies77}
 one can see many changes."
 The strongest component is the RCP feature near 12 | with the same flux density. but with the radial velocity shifted to 12.45 +.," The strongest component is the RCP feature near 12 $^{-1}$ with the same flux density, but with the radial velocity shifted to 12.45 $^{-1}$."
 A nearby weaker component at 13 ! has disappeared since 1975., A nearby weaker component at 13 $^{-1}$ has disappeared since 1975.
 Two new RCP features have appeared. at 0.65 ! and 3.0 +.," Two new RCP features have appeared, at 0.65 $^{-1}$ and 3.0 $^{-1}$."
 In the LCP spectrum a new component at 9.4 ! has appeared. as well as a component at 0.65 !.," In the LCP spectrum a new component at 9.4 $^{-1}$ has appeared, as well as a component at 0.65 $^{-1}$."
 The 1667 MIIz spectrum shows less spectral features. and they occupy a smaller velocity range.," The 1667 MHz spectrum shows less spectral features, and they occupy a smaller velocity range."
 The total intensity eross-correlated spectra αἱ 1665 MIIZ and 1667 MllIz are shown in Fig., The total intensity cross-correlated spectra at 1665 MHz and 1667 MHz are shown in Fig.
 dab. There seems (o be no one-to-one correspondence between (he spectral features in the (wo main OLI lines.," 4a,b. There seems to be no one-to-one correspondence between the spectral features in the two main OH lines."
 No emission leatures were found in the OIL satellite lines at 1612 MIIz and 1720 MIIz., No emission features were found in the OH satellite lines at 1612 MHz and 1720 MHz.
The measurement of dillerential image motion by using Dillerential Image Motion. Monitor (DIMM) is widely spread. method for characterization of the atmospheric (OV) (27)...,"The measurement of differential image motion by using Differential Image Motion Monitor (DIMM) is widely spread method for characterization of the atmospheric (OT) \citep{dimm,iac-dimm}."
 The cdilferential image motion corresponds to the Ductuations of the dillerence of wavelron tilts (or cillerence of angles-of-arrival) in two close apertures., The differential image motion corresponds to the fluctuations of the difference of wavefront tilts (or difference of angles-of-arrival) in two close apertures.
 The DIMAL is simple and robust. instrument which suits well for long-term and. field. campaigns., The DIMM is simple and robust instrument which suits well for long-term and field campaigns.
 Interpretation of its output data based on the theory of light propagation through turbulent media is quite straightforward: variance ofthe differential image motion is unambiguously connectec with seeing Jo (?2)..," Interpretation of its output data based on the theory of light propagation through turbulent media is quite straightforward: variance of the differential image motion is unambiguously connected with seeing $\beta_0$ \citep{dimm,Martin1987}."
 However. there are several effects introducing biases aux random errors in DIMM results (?2)..," However, there are several effects introducing biases and random errors in DIMM results \citep{Toko2002,KT2007}."
 One of them arises due to the fact that CCD camera captures star images which are used for centroids estimation with non-zero exposure., One of them arises due to the fact that CCD camera captures star images which are used for centroids estimation with non-zero exposure.
 During typical exposure time 7LO ms. a wind shifts the turbulence by distance comparable or more than aperture £2.," During typical exposure time $\tau \sim 10$ ms, a wind shifts the turbulence by distance comparable or more than aperture $D$."
 Therefore estimated centroids inevitably undergo the wind smoothing., Therefore estimated centroids inevitably undergo the wind smoothing.
 Necessity of reduction for this elfect led to appearance of both theoretical (?) and experimental (7) studies. of impact of the exposure on power of the differential image motion measured with DIAIA, Necessity of reduction for this effect led to appearance of both theoretical \citep{Martin1987} and experimental \citep{Soules1996} studies of impact of the exposure on power of the differential image motion measured with DIMM.
L The methods of correction developed: in these works require an a priori information about wind or some temporal characteristics of the image motion estimated during the measurement process., The methods of correction developed in these works require an a priori information about wind or some temporal characteristics of the image motion estimated during the measurement process.
 Usually DIMAM data is corrected for this. bias using interlaced exposures., Usually DIMM data is corrected for this bias using interlaced exposures.
 This method consists of alternation of single and double exposures (for instance 5 ancl 10 ms)., This method consists of alternation of single and double exposures (for instance 5 and 10 ms).
 Secings are nieasured separately for these two sets and final estimation calculated as certain combination of them (??)..," Seeings are measured separately for these two sets and final estimation calculated as certain combination of them \citep{Sarazin1997,Toko2002}."
 Realization of the measurements. with modern CCD cameras changed situation significantly., Realization of the measurements with modern CCD cameras changed situation significantly.
 It became possible to use shorter exposures and faster frame rates., It became possible to use shorter exposures and faster frame rates.
 On the other hand. the method of the interlaced exposures can be hardly used. with these cameras.," On the other hand, the method of the interlaced exposures can be hardly used with these cameras."
 Alternative solution of the problem. consists in a nmieasurement of covariance of the image motion between adjacent frames. which 15 defined bv characteristic timescales of the process.," Alternative solution of the problem consists in a measurement of covariance of the image motion between adjacent frames, which is defined by characteristic timescales of the process."
 In this paper we describe this method and. prove theoretically its applicability., In this paper we describe this method and prove theoretically its applicability.
 Also. it is shown that the covariance can be used to estimate some ellective wind in the atmosphere.," Also, it is shown that the covariance can be used to estimate some effective wind in the atmosphere."
 In two first sections. we consider general expressions defining power of dillerential image motion and elfect of wind smoothing for a single. turbulent laver.," In two first sections, we consider general expressions defining power of differential image motion and effect of wind smoothing for a single turbulent layer."
 The dependencies of the power (variance of cdillerential. image motion) on wind speed and wind direction and instrumental parameters are studied as well., The dependencies of the power (variance of differential image motion) on wind speed and wind direction and instrumental parameters are studied as well.
 In the section 4.. we show that inregime. when the wind shear ez«2. where e is the wind speed. the wind cllects can be described as a quadratie functions of the shear.," In the section \ref{sec:short_approx}, , we show that in, when the wind shear $v \tau \ll D$, where $v$ is the wind speed, the wind effects can be described as a quadratic functions of the shear."
 This approximation allows us to establish simple relation between the wind elfeets on variance and covariance., This approximation allows us to establish simple relation between the wind effects on variance and covariance.
 The section 5. generalizes obtained results for the case of the whole atmosphere., The section \ref{sec:whole_atm} generalizes obtained results for the case of the whole atmosphere.
 Expressions of the wind effectson variance and covariance depending on some ellective wind, Expressions of the wind effectson variance and covariance depending on some effective wind
"ouce the distance is suiall as Planck leugth. 7,=(AGES~10.Pm. we have to cousider he quautui effect of eravity.","once the distance is small as Planck length, $l_p=(\hbar G/c^3)^{1\over
2}\approx 10^{-35} m$, we have to consider the quantum effect of gravity."
 Two of the most mniportaut physical effects related to the quantum eravity are the Tawkine-raciiation aud the black hole eutropy., Two of the most important physical effects related to the quantum gravity are the Hawking-radiation and the black hole entropy.
 In classical CR. black hole's properties cau )o precisely. calculated. aud the holes may be thought of astronomical objects with masses about several times of our Sun.," In classical GR, black hole's properties can be precisely calculated, and the holes may be thought of astronomical objects with masses about several times of our Sun."
 In this classical case. event horizon emerges. aud anything Gucluding the light) can not escape frou it to arrive at a particular observer who is outside he horizon.," In this classical case, event horizon emerges, and anything (including the light) can not escape from it to arrive at a particular observer who is outside the horizon."
" But a surprise happened: When ο, Πανάς studied the Dekeustem-clatiou of lack hole. he found that black hole cuits radiations aud the radiation spectra is just the dack bodys."," But a surprise happened: When S. Hawking studied the Bekenstein-relation of black hole, he found that black hole emits radiations and the radiation spectra is just the black body's."
 This fact iudicates that the black hole las a temperature whose expression where # is surface eravity on the horizon., This fact indicates that the black hole has a temperature whose expression where $\kappa$ is surface gravity on the horizon.
 Thus the Dekensteim-relation becoues the “real” thermodvuaiical relationship of the black hole.," Thus the Bekenstein-relation becomes the ""real"" thermodynamical relationship of the black hole."
 Namely the black hole has a thermodvuaimical eutropy as follows where Spy is the Bekeustein-Wawking entropy and Ais the area of the event horizon., Namely the black hole has a thermodynamical entropy as follows where $S_{BH}$ is the Bekenstein-Hawking entropy and $A$ is the area of the event horizon.
 The Quantum Mechanices(Q\D and the Quantmn Field Theory (OFT) in the curved space-time with (classical) event horizon provides a framework to wuderstand the statistic nechanics origin of the thermmodvuamiics of non-extreme black hole. aud serves as a powerful ool to calculate its entropy aud the asking radiationikazinarkinainBodyCitationStart305|6|.," The Quantum Mechanics(QM) and the Quantum Field Theory (QFT) in the curved space-time with (classical) event horizon provides a framework to understand the statistic mechanics origin of the thermodynamics of non-extreme black hole, and serves as a powerful tool to calculate its entropy and the Hawking radiation."
. In this framework. the gravity fields are backeround Ποια».," In this framework, the gravity fields are background fields."
 The quautuim fields move ou this classical backeround which will be not affected by he quautum fields., The quantum fields move on this classical background which will be not affected by the quantum fields.
 It is essential that the black hole backeround has coordinate singularity at the horizon., It is essential that the black hole background has coordinate singularity at the horizon.
" This singularity will seriously plague our understanding of the black ho""BS hermodvuauics:Firstly. in the eutropy calculations. an ultraviolet cutoff (or brick wall) las o be put iu by (even though it could be done in a proper way [2]))."," This singularity will seriously plague our understanding of the black hole's thermodynamics:Firstly, in the entropy calculations, an ultraviolet cutoff (or brick wall) has to be put in by (even though it could be done in a proper way )."
 Otherwise. the eutropv will be divergeut: Secoudly. in the Wawhking radiation derivations. one should also artificially use au analytic ουήλιο trick to eo over the singularity associating with that[19/7].," Otherwise, the entropy will be divergent; Secondly, in the Hawking radiation derivations, one should also artificially use an analytic continuing trick to go over the singularity associating with that."
.. This disease spoils our understanding of Tawking radiation as quautuni tunnelling effects. and some ambiguities will be left due to it.," This disease spoils our understanding of Hawking radiation as quantum tunnelling effects, and some ambiguities will be left due to it."
 Thus we face an unsatisfactory situation that if once set up a brick wall for ectting the entropy rightly. we will have no wav to do analytic continuing aud hence no wav for deriviug the Tawking temperature in this formalism.," Thus we face an unsatisfactory situation that if once set up a brick wall for getting the entropy rightly, we will have no way to do analytic continuing and hence no way for deriving the Hawking temperature in this formalism."
 Ou the other haud. if once withdrew the wall (or the ultraviolet cutoff). we will have no wax to handle the cutropy even though the Tawking temperature deriving become doable.," On the other hand, if once withdrew the wall (or the ultraviolet cutoff), we will have no way to handle the entropy even though the Hawking temperature deriving become doable."
 Cousequeutly. a self-consisteut QET (or OAT) model in curved space to calculate both the black hole's cutropy aud the Tawking temperature sinultaneouslv aud rightly is still iu absence.," Consequently, a self-consistent QFT (or QM) model in curved space to calculate both the black hole's entropy and the Hawking temperature simultaneously and rightly is still in absence."
,.
" This way, the number of galaxies of a given morphological type in a mass or luminosity bin is simply given by its mathematical expectation, and the |—o error is the square root of the variance: All the galaxies contribute to the mass function of a given morphological type weighted by its probability."," This way, the number of galaxies of a given morphological type in a mass or luminosity bin is simply given by its mathematical expectation, and the $1-\sigma$ error is the square root of the variance: All the galaxies contribute to the mass function of a given morphological type weighted by its probability."
" As a result, a galaxy that is Sd and E will still contribute to the mass function of elliptical galaxies with a weight of 0.005."," As a result, a galaxy that is Sd and E will still contribute to the mass function of elliptical galaxies with a weight of 0.005."
 Another approach is to make probability cuts., Another approach is to make probability cuts.
" This way, we decide that galaxies belong to a given class by applying a probability threshold."," This way, we decide that galaxies belong to a given class by applying a probability threshold."
This approach (even if not optimal) should be closer to the classical approach from visual classifications in which galaxies only contribute in one given class.,This approach (even if not optimal) should be closer to the classical approach from visual classifications in which galaxies only contribute in one given class.
 The threshold to apply depends on the application., The threshold to apply depends on the application.
" For example, it is interesting to determine which threshold is the best to get similar distributions than with visual classifications."," For example, it is interesting to determine which threshold is the best to get similar distributions than with visual classifications."
" In figure we compare the two estimations of the observed distribution[I2], of stellar masses with the ones obtained from the visual classification of ?.."," In figure \ref{fig:mass_counts_nair}, we compare the two estimations of the observed distribution of stellar masses with the ones obtained from the visual classification of \cite{Nair10}."
" We use a threshold of Pr;,5,,>0.45 in each type and obtain similar distributions for all morphological types.", We use a threshold of $P_{Ttype}>0.45$ in each type and obtain similar distributions for all morphological types.
evaluated with a fixed smoothing of«=0.0125 length units. which is small enough. to follow the overall dynamics of 10 interacting galaxies: the nuclei of the resulting merecr remnants are poorly resolved. but the cisks formed in these simulations are an order of magnitude larger than ο and lus not severely compromised by gravitational smoothing.,"evaluated with a fixed smoothing of $\epsilon = 0.0125$ length units, which is small enough to follow the overall dynamics of the interacting galaxies; the nuclei of the resulting merger remnants are poorly resolved, but the disks formed in these simulations are an order of magnitude larger than $\epsilon$ and thus not severely compromised by gravitational smoothing."
" IIvdrodynamic forces were calculated by smoothing. over 40 gas particles: cach gas particle has a mass of 2.I""x1.0.10 mass units. so the SPILL calculation has a mass resolution of 7.610. 7."," Hydrodynamic forces were calculated by smoothing over $40$ gas particles; each gas particle has a mass of $2^{-19} \simeq 1.9
\times 10^{-6}$ mass units, so the SPH calculation has a mass resolution of $7.6 \times 10^{-5}$ ."
" For gas with a density p and sound speed e,=0.0966. the Jeans length is Ayc0.1712p.7 and the Jeans mass is M4c0.002630> the gravitational force calculation resolves Ay if p£i190. while the SPIEL calculation resolves Ma if pX1200."," For gas with a density $\rho$ and sound speed $c_{\rm s} = 0.0966$, the Jeans length is $\lambda_{\rm J}
\simeq 0.1712 \rho^{-1/2}$ and the Jeans mass is $M_{\rm J} \simeq
0.00263 \rho^{-1/2}$; the gravitational force calculation resolves $\lambda_{\rm J}$ if $\rho \la 190$, while the SPH calculation resolves $M_{\rm J}$ if $\rho \la 1200$."
 Thus the collapse of small-scale structure in the gas is limited. by the spatial resolution of the gravitational force calculation. with hvedrocdvnamic resolution plaving only a seconcary role.," Thus the collapse of small-scale structure in the gas is limited by the spatial resolution of the gravitational force calculation, with hydrodynamic resolution playing only a secondary role."
 Suppression of small-scale collapse. which woutleL be fatal for studies of fragmentation via Jeans instabilities ος. Bate Burkert 1997). may even benefit these simulations by preserving the smooth structure of the gas.," Suppression of small-scale collapse, which would be fatal for studies of fragmentation via Jeans instabilities (e.g. Bate Burkert 1997), may even benefit these simulations by preserving the smooth structure of the gas."
 The code does suppress the Jeans instability on scales Ay2e: on such scales. the gas is usually stabilized by the combined effects of pressure and rotation (Toomre 1964).," The code does suppress the Jeans instability on scales $\lambda_{\rm J} > \epsilon$; on such scales, the gas is usually stabilized by the combined effects of pressure and rotation (Toomre 1964)."
 Nonetheless. these simulations are in fact quite crude.," Nonetheless, these simulations are in fact quite crude."
 The SPL smoothing length often exceeds the vertical scale height of the simulated gas disks: in elfect. these structures are resolved in the radial direction. but not in the vertical direction.," The SPH smoothing length often exceeds the vertical scale height of the simulated gas disks; in effect, these structures are resolved in the radial direction, but not in the vertical direction."
 Under such circumstances the caleulation provides only a rough approximation to the dynamics of a smooth eas. and it seems better to view the SPII code as a locally momentum-conserving scheme in which particles representing gas tend to seek out closed. non-intersecting orbits.," Under such circumstances the calculation provides only a rough approximation to the dynamics of a smooth gas, and it seems better to view the SPH code as a locally momentum-conserving scheme in which particles representing gas tend to seek out closed, non-intersecting orbits."
 At least an order of magnitude more gas particles are needed to significantly improve this situation., At least an order of magnitude more gas particles are needed to significantly improve this situation.
 While there'sno simple way to anticipate the results of such ambitious, While there'sno simple way to anticipate the results of such ambitious
The GeV flares produced in this case correlate tightly with the X-ray flare. i.e. N-vav aud high cucrey flares have simular temporal profiles and durations. because they come from) the same cussion region and electron population.,"The GeV flares produced in this case correlate tightly with the X-ray flare, i.e., X-ray and high energy flares have similar temporal profiles and durations, because they come from the same emission region and electron population."
 Ou the other hand. in the previous case of an afterglow IC enission due to inner-origin X-ray flares illuminating the afterelow electrous. the ταν flare aud IC omission are produced in different regions. aud the duration of the high energv flare is determined by the outer afterglow shock ecometry. R/2T?e~f. leading to a duration longer than that of the N-rav flare (0f< #).," On the other hand, in the previous case of an afterglow IC emission due to inner-origin X-ray flares illuminating the afterglow electrons, the X-ray flare and IC emission are produced in different regions, and the duration of the high energy flare is determined by the outer afterglow shock geometry, $\sim R/2\Gamma^2c\sim t$, leading to a duration longer than that of the X-ray flare $\delta t<t$ )."
 This difference provides a useful approach to distinguish these two different mo0dels of the X-ray flares. using future observations of (ιοΤον flares.," This difference provides a useful approach to distinguish these two different models of the X-ray flares, using future observations of GeV-TeV flares."
 Several imiodels have been proposed to explain the delaved CeV. ciission from GBRD910217.. including au electron IC cinission scenario in the forward shocks (Mésszárros Rees 1991. Zhang Alésszárros 2001). a. hadrou process scenario (c.g. hatz 1991: Totani 1998) aud a scenario invoking electromagnetic cascade processes of TeV οσατας in the infrared/1ücrowave backerouud (Cheng Cheng 1996: Dai Lu 2002: Wang ct al.," Several models have been proposed to explain the delayed GeV emission from GRB940217, including an electron IC emission scenario in the forward shock (Mésszárros Rees 1994, Zhang Mésszárros 2001), a hadron process scenario (e.g. Katz 1994; Totani 1998) and a scenario invoking electromagnetic cascade processes of TeV gamma-rays in the infrared/microwave background (Cheng Cheng 1996; Dai Lu 2002; Wang et al."
 20014: Razzaque. AMésszárros aud Zhang 2001).," 2004; Razzaque, Mésszárros and Zhang 2004)."
 Here we have suggested a new model. ic. the electrou IC emission accompanying the late time N-rav flares recently discovered in GRBs.," Here we have suggested a new model, i.e. the electron IC emission accompanying the late time X-ray flares recently discovered in GRBs."
 This model can be tested with high euergv detectors such as ACGILE or GLAST operating simultancously withSwift., This model can be tested with high energy detectors such as AGILE or GLAST operating simultaneously with.
 TeV eununseeavs accompanying bright N-rayv flares are also expected to be detected by ground-based telescope such as ILE.S.S.. MAGIC. VERITAS. ARGO.," TeV gamma-rays accompanying bright X-ray flares are also expected to be detected by ground-based telescope such as H.E.S.S., MAGIC, VERITAS, ARGO."
 Due to the delay times up to 100101 seconds of the N-rav flares after the trieeer. ground-based telescopes mav have suffiicicut time to slew to the directions of the bursts.," Due to the delay times up to $10^3-10^4$ seconds of the X-ray flares after the trigger, ground-based telescopes may have sufficient time to slew to the directions of the bursts."
" Tu παν, we lave studied the IC emission associated with N-rav fares in GRBs."," In summary, we have studied the IC emission associated with X-ray flares in GRBs."
 We fined the following (1) The fluence of such high cucrey flares is suffücieutlv hieh for detection bv high cucerey detectors such as GLAST., We find the following (1) The fluence of such high energy flares is sufficiently high for detection by high energy detectors such as GLAST.
 The delaved GeV cussion from GRD910217 could have becu produced in this process in a strong GRB., The delayed GeV emission from GRB940217 could have been produced in this process in a strong GRB.
 The light curves of such CeV-TeV flares may not correlate with that of the N-ray flares. but will instead become smoother aud (2) Strong N-rav. flare fluxes can enlace the cooling of the forward shock electrons. lence suppressing the svuchrotrou afterglow enüssiou during the periods of the flares.," The light curves of such GeV-TeV flares may not correlate with that of the X-ray flares, but will instead become smoother and (2) Strong X-ray flare fluxes can enhance the cooling of the forward shock electrons, hence suppressing the synchrotron afterglow emission during the periods of the flares."
 This will cause diuuumg or changing of the decay slopes of the optical afterglow., This will cause dimming or changing of the decay slopes of the optical afterglow.
 Moreover. late-time CeV-TeV detection during N-rav flares are useful for coustraining the microphysics in the afterglow (3)GeV-TeV. flares could also be produced by the selt-svuchrotron Compton emission of the flare photous in au external shock scenario of the N-rav flares.," Moreover, late-time GeV-TeV detection during X-ray flares are useful for constraining the microphysics in the afterglow (3)GeV-TeV flares could also be produced by the self-synchrotron Compton emission of the flare photons in an external shock scenario of the X-ray flares."
 Iu this case. the helt curves of the CeV-TeV flares iav. correlate with the X-ray flares themselves.," In this case, the light curves of the GeV-TeV flares may correlate with the X-ray flares themselves."
 So the different temporal behavior of the GeV-TeV flares for the two IC processes provides a potential way to distinguish whether X-ray flares originate from a late central engine activity or from an external shock scenario., So the different temporal behavior of the GeV-TeV flares for the two IC processes provides a potential way to distinguish whether X-ray flares originate from a late central engine activity or from an external shock scenario.
 We are grateful to E. Waxman. D. Zhaug. Z. C. Dai. S. Razzaque. L. J. Cou. X. Εν Wu aud Y. Z. Fan for useful discussions. aud the referee for valuable conuuceuts.," We are grateful to E. Waxman, B. Zhang, Z. G. Dai, S. Razzaque, L. J. Gou, X. F. Wu and Y. Z. Fan for useful discussions, and the referee for valuable comments."
 This work is supported by NASA NÀGB5-13286 and NSF AST 0307376. an ISE eraut and a Minerva graut (for ZL). the National Natural Science Foundation of China under grants 10103002. 10233010 and 10221001. and the Foundation for the Authors of National Excellent Doctoral Dissertations of China (for NYW).," This work is supported by NASA NAG5-13286 and NSF AST 0307376, an ISF grant and a Minerva grant (for ZL), the National Natural Science Foundation of China under grants 10403002, 10233010 and 10221001, and the Foundation for the Authors of National Excellent Doctoral Dissertations of China (for XYW)."
"Combining all these uncertainties together. the total systematic uncertainty in the LEs of the bright colmpouents of dACNSs is likely to be a factor of 3.1 over the huuiuositv range 10°105, ","Combining all these uncertainties together, the total systematic uncertainty in the LFs of the bright components of dAGNs is likely to be a factor of $3-4$ over the luminosity range $10^6-10^{8.5}\lsun$ ."
Figure 3. shows the evolution of the dAGNv frequency. e.. the ratio of the umber deusitv of simulated dLACNs selected. through the double-peaked narrow line features to the umuuber deusitv of ACNs derived from the observational total ACN LF2010).," Figure \ref{fig:f3} shows the evolution of the dAGN frequency, i.e., the ratio of the number density of simulated dAGNs selected through the double-peaked narrow line features to the number density of AGNs derived from the observational total AGN LF."
. Tere we do not distinguish between type 1 aud type 2 AGNs. but include both.," Here we do not distinguish between type 1 and type 2 AGNs, but include both."
 This ratio is equal to the ratio of the cumulative LF of dAGNs to that of ACNS. and depends weakly ou the lower limit of the AGN huninosities.," This ratio is equal to the ratio of the cumulative LF of dAGNs to that of AGNs, and depends weakly on the lower limit of the AGN luminosities."
 These estimates could be off bv a factor of 3 las previously discussed., These estimates could be off by a factor of $3-4$ as previously discussed.
 At redshift 2~0.1 0.3. the estimated dACGN frequency is ~0.1%]from (," At redshift $z\sim
0.1-0.3$ , the estimated dAGN frequency is $\sim 0.1\%-1\%$. ("
This value couldbe hieher than the estinates observed dACNs selected through the double-peaked narrow lines. considering that the signal-to-noise of some dAGN spectra are so low that the dACNs are not detected observationally.),"This value could be higher than the estimates from observed dAGNs selected through the double-peaked narrow lines, considering that the signal-to-noise of some dAGN spectra are so low that the dAGNs are not detected observationally.)"
 There are quite a nuniber of dAGNs which cannot be detected. by using double-peaked narrow lines. due to the projection effect aud the velocity cutoff.," There are quite a number of dAGNs which cannot be detected by using double-peaked narrow lines, due to the projection effect and the velocity cutoff."
 Tncliding the undetected dACNs. our model sugeests that the real frequency of dACNS is : larecr than that of those dAGNs selected through the double-peaked narrow line features by a factor of ~2.," Including the undetected dAGNs, our model suggests that the real frequency of dAGNs is $\sim 0.2\%-2\%$, larger than that of those dAGNs selected through the double-peaked narrow line features by a factor of $\sim 2$."
 These values are fully consistent with the observational estimates obtained by(2010).. ¢0.5 without completeness correction or 2.nA after the correction of in-conipleteuess due to the projection effect. velocity cutoff and low to-uolse spectra of some LAGNs ((sce also2011.. iu which the dAGNfrequency is ~0.1% without the completeness correction).," These values are fully consistent with the observational estimates obtained by, i.e., $\sim 0.1\%-0.5\%$ without completeness correction or $\sim 0.5\%-2.5\%$ after the correction of in-completeness due to the projection effect, velocity cutoff and low signal-to-noise spectra of some dAGNs (see also, in which the dAGN frequency is $\sim 0.4\%$ without the completeness correction)."
 The reasons that ai lower frequency of dACNs is obtained from our model (compared to the simple estimates: eg... 155 bv 2011)) ave three olds: (1) dAGNs can be produced only if both of the wo progenitor galaxies are gas rich so that sufficient accretion materials could be provided: (2) MDIIs in hose easvich progenitor galaxies may be initially substantially smaller than their final masses. aud thus he ACN phenomena triggered before AIBI inergers nay be substautiallve less luminous compared with those rieeered after MDII mergers: aud (3) in a siguificaut raction of major niergeers. the masses of the MDIIs rosted iu the two progenitor galaxies wav differ bv a aree factor and thus the induced dACNs mav have too arge huuinosity cdiffercuce to be included.," The reasons that a lower frequency of dAGNs is obtained from our model (compared to the simple estimates; e.g., $15\%$ by ) are three folds: (1) dAGNs can be produced only if both of the two progenitor galaxies are gas rich so that sufficient accretion materials could be provided; (2) MBHs in those gas-rich progenitor galaxies may be initially substantially smaller than their final masses, and thus the AGN phenomena triggered before MBH mergers may be substantially less luminous compared with those triggered after MBH mergers; and (3) in a significant fraction of major mergers, the masses of the MBHs hosted in the two progenitor galaxies may differ by a large factor and thus the induced dAGNs may have too large luminosity difference to be included."
 As shown in Figure 3.. the modeled frequency of spe-scale αλδν that can be detected through the double-peaked narrow line features declines to even lower values ou the order of several teu-thousanudths. c.e.. 020.06% at redshift τς~0.512.," As shown in Figure \ref{fig:f3}, the modeled frequency of kpc-scale dAGNs that can be detected through the double-peaked narrow line features declines to even lower values on the order of several ten-thousandths, e.g., $0.02\%-0.06\%$ at redshift $z\sim 0.5-1.2$."
 And the requencv of all kpc-scale dAGNs declines to 0.01%] at redshift 2—0.5, And the frequency of all kpc-scale dAGNs declines to $0.04\%-0.1\%$ at redshift $z\sim 0.5-1.2$.
 This decline is mainly due to the redshitt evolution of the galaxy morphology distribution2006)., This decline is mainly due to the redshift evolution of the galaxy morphology distribution.
. Correspoucdinely. he major mergers at such redshifts are douunated by very late-type galaxies (0.g.. Sc-Sd aud Bb galaxies) aud he fraction of mergers that host lurge initial MDITIs )econies substantialle smaller at hieh redshifts.," Correspondingly, the major mergers at such redshifts are dominated by very late-type galaxies (e.g., Sc-Sd and Irr galaxies) and the fraction of mergers that host large initial MBHs becomes substantially smaller at high redshifts."
 Euture observations which would specifically search for biel redshitt dACNs could test this prediction., Future observations which would specifically search for high redshift dAGNs could test this prediction.
 Tn addition. we note here that the estimates of the dAGN LF and the separation distribution may be affected by some sinunplified. assumptions made above.," In addition, we note here that the estimates of the dAGN LF and the separation distribution may be affected by some simplified assumptions made above."
 Below we diseuss how our results would be affected by some other mechanisms. which ma trigecr nuclear activities. but is ignored in the model above.," Below we discuss how our results would be affected by some other mechanisms, which may trigger nuclear activities, but is ignored in the model above."
 First. we have ΟΤΑ assuned that siguificaut unclear activities can onlybe triggered iu the eas-rich conrponeut of major merecr.," First, we have explicitly assumed that significant nuclear activities can only be triggered in the gas-rich component of a major merger."
" If we relax this assuuption and allow sieuificanta uuclear activities to be triggered in the gaspoor coniponenut of major merger. then the estimated αλ frequeney would rise to about ten perceut or higher at luminosity LiiZLOL, but is not affected much at lower bhuninosities."," If we relax this assumption and allow significant nuclear activities to be triggered in the gas-poor component of a major merger, then the estimated dAGN frequency would rise to about ten percent or higher at luminosity $L_{\rm [OIII]}\ga 10^{8.5}L\sun$, but is not affected much at lower luminosities."
 Under this assunrption. the dACNs at the high-bIunünositv end are primarily formed byw major merecrs of two red progenitor galaxies as the initial AIBUs in red ealaxies are substantialle big compared to those in blue galaxies with simular total stellar masses.," Under this assumption, the dAGNs at the high-luminosity end are primarily formed by major mergers of two red progenitor galaxies as the initial MBHs in red galaxies are substantially big compared to those in blue galaxies with similar total stellar masses."
 Mixed morecrs of a red progenitor galaxy with a blue progenitor ealaxy contribute little to the population of dACNs with comparable luminosities. as the masses of the initial AIDIIs iu these two progenitor galaxies are less likely to be comparable.," Mixed mergers of a red progenitor galaxy with a blue progenitor galaxy contribute little to the population of dAGNs with comparable luminosities, as the masses of the initial MBHs in these two progenitor galaxies are less likely to be comparable."
 Second. a dACN can also be produced atf larger separations if nuclear activities could be triggeredby tidal iuteractious when the two progenitor galaxies of a majormerecr are still far awav from cach other.," Second, a dAGN can also be produced at larger separations if nuclear activities could be triggeredby tidal interactions when the two progenitor galaxies of a majormerger are still far away from each other."
 Under, Under
 The fforest in quasar spectra (see c.g. Rauch 1998 for a review) offers a useful means to probe the deusity aud ionization structure of the Ineh redshift intergalactic medium.,] The forest in quasar spectra (see e.g. Rauch 1998 for a review) offers a useful means to probe the density and ionization structure of the high redshift intergalactic medium.
 The predictions of cosmological simulations (Con 1991: Zhaug. Auninos. Norman 1995: Petitjean. Müccket. Kates 1995: Hernauist 1996: Katz 1996: Wadsley Boud 1997: Theuus 19985: Dave et al.," The predictions of cosmological simulations (Cen 1994; Zhang, Anninos, Norman 1995; Petitjean, Müccket, Kates 1995; Hernquist 1996; Katz 1996; Wadsley Bond 1997; Theuns 1998; Davé et al."
 1999) have been secu to match observational data well foe. Viel 2001).," 1999) have been seen to match observational data well (e.g., Viel 2004)."
" In current cosmological theories. the fforest is produced by the remnant ueutral hydrogen in a smoothly fluctuating density field of largely photoionized eas,"," In current cosmological theories, the forest is produced by the remnant neutral hydrogen in a smoothly fluctuating density field of largely photoionized gas."
 The fluctuations in the ooptical depth 7 are predicted (sce c.g... Weinberg 1997. Croft 1997) to be related to the eas density (p) at cach point (x) aud inversely proportional to the intensity J of the ionizing backeround radiation field: Close to the brightest sources. quasars. the effect of increased J dis strong enough that it has been possible to detect the decrease in fforest absorption alone the same quasar sightline. the called proximity effect (Bajtlik 1988. Scott 20003.," The fluctuations in the optical depth $\tau$ are predicted (see e.g., Weinberg 1997, Croft 1997) to be related to the gas density $\rho$ ) at each point ${\bf x}$ ) and inversely proportional to the intensity $J$ of the ionizing background radiation field: Close to the brightest sources, quasars, the effect of increased $J$ is strong enough that it has been possible to detect the decrease in forest absorption along the same quasar sightline, the so-called proximity effect (Bajtlik 1988, Scott 2000)."
 Quasar radiation is expected to travel not only along the line of sight. but also trausverse to it. aud affect the ionization state of the IGAL which can be probed by adjacent sightlines.," Quasar radiation is expected to travel not only along the line of sight, but also transverse to it, and affect the ionization state of the IGM which can be probed by adjacent sightlines."
 This effect. known as the trausverse. or foreground proximity effect has not vet been seen convincingly in bydrogen fforest data (e.g... Schirber 200[. Croft 2001. although see Dobrzveld Bechtold 1991).," This effect, known as the transverse, or foreground proximity effect has not yet been seen convincingly in hydrogen forest data (e.g., Schirber 2004, Croft 2004, although see Dobrzycki Bechtold 1991)."
 Iu general. the presence ofdiscrete sources of photoionizius radiation will lead to spatial Suctuations in J(x).," In general, the presence of discrete sources of photoionizing radiation will lead to spatial fluctuations in $J({\bf x})$."
 The statistical properties of these have been explored by may authors. including for example Zuo (1992) aud Moeiksn White (2003).," The statistical properties of these have been explored by many authors, including for example Zuo (1992) and Meiksin White (2003)."
 Stroug sources such as quasars will leave particularly recognizable imprints in the intergalactic radiation field (see Figure 5 of Croft 2001)., Strong sources such as quasars will leave particularly recognizable imprints in the intergalactic radiation field (see Figure 5 of Croft 2004).
 Their size. shape and streugth can be probed using multiple sspeetruni sightlines. in à generalization of the trausverse sroxtnuity effect.," Their size, shape and strength can be probed using multiple spectrum sightlines, in a generalization of the transverse proximity effect."
 Adelberger (2001) has shown how sightlines chosen avound observed quasars eau be used to measure he radiative histories of quasars through the trausverse oxinity effect., Adelberger (2004) has shown how sightlines chosen around observed quasars can be used to measure the radiative histories of quasars through the transverse proximity effect.
" We choose to use the term ""light. echo” o describe these features because although they are not directly analogous to supernova light echoes (c.g.. Crotts 1989) the light from quasars is still detectable hrough them after the quasars themselves have become ""quiet."," We choose to use the term “light echo” to describe these features because although they are not directly analogous to supernova light echoes (e.g., Crotts 1989) the light from quasars is still detectable through them after the quasars themselves have become “quiet”."
 We will explün in 822 how we can find the light echoes aud so detect quasars aud vield coustraiuts ou their xoperties even when the Leht from them is uot directly observable., We will explain in 2 how we can find the light echoes and so detect quasars and yield constraints on their properties even when the light from them is not directly observable.
 There are several examples of extended objects which rave been searched for iu cosiological datasets with some vpe of matched filtering., There are several examples of extended objects which have been searched for in cosmological datasets with some type of matched filtering.
 These include galaxy. clusters at Heh redshift (Postinan 1996)., These include galaxy clusters at high redshift (Postman 1996).
 Circles in the CMD sky. which would be a sign of a universe with a small opolosv scale have also been looked for (Cornish 2001).," Circles in the CMB sky, which would be a sign of a universe with a small topology scale have also been looked for (Cornish 2004)."
 We will use a simular idea to fiud aud measure light echoes in forest data. searching for the signature of a deficit of aabsorption with a particular geometry.," We will use a similar idea to find and measure light echoes in forest data, searching for the signature of a deficit of absorption with a particular geometry."
 The of lifetimesquasars are poorly relativelycoustrained. aud observatioually the evideuce poiuts to the range 109LO Yr (see Martini 2003 for a review).," The lifetimes of quasars are relatively poorly constrained, and observationally the evidence points to the range $10^6-10^8$ Yr (see Martini 2003 for a review)."
 Theoretical predictions of quasar lifetimes have receuthly become available from hydrodvuaiic simulatious of galaxy formation which inclide, Theoretical predictions of quasar lifetimes have recently become available from hydrodynamic simulations of galaxy formation which include
ionization front arrives at the proto-LG position from the direction of Virgo and quickly sweeps through it as illustrated by the next three images.,ionization front arrives at the proto-LG position from the direction of Virgo and quickly sweeps through it as illustrated by the next three images.
 This corresponds to the dramatic jump of the LG local ionized fraction seen in Figure 4 at z=10.5— 10.25.," This corresponds to the dramatic jump of the LG local ionized fraction seen in Figure \ref{reion_hist_fig}
 at $z=10.5-10.25$ ."
" By z=10 the Local Group material is almost fully ionized, reaching ionized fraction by mass."," By $z=10$ the Local Group material is almost fully ionized, reaching ionized fraction by mass."
" In this case, therefore the Local Group is largely ionized from the outside, primarily by the Virgo progenitors."," In this case, therefore the Local Group is largely ionized from the outside, primarily by the Virgo progenitors."
" The evolution proceeds quite differently in the photon-poor, Model 2, case, as illustrated in Figure 7.."," The evolution proceeds quite differently in the photon-poor, Model 2, case, as illustrated in Figure \ref{3D_cross_HI_wmap5_fig}."
" While the nearby clusters, in this case both Virgo and Fornax, once again reionize themselves from the inside and relatively earlier than the Local Group, there are no clear ionization fronts to arrive from them and sweep over the LG material."," While the nearby clusters, in this case both Virgo and Fornax, once again reionize themselves from the inside and relatively earlier than the Local Group, there are no clear ionization fronts to arrive from them and sweep over the LG material."
 The Local Group internal sources carve ionized bubbles from the inside and eventually manage to reionize all the LG material mostly by themselves., The Local Group internal sources carve ionized bubbles from the inside and eventually manage to reionize all the LG material mostly by themselves.
" While we cannot exclude modest contributions from Virgo and Fornax, the local LG sources appear to dominate the evolution in this case."," While we cannot exclude modest contributions from Virgo and Fornax, the local LG sources appear to dominate the evolution in this case."
" In order to evaulate the local reionization process for each structure in a more quantitative way, we counted and added up all ionizing photons emitted by sources within the same Lagrangean volume, normalized by the total number of atoms belonging to that object."," In order to evaulate the local reionization process for each structure in a more quantitative way, we counted and added up all ionizing photons emitted by sources within the same Lagrangean volume, normalized by the total number of atoms belonging to that object."
" We also counted and added together the cumulative number of ionizing photons used up to ionize each object of interest and also to keep it ionized (i.e. recombinations, since after each recombination back to neutral state that atom would need to be ionized again), again normalized per atom in that object."," We also counted and added together the cumulative number of ionizing photons used up to ionize each object of interest and also to keep it ionized (i.e. recombinations, since after each recombination back to neutral state that atom would need to be ionized again), again normalized per atom in that object."
 Comparing the values of these two numbers over time shows if that particular object by itself produced enough photons up to that point in time to fully account for its current ionization state., Comparing the values of these two numbers over time shows if that particular object by itself produced enough photons up to that point in time to fully account for its current ionization state.
 Results for both Models are shown in Figure 8.., Results for both Models are shown in Figure \ref{photon_budget_fig}.
 The results confirm our conclusions based on the visual examination of the reionization process we discussed above., The results confirm our conclusions based on the visual examination of the reionization process we discussed above.
 The galaxy clusters initially produce most of the photons needed for their own reionization., The galaxy clusters initially produce most of the photons needed for their own reionization.
" Until z10.5 (z9— 9.5) for Model 1 (Model 2) there is some deficit of photons, i.e. a little more are used up than produced by the progenitors of that cluster."," Until $z\sim10.5$ $z\sim9-9.5$ ) for Model 1 (Model 2) there is some deficit of photons, i.e. a little more are used up than produced by the progenitors of that cluster."
" The reason for this is that clusters are located at high peaks of the density field, and it is well established that CDM halos, and thus our ionizing sources, strongly cluster around such high peaks."," The reason for this is that clusters are located at high peaks of the density field, and it is well established that CDM halos, and thus our ionizing sources, strongly cluster around such high peaks."
 As a consequence of that there are many nearby sources which surround the proto-cluster region and contribute to its reionization., As a consequence of that there are many nearby sources which surround the proto-cluster region and contribute to its reionization.
" However, we note that the final overlap in each case is only reached after the cluster itself has produced sufficient number of photons."," However, we note that the final overlap in each case is only reached after the cluster itself has produced sufficient number of photons."
" We also note that recombinations have a very significant effect for proto-clusters, yielding usage of 1-1.5 additional ionizing photons per atom in addition to the one photon needed to achieve the initial ionization of that atom."," We also note that recombinations have a very significant effect for proto-clusters, yielding usage of 1-1.5 additional ionizing photons per atom in addition to the one photon needed to achieve the initial ionization of that atom."
" On the other hand, the proto-Local Group results are different from the cluster results in both cases."," On the other hand, the proto-Local Group results are different from the cluster results in both cases."
 In the photon-rich Model 1 up to z10.5 there are exactly as many photons produced by the LG progenitors as areused up for its own reionization., In the photon-rich Model 1 up to $z\sim10.5$ there are exactly as many photons produced by the LG progenitors as areused up for its own reionization.
 After that, After that
two in one dinmensiou.,two in one dimension.
 We discuss the streueth of our scheme while comparing will other sclemes eiven in literatures., We discuss the strength of our scheme while comparing with other schemes given in literatures.
 Suisalu&Saar(1995) took another approach when solving the Poisson equation in their AMR computation., \citet{suisalu95} took another approach when solving the Poisson equation in their AMR computation.
 They adopted au cubic interpolation formula at the grid level botucdary to eusure the continuity of the solution., They adopted an cubic interpolation formula at the grid level boundary to ensure the continuity of the solution.
 They have noticed that a linear interpolation eusures only the coutinuity in the potenti but not that in the gravitational force., They have noticed that a linear interpolation ensures only the continuity in the potential but not that in the gravitational force.
 Discontinuity iu the force may cause a spurious feature across the grid level boundary., Discontinuity in the force may cause a spurious feature across the grid level boundary.
 Although not mentioned explicitly iu (1995).. the discontinuity causes a more serious problem: it introduces spurious iuass on the erid level boundary.," Although not mentioned explicitly in \citet{suisalu95}, the discontinuity causes a more serious problem; it introduces spurious mass on the grid level boundary."
 [f the gravity g lias ciflerent values at a erid level boundary. the solution does uot satisfies the Ciausss theorem.," If the gravity $ \mbox{\boldmath$ $} $ has different values at a grid level boundary, the solution does not satisfies the Gauss's theorem."
 The spurious mass ou the erid boundary. gives a serious error ou the solution in a long rauge., The spurious mass on the grid boundary gives a serious error on the solution in a long range.
 The error decreases ouly inversely proportioual to the distance [rom the spurious mass., The error decreases only inversely proportional to the distance from the spurious mass.
 This error domiuates over the gravitational torque which decreases more steeply., This error dominates over the gravitational torque which decreases more steeply.
 Compared with adoption of a higher order interpolation formula at the grid level bouudary. our scheme lias several advantages.," Compared with adoption of a higher order interpolation formula at the grid level boundary, our scheme has several advantages."
 First our approach ensures abseuce of a spurious mass on the erid level boundary while it is not guarauteed in a higher oder interpolation formula., First our approach ensures absence of a spurious mass on the grid level boundary while it is not guaranteed in a higher oder interpolation formula.
 Secoud higher order interpolation formula increases computation load., Second higher order interpolation formula increases computation load.
 Third higher order interpolation may slow down ellicieucy of convergence., Third higher order interpolation may slow down efficiency of convergence.
 As [ar as in our experience. a simple multi-erid iteration does not work when Equatious (19)) aud (20)) are replaced with a higher order interpolation formula.," As far as in our experience, a simple multi-grid iteration does not work when Equations \ref{interpolation1}) ) and \ref{interpolation2}) ) are replaced with a higher order interpolation formula."
 We discuss the thircl potut further in the following., We discuss the third point further in the following.
 Our discrete Polsson equation has the following favorable character., Our discrete Poisson equation has the following favorable character.
 When it is rewritten in the form. the coelficient ας is always positive except for 7 = j.," When it is rewritten in the form, the coefficient $ a _{i,j} $ is always positive except for $ i $ = $ j $."
 Here symbols 7 and j denote the cell ntunber after sorted in oue dimeusion., Here symbols $ i $ and $ j $ denote the cell number after sorted in one dimension.
 This property eusures that the Catss-Seicdel iteration converges always since it always uuderestimates the correction., This property ensures that the Gauss-Seidel iteration converges always since it always underestimates the correction.
 Iu otler words. the couvergeuce of the Gauss-Seidel iteration is ensured since the matrix aj; is diagonally dominant. When we use a higher order interpolation formula. this the dominauce of the diagonal element [Eq. (23))]," In other words, the convergence of the Gauss-Seidel iteration is ensured since the matrix $ a _{i,j} $ is diagonally dominant, When we use a higher order interpolation formula, this the dominance of the diagonal element [Eq. \ref{Ddominance}) )]"
 is lost., is lost.
" The coefficient «;,; is no longer always positive."," The coefficient $ a _{i,j} $ is no longer always positive."
 Then the Gauss-Seidel iteration may overestimate the correction aud may not converge., Then the Gauss-Seidel iteration may overestimate the correction and may not converge.
 If we apply successive tuder-relaxatiou. the couvergeuce slows dow.," If we apply successive under-relaxation, the convergence slows down."
 Rickeretal.(2000) have proposed another idea for solving the Poisson equation in a nonuuilorui eric., \citet{ricker00} have proposed another idea for solving the Poisson equation in a nonuniform grid.
 They have used a uouuniform grid in which the spacing ina certain direction varies but ouly in the direction., They have used a nonuniform grid in which the spacing in a certain direction varies but only in the direction.
 For example. the spacing in the sr-direction varies with or but not in the y aud," For example, the spacing in the $ x $ -direction varies with $ x $ but not in the $ y $ - and"
"and NSF grant AST-1041590, and by NSF through TeraGrid resources provided by QueenBee of the Louisiana Optical Network Initiative (http://www.loni.org) and Kraken of the National Institute for Computational Sciences under grant number TG-AST080026N (RN AT) and TG-AST100040 (AT).","and NSF grant AST-1041590, and by NSF through TeraGrid resources provided by QueenBee of the Louisiana Optical Network Initiative ) and Kraken of the National Institute for Computational Sciences ) under grant number TG-AST080026N (RN AT) and TG-AST100040 (AT)."
 AT was supported in part by a Princeton Center for Theoretical Science Fellowship., AT was supported in part by a Princeton Center for Theoretical Science Fellowship.
 PK was supported in part by NSF grant AST-0909110., PK was supported in part by NSF grant AST-0909110.
" A cold hydrodynamic flow can have a shock for any choice of the upstream velocity u,.", A cold hydrodynamic flow can have a shock for any choice of the upstream velocity $u_u$.
 This is because no signals can propagate in the cold gas and so the upstream gas is always supersonic., This is because no signals can propagate in the cold gas and so the upstream gas is always supersonic.
 A magnetized fluid is different., A magnetized fluid is different.
" Even if the gas is cold, Alfven and fast magnetosonic waves can still propagate in the fluid."," Even if the gas is cold, Alfven and fast magnetosonic waves can still propagate in the fluid."
 Thus a shock is possible only if the upstream gas moves faster than these waves., Thus a shock is possible only if the upstream gas moves faster than these waves.
" For the particular geometry we have considered, viz.,"," For the particular geometry we have considered, viz.,"
" a perpendicular shock with magnetic field perpendicular to the velocity vector, the relevant wave speed is that of the fast magnetosonic wave, which is given (in the comoving frame of the gas) by Thus we can have a shock only if We now consider a number of limiting cases."," a perpendicular shock with magnetic field perpendicular to the velocity vector, the relevant wave speed is that of the fast magnetosonic wave, which is given (in the comoving frame of the gas) by Thus we can have a shock only if We now consider a number of limiting cases."
"When c, we expect to have a strongshock, whereas when iu is only marginally greater than c, we expect a weak shock.","When $u_u^2 \gg \sigma$ , we expect to have a strongshock, whereas when $u_u^2$ is only marginally greater than $\sigma$, we expect a weak shock."
" In addition, we have different results depending on whether i2>>1 (ultrarelativistic) or u2«1 (nonrelativistic), and on whether o>1 (strongly magnetized) or 7«1 (weakly magnetized)."," In addition, we have different results depending on whether $u_u^2 \gg 1$ (ultrarelativistic) or $u_u^2\ll 1$ (nonrelativistic), and on whether $\sigma \gg 1$ (strongly magnetized) or $\sigma \ll 1$ (weakly magnetized)."
" considered a couple of important cases, but here we present [KC84]scalings for all the different regimes."," considered a couple of important cases, but here we present scalings for all the different regimes."
 Figure identifies the regimes and labels them by the respective subsection where each is discussed., Figure \ref{fig4} identifies the regimes and labels them by the respective subsection where each is discussed.
" This case has been considered by who give the following results: where o(x) Tas)denotes lam)terms of higher order than x, ie., o(x)/x—O0 as x- 0."," This case has been considered by who give the following results: where $\os{x}$ denotes terms of higher order than $x$, i.e., $\os{x}/x\to0$ as $x\to0$ ."
" In the above, we remind the reader that the magnetic field strengths B, and By are measured in the shock frame."," In the above, we remind the reader that the magnetic field strengths $B_u$ and $B_d$ are measured in the shock frame."
 These results are obtained by expanding the jump conditions as power series in the small quantity 1/σ. and matching terms of similar order.," These results are obtained by expanding the jump conditions as power series in the small quantity $1/\sigma$, and matching terms of similar order."
" We have set h(6;)= 4, as appropriate for relativistically hot downstream gas."," We have set $h(\theta_d)=4$ , as appropriate for relativistically hot downstream gas."
" Our result for 0;/u, differs from that given in[", Our result for $\theta_d/u_u$ differs from that given in.
KC84.. As i? approaches c the shock becomes progressively weaker.," As $u_u^2$ approaches $\sigma$, the shock becomes progressively weaker."
" Let us write u?=o(1+A), with A«1."," Let us write $u_u^2=\sigma (1+\Delta)$, with $\Delta \ll 1$."
" As A becomes progressively smaller, less and less of the kinetic energy of the upstream gas is thermalized in the shock."," As $\Delta$ becomes progressively smaller, less and less of the kinetic energy of the upstream gas is thermalized in the shock."
" In this limit, the downstream temperature 6, becomes non-relativistic and so we set (07)=5/2.In this limit, the leading terms in the solution are as follows: Note that the leading term in the temperature of the shocked gas goes as A?, i.e., the shock is extremely inefficient."," In this limit, the downstream temperature $\theta_d$ becomes non-relativistic and so we set $h(\theta_d)=5/2$.In this limit, the leading terms in the solution are as follows: Note that the leading term in the temperature of the shocked gas goes as $\Delta^3$ , i.e., the shock is extremely inefficient."
 This is a characteristic featureof weak shocks., This is a characteristic featureof weak shocks.
" Thiscase has again been considered by KC84]. With A(@,)= 4, the"," Thiscase has again been considered by With $h(\theta_d)=4$ , the"
data reduction.,data reduction.
 Section. ?? includes the object detection and candidate selection and a discussion about possible interlopers 1n. the sample., Section \ref{sec:sel} includes the object detection and candidate selection and a discussion about possible interlopers in the sample.
 In Sect., In Sect.
 ?? we present all the basic characteristics of the sample. including photometry. AGN contribution. surface density. and sizes. and the equivalent width distribution of the candidates.," \ref{sec:basic} we present all the basic characteristics of the sample, including photometry, AGN contribution, surface density, and sizes, and the equivalent width distribution of the candidates."
 We summarise the results in Sect. ??.., We summarise the results in Sect. \ref{sec:disc}.
" 5mm Throughout this paper. we assume à cosmology with Il, =72kms + + (Freedman et al."," 5mm Throughout this paper, we assume a cosmology with $H_0=72$ km $^{-1}$ $^{-1}$ (Freedman et al."
" 2001). O4,=0.3 and O4—0.7."," 2001), $\Omega _{\rm m}=0.3$ and $\Omega _\Lambda=0.7$."
 Magnitudes are given in the AB system., Magnitudes are given in the AB system.
 In March. 2007. à 35\34 aremin? section of the COSMOS field. centred on R.A. —1(0007275 and Dec = (02712/2277(.J2000). wasobscrecdwiththeW ik Ficlllinagev(WFI: Baadeetal," In March 2007, a $\times$ 34 $^2$ section of the COSMOS field, centred on R.A. $ = 10^h 00^m 27^s$ and Dec $ = 02^{\circ} 12' 22$ $7$ (J2000), was observed with the Wide-Field Imager (WFI; Baade et al."
 1999 onthe MPG ESO2 2intelescopeonLaSilla., 1999) on the MPG/ESO 2.2m telescope on La Silla.
AMogoftheobservationsceanbefoundinTuble ?2.., A log of the observations can be found in Table \ref{obslog}.
 The total dithered image consisted of 29 exposures with a total exposure time of 99624 seconds. or 27.7 hours.," The total dithered image consisted of 29 exposures with a total exposure time of 99624 seconds, or 27.7 hours."
 The observations were made with narrow-band filter N396/12 with a central wavelength of 396.3 nm and a FWHM of 12.9 nm., The observations were made with narrow-band filter N396/12 with a central wavelength of 396.3 nm and a FWHM of 12.9 nm.
 This wavelength range corresponds to 2=2.206—2.512 for Lya. ;:=(L016—0.081 for [OI]. and :=1.521.59 for CIV and the surveyed comoving volumes (after masking. see Sect. 22))," This wavelength range corresponds to $z = 2.206 - 2.312$ for $\alpha$ , $z = 0.046 - 0.081$ for [OII], and $z = 1.52 - 1.59$ for CIV and the surveyed comoving volumes (after masking, see Sect. \ref{selection}) )"
 are ~329 300 Ὁ. 820 ? and ~ 000 7. respectively.," are $\sim 329$ $300$ $^{-3}$, $820$ $^{-3}$ and $\sim 225$ $000$ $^{-3}$, respectively."
 The narrow-band filter curve is shown with the filter curves of the filters used for the selection of candidates in Fig. 1.., The narrow-band filter curve is shown with the filter curves of the filters used for the selection of candidates in Fig. \ref{filtercurve}.
 The data were reduced using the pipeline developed in MIDAS with specific routines to handle WFI data., The data were reduced using the pipeline developed in MIDAS with specific routines to handle WFI data.
 We briefly describe the reduction steps performed on the data., We briefly describe the reduction steps performed on the data.
 The data from the individual CCDs were converted into MIDAS format and bias-corrected on each separate CCD., The data from the individual CCDs were converted into MIDAS format and bias-corrected on each separate CCD.
 The individual CCD frames were then placed into placeholders in an empty mosaic. thus creating a full mosaic for each frame including the gaps between the CCDs.," The individual CCD frames were then placed into placeholders in an empty mosaic, thus creating a full mosaic for each frame including the gaps between the CCDs."
 To correct for bad pixels and columns as well as correeting columns with a constant offset compared to the surrounding pixels. lamp images with exposure times ranging between 1.—220 s were downloaded and analysed.," To correct for bad pixels and columns as well as correcting columns with a constant offset compared to the surrounding pixels, lamp images with exposure times ranging between $1 - 220$ s were downloaded and analysed."
 The science frames were then corrected for both the bad columns and the offset columns., The science frames were then corrected for both the bad columns and the offset columns.
 The frames were then flat-field corrected using a master flat created from sky flats taken at the time of observations., The frames were then flat-field corrected using a master flat created from sky flats taken at the time of observations.
 To determine the offsets between the images. an SDSS catalogue of known sources was used. combined with an algorithm that identifies sources anc matches them with catalogue entries.," To determine the offsets between the images, an SDSS catalogue of known sources was used, combined with an algorithm that identifies sources and matches them with catalogue entries."
 The information about the shift is then entered into the header., The information about the shift is then entered into the header.
 In the final step. the images were shifted and rebinned onto a gnomonic projectior (Le. by de-projecting all great circles onto straight lines). cosmic-ray hits were removed and a final. co-added image created.," In the final step, the images were shifted and rebinned onto a gnomonic projection (i.e. by de-projecting all great circles onto straight lines), cosmic-ray hits were removed and a final, co-added image created."
 Significant residuals from the background subtractior were apparent in the mosaiced image. in. particular. around the edges of the individual images used to create the mosaic.," Significant residuals from the background subtraction were apparent in the mosaiced image, in particular around the edges of the individual images used to create the mosaic."
 We used the following procedure to remove it., We used the following procedure to remove it.
" First. we used the ""clean"" option of the IRAF task to create a version of the image in which objects had been removed and interpolated over."," First, we used the “clean” option of the IRAF task to create a version of the image in which objects had been removed and interpolated over."
 We then edited this image interactively to remove any residual flux from objects., We then edited this image interactively to remove any residual flux from objects.
 Subsequently. we fitted a 2020 piece bicubic spline to this image.," Subsequently, we fitted a $\times$ 20 piece bicubic spline to this image."
 We ther subtracted this fit from the original image., We then subtracted this fit from the original image.
 This procedure reduced the background variations to a small fraction of the background noise., This procedure reduced the background variations to a small fraction of the background noise.
 We then flux calibrated the narrow-band image by calculating the fluxes of all narrow-band selected objects (see Sect. 22)), We then flux calibrated the narrow-band image by calculating the fluxes of all narrow-band selected objects (see Sect. \ref{selection}) )
 in the CFHT « and the SUBARU BJ images (see Table 3)) and interpolated the fluxes in the narrow- image., in the CFHT $u^*$ and the SUBARU Bj images (see Table \ref{cosmodata}) ) and interpolated the fluxes in the narrow-band image.
 We can then calculate the zero-point of the image. assuming that the median equivalent width of all objects i5 zero.," We can then calculate the zero-point of the image, assuming that the median equivalent width of all objects is zero."
 The 5c detection limit in a 3” diameter aperture in the image Is 25.3 AB magnitude and the 90% completeness limit is 25.1 AB magnitude. corresponding to Lya luminosities of logL= 12.36ergs bandlog£= I2.1lergs |. respectively. at;=2.25.," The $5 \sigma$ detection limit in a $3''$ diameter aperture in the image is $25.3$ AB magnitude and the $90$ completeness limit is $25.1$ AB magnitude, corresponding to $\alpha$ luminosities of $\log{L} = 42.36$ erg $^{-1}$ and $\log{L} = 42.44$ erg $^{-1}$, respectively, at $z = 2.25$."
 For object detection. we used the SExtractor software (Bertin Arnouts 1996).," For object detection, we used the SExtractor software (Bertin Arnouts 1996)."
 The narrow-band image was used as a detection image. and objects with a minimum of 8 adjoining pixels and a threshold of 2o per pixel were selected.," The narrow-band image was used as a detection image, and objects with a minimum of 8 adjoining pixels and a threshold of $2\sigma$ per pixel were selected."
 For the selection. we used two broad-band filters corresponding to wavelengths to both the blue and red side of the narrow-band filter. the CFHT «& band image and the SUBARU Bj band image. see also Fig. I..," For the selection, we used two broad-band filters corresponding to wavelengths to both the blue and red side of the narrow-band filter, the CFHT $u^*$ band image and the SUBARU Bj band image, see also Fig. \ref{filtercurve}."
 Both images weretaken from the public data in the COSMOS field (Capak et al., Both images weretaken from the public data in the COSMOS field (Capak et al.
 2007)., 2007).
 The « and By band images were created from combining several smaller sub- to match perfectly the field of the narrow-band image.," The $u^*$ and Bj band images were created from combining several smaller sub-images to match perfectly the field of the narrow-band image,"
where cosmological analyses usually restrict themselves.,where cosmological analyses usually restrict themselves.
 Moreover. we have shown that this correction changes only slightly between the NEAR. MID. and. FAR LRG reconstinetec halo densitv fields. while (he FOG compressed mocks have a much larger variation between samples.," Moreover, we have shown that this correction changes only slightly between the NEAR, MID, and FAR LRG reconstructed halo density fields, while the FOG compressed mocks have a much larger variation between samples."
 Therefore. we can hope to push cosmological analvses to larger A using (he reconstructed halo density field as a tracer of the underlying matter density field (Iuctiations— particularly for 5galaxy samples like the LRGs which are spread over a large5 redshift range5 and are not volume-limitecl. and thus have substantial variation in (he satellite contribution to the power spectrum wilh While we have not addressed the variation of the nonlinear correction to the reconstructed halo density field as a function of cosmological parameters. we have designed the form of our correction to minimize the variation with cosmology.," Therefore, we can hope to push cosmological analyses to larger $k$ using the reconstructed halo density field as a tracer of the underlying matter density field fluctuations— particularly for galaxy samples like the LRGs which are spread over a large redshift range and are not volume-limited, and thus have substantial variation in the satellite contribution to the power spectrum with While we have not addressed the variation of the nonlinear correction to the reconstructed halo density field as a function of cosmological parameters, we have designed the form of our correction to minimize the variation with cosmology."
 Other researchers (e.g..?7) are studying the dark matter power spectrum as a function of cosmology.," Other researchers \citep[e.g., ][]{habib/etal:2007} are studying the dark matter power spectrum as a function of cosmology."
" We expect (hat our nonlinear correction Prjc;/Poa, will remain small (of order below fk=0.2bh\Ipe+) as the cosmology is varied. so the variation of this small correction shoulcl be even smaller."," We expect that our nonlinear correction $P_{LRG}/P_{DM}$ will remain small (of order below $k = 0.2 \; h \; {\rm Mpc}^{-1}$ ) as the cosmology is varied, so the variation of this small correction should be even smaller."
 Therefore. instead of introducing a nuisance parameter lor the nonlinear correction. we propose that the amplitude of the correction should be taken as the error on ils value in cosmological parameter analyses. or relatively strong priors on the amplitude of (he correction to the reconstructed halo densitv field be In (his work we have also investigated (he properties of (he power spectrum covariance matrix for the dark matter. as well as the mock galaxy catalogs divided into 6 different saniples: central galaxies in real and redshift space. central and satellite galaxies in real and redshift space. our reconstructed halo density field in redshilt space. and the ? FOG compressed galaxy densitv field in redshift space.," Therefore, instead of introducing a nuisance parameter for the nonlinear correction, we propose that the amplitude of the correction should be taken as the error on its value in cosmological parameter analyses, or relatively strong priors on the amplitude of the correction to the reconstructed halo density field be In this work we have also investigated the properties of the power spectrum covariance matrix for the dark matter, as well as the mock galaxy catalogs divided into 6 different samples: central galaxies in real and redshift space, central and satellite galaxies in real and redshift space, our reconstructed halo density field in redshift space, and the \citet{tegmark/etal:2006} FOG compressed galaxy density field in redshift space."
 All of these samples were well-modeled bv a diagonal matrix with (he usual Gaussian aud Poisson shot noise terms plus the term presented in ?.., All of these samples were well-modeled by a diagonal matrix with the usual Gaussian and Poisson shot noise terms plus the beat-coupling term presented in \citet{hamilton/rimes/scoccimarro:2006}.
 We expect that (this will be a usefulmodel for fitting the survey covariance matrix. where (75) is replaced by Ανν) with Fi; determined by some elLective survey Finally we examined (he structure of the redshift space distortions as a function of 7 using the quadrupole.," We expect that this will be a usefulmodel for fitting the survey covariance matrix, where $P(k_b)$ is replaced by $P(k_{survey})$ with $k_{survey}$ determined by some effective survey Finally we examined the structure of the redshift space distortions as a function of $k$ using the quadrupole."
 Both the reconstructed halo density field and the FOG-compressed mock LRG samples reproduce the modest /: dependence of the halo density field cuadrupole io monopole ratio., Both the reconstructed halo density field and the FOG-compressed mock LRG samples reproduce the modest $k$ dependence of the halo density field quadrupole to monopole ratio.
 Since the LRGs are so highly. biased. this scale dependence causes a <3% deviation in the redshift space monopole to real space power spectrum ratio out io k= 0.2.," Since the LRGs are so highly biased, this scale dependence causes a $\lesssim 3\%$ deviation in the redshift space monopole to real space power spectrum ratio out to $k=0.2$ ."
 When satellites are included in the sample without FOG compression. the," When satellites are included in the sample without FOG compression, the"
are given.,are given.
 Then we describe the correction. of. measured variances to zero exposure and define an applicability aud accuracy the method., Then we describe the correction of measured variances to zero exposure and define an applicability and accuracy the method.
 Obtained formulae permit to evaluate square averaged wind Y» weighted with turbulence intensity and calculate the atmospheric coherence time 75., Obtained formulae permit to evaluate square averaged wind $\bar V_2$ weighted with turbulence intensity and calculate the atmospheric coherence time $\tau_0$.
 This is demonstrated in the Section 6.3.J) on the base of analysis of real data., This is demonstrated in the Section \ref{sec:MASSwind} on the base of analysis of real data.
 The last section contains discussion of the results and recommendations on optimal measurements ancl processing., The last section contains discussion of the results and recommendations on optimal measurements and processing.
 Details of evaluation of needed integrals over dimensional spatial frequcney are given in appenclixes., Details of evaluation of needed integrals over two-dimensional spatial frequency are given in appendixes.
 The numerical estimates are. computed. for. DIAIAL device which we use for the OT. measurement on Mount Shatdjatmaz (?).., The numerical estimates are computed for DIMM device which we use for the OT measurement on Mount Shatdjatmaz \citep{kgo2010}.
 Main. parameters of the instrument are aperture diameter. 2=0.09 m. cimensionless baseline b= 2.18. exposure 7=4 ms before 2009 December 13 and 2.5 ms after. period between exposures 7=5 ms.," Main parameters of the instrument are aperture diameter $D = 0.09$ m, dimensionless baseline $b = 2.18$ , exposure $\tau = 4$ ms before 2009 December 13 and $2.5$ ms after, period between exposures $T = 5$ ms."
 Equations for spatial spectrum. of the cilferential image motion were given in classical works by ? and 7.., Equations for spatial spectrum of the differential image motion were given in classical works by \citet{Fried1965} and \citet{Martin1987}.
 In the paper of ?.. these equations were developed: by successive filtering of initial spatial spectrum. of light wave phase distortions on the assumption of WKolmogoroy OT. model.," In the paper of \citet{Martin1987}, these equations were developed by successive filtering of initial spatial spectrum of light wave phase distortions on the assumption of Kolmogorov OT model."
 lt is known that spectral power densitv ® of phase fluctuations is proportional to the intensity AS=C2AA in homogeneous and isotropic turbulent [aver: where f is the modulus of the 2D spatial frequency. A is the wavelength and ro is the Fried parameter.," It is known that spectral power density $\Phi$ of phase fluctuations is proportional to the intensity $\Delta J = C_n^2\,\Delta h$ in homogeneous and isotropic turbulent layer: where $f$ is the modulus of the 2D spatial frequency, $\lambda$ is the wavelength and $r_0$ is the Fried parameter."
 For practical purposes AJ is more appropriate because this value is laver-byv-laver additive and used. for. characterization of vertical OT profile., For practical purposes $\Delta J$ is more appropriate because this value is layer-by-layer additive and used for characterization of vertical OT profile.
" Phe phase spectrum is axisvmmetric. i.e. depends only on the f and has the dimension m.""7."," The phase spectrum is axisymmetric, i.e. depends only on the $f$ and has the dimension $\mathrm{m}^{-2}$."
 Eq. (1)), Eq. \ref{eq:KolmPSD}) )
 is valid on condition that propagation cllects of disturbed wavelront can be neglected (near-fickel approximation)., is valid on condition that propagation effects of disturbed wavefront can be neglected (near-field approximation).
 llereinafter we use coordinate svstem cilferent from one used by 2..Phe c axis is placed along line connecting centres ofapertures., Hereinafter we use coordinate system different from one used by \citet{Martin1987}.The $x$ axis is placed along line connecting centres of apertures.
 In frequency domain. coordinates of a point are defined by modulus f. and position angle © being reckoned [rom the wr axis.," In frequency domain, coordinates of a point are defined by modulus $f$ and position angle $\phi$ being reckoned from the $x$ axis."
 As usual two components of the differential image motion are considered: longitudinal is measured along the in axis (hereafter. ο) and transverse is measured along yaris (/, As usual two components of the differential image motion are considered: longitudinal is measured along the $x$ axis (hereafter ) and transverse is measured along $y$ axis ).
-3nolion) The spectrum £ of the image motion can be represented as a product of phase spectrum. gradient. filter. aperture filter ane cüllerential filter (??)..," The spectrum $F$ of the image motion can be represented as a product of phase spectrum, gradient filter, aperture filter and differential filter \citep{Martin1987,Toko2002}."
 Formulae for these filters were traüsformed to used coordinate svstem., Formulae for these filters were transformed to used coordinate system.
 Cradient filter C for f- and (motions has the form in m+ to the anegle-of-arrival in radians. measured in the methoc.," Gradient filter $G$ for $l$ - and $t$ -motions has the form where factor $(\lambda/2\pi)^2$ is needed to pass from the slope in ${\rm m}^{-1}$ to the angle-of-arrival in radians, measured in the method."
 Expression for the aperture filter ;d depends on how image centroids are evaluated., Expression for the aperture filter $A$ depends on how image centroids are evaluated.
 This issue is quite important because the wavelront cannot. be. considered. Lat within aperture of the diameter D. ancl corresponding image in focal plane is neither dilfractive nor even axisvmmoetrical., This issue is quite important because the wavefront cannot be considered flat within aperture of the diameter $D$ and corresponding image in focal plane is neither diffractive nor even axisymmetrical.
 Normally two models are considered: gravity centre of the image (g-tilt.see7) or position corresponding to the normal of plane approximating wavefront within aperture (z-tilt.see7)..," Normally two models are considered: gravity centre of the image \citep[g-tilt, see][]{Martin1987} or position corresponding to the normal of plane approximating wavefront within aperture \citep[z-tilt, see][]{Fried1975}."
 The dillerence between these approaches was analyzed in the paper of ?.., The difference between these approaches was analyzed in the paper of \citet{Toko2002}.
" Both filters are axisvmmetrical and can be expressed through Bessel functions J,, as We will use A""(Df) by default since an exact form of aperture filter is not eritical for analysis.", Both filters are axisymmetrical and can be expressed through Bessel functions $J_n$ as We will use $A^g(Df)$ by default since an exact form of aperture filter is not critical for analysis.
 The dillerential filter corresponding to two identical apertures placed. at the distance 2 from each other along 4A axis is given by the expression: The final spectral power density £(f.ó) of the dillerential image motion measured with DIMAM is defined bv the relation: lt is worth noting that after substitution of (1)) and 2)). the spectral density. will not depend on wavelength A of incoming light.," The differential filter corresponding to two identical apertures placed at the distance $B$ from each other along $x$ axis is given by the expression: The final spectral power density $F(f,\phi)$ of the differential image motion measured with DIMM is defined by the relation: It is worth noting that after substitution of \ref{eq:KolmPSD}) ) and \ref{eq:grad}) ), the spectral density will not depend on wavelength $\lambda$ of incoming light."
" Variance of the dillerential image motion 67, can be computed. as integral. of the fuf.o) over all spatial frequencies."," Variance of the differential image motion $\sigma_{l,t}^2$ can be computed as integral of the $F_{l,t}(f,\phi)$ over all spatial frequencies."
 Because integration is being carried out in polar coordinates. factor f has to be added.," Because integration is being carried out in polar coordinates, factor $f$ has to be added."
 Let us group the spectral filters in order to separate terms depending: aud not depending on polar angle ὦ for consequent averaging by this angle., Let us group the spectral filters in order to separate terms depending and not depending on polar angle $\phi$ for consequent averaging by this angle.
 Also. it is convenientto use the dimensionless frequency dq=fD anc the dimensionless baseline b— B/D. In this case. the variance of the differential image motion is where the factor Vr;(q) is a result of averaging of all terms depending on angle © over this angle.," Also, it is convenientto use the dimensionless frequency $q=fD$ and the dimensionless baseline $b=B/D$ In this case, the variance of the differential image motion is where the factor $\Psi_{l,t}(q)$ is a result of averaging of all terms depending on angle $\phi$ over this angle."
 In case of the I-motion. it becomes and for the f-motion:," In case of the $l$ -motion, it becomes and for the $t$ -motion:"
attempt to measure morphological parameters of a large sample of LAEs as previous work has concentrated primarily on measuring sizes (2?).. used small (77S12) sample sizes (??7).. or measured the profiles of a stacked sample due to low signal-to-noise (7)..,"attempt to measure morphological parameters of a large sample of LAEs as previous work has concentrated primarily on measuring sizes \citep{Venemans05,Bond09}, used small $n \lesssim 12$ ) sample sizes \citep{Overzier08,Pirzkal07}, or measured the profiles of a stacked sample due to low signal-to-noise \citep{taniguchi}."
 The advent ol larger samples of LAEs from redshifts ranging trom z=2.1 to z>7 should allow or meaningful comparisons across redshifts down (o similar huminositw limits to be mace., The advent of larger samples of LAEs from redshifts ranging from $z=2.1$ to $z>7$ should allow for meaningful comparisons across redshifts down to similar luminosity limits to be made.
 Without such data. it is impossible to form a clear picture of how the process of galaxy omnation proceeds over (ime. (e.g..2)..," Without such data, it is impossible to form a clear picture of how the process of galaxy formation proceeds over time. \citep[e.g.,][]{overzier10}."
A wealth of observations lave shown that ealaxyv merecrs are common and that nearly all galaxies host a central supermassive black hole (SMDBITD.,A wealth of observations have shown that galaxy mergers are common and that nearly all galaxies host a central supermassive black hole (SMBH).
 Cousequeutly. sole galaxies nist host two SMDIIS as the result of recent mergers.," Consequently, some galaxies must host two SMBHs as the result of recent mergers."
 These are known as dual SMDIIs for he first ~100 Myr after the merger when they are at separationsZ1 kpc (Beechnanetal.1950:Milosavljevié&Meritt 2001).," These are known as dual SMBHs for the first $\sim 100$ Myr after the merger when they are at separations$\simgt 1$ kpc \citep{BE80.1,MI01.1}."
. Dual-SMDIT systeuis are an important esting ground. for theories of galaxw formation and evolution., Dual-SMBH systems are an important testing ground for theories of galaxy formation and evolution.
 For example. simulations predict that quasar ecdhack in imerecrs can have extreme effects on star ormation (Springeletal.2005) aud that the core-cusp division in nuclear stellar distributions may be caused * the scouring effects of dual SMDIIS (Milosavljevicetal.2002:Lauer 2007).," For example, simulations predict that quasar feedback in mergers can have extreme effects on star formation \citep{SP05.1} and that the core-cusp division in nuclear stellar distributions may be caused by the scouring effects of dual SMBHs \citep{MI02.2,LA07.1}."
 A statistical study of dual SMDBIIs aud their host galaxies would thus have iuportant duplications for theories of galaxy formation and SMDIT growth., A statistical study of dual SMBHs and their host galaxies would thus have important implications for theories of galaxy formation and SMBH growth.
 Dual SABES are observable when sufficient gas accretes. onto. them to power dual active galactic nuclei (ACN)., Dual SMBHs are observable when sufficient gas accretes onto them to power dual active galactic nuclei (AGN).
" While there have been identifications of lnudveds of binary quasar pairs at separations 210 kpe (οιο,, Tennawietal.2006:λίπους2007.2008:Ctireenetal. 20103). as well as a handful of galaxy pairs where each galaxy hosts au AGN (Balloctal.2001:Cuainazzietal.2005:Piconcelli 2010)... very few ACN pairs have heen observed iu the next evolutionary stage where they coexist at kpc-scale separations im the same mereor-remnant galaxy."," While there have been identifications of hundreds of binary quasar pairs at separations $>10$ kpc (e.g., \citealt{HE06.1,MY07.1,MY08.1,GR10.2}) ), as well as a handful of galaxy pairs where each galaxy hosts an AGN \citep{BA04.2,GU05.1,PI10.1}, very few AGN pairs have been observed in the next evolutionary stage where they coexist at kpc-scale separations in the same merger-remnant galaxy."
 To date. the confirmed dual ACN were identified bv radio or N-ray resolution of two ACN with separations of 7 kpe in the :=0.02 double radio source 3C 75 at the couter of the galaxy cluster Abel 100 (ITudsonetal. 2006).. {Προ αι the +=0.05 ultraluniuous infrared galaxy ας 163 (Bianchietal.2008).. and 0.7 kpc in the 2=0.02 ultralininous infrared galaxy NGC 62[0 (Ixo1iossaetal.2003)..," To date, the confirmed dual AGN were identified by radio or X-ray resolution of two AGN with separations of 7 kpc in the $z=0.02$ double radio source 3C 75 at the center of the galaxy cluster Abell 400 \citep{HU06.1}, , 4 kpc in the $z=0.05$ ultraluminous infrared galaxy Mrk 463 \citep{BI08.1}, and 0.7 kpc in the $z=0.02$ ultraluminous infrared galaxy NGC 6240 \citep{KO03.1}."
 Iun recent vears dual AGN candidates have Όσοι selected as galaxies with double-peaked AGN eniissiou lines. first iu the DEEP? Calaxy Redshift Survey (Corkeetal.2007:Comerford20098). and later in the Sloan Digital Sky Survey (SDSS: Wanectal.2009:Liuetal.2010b:Sith 2010)).," In recent years dual AGN candidates have been selected as galaxies with double-peaked AGN emission lines, first in the DEEP2 Galaxy Redshift Survey \citep{GE07.2,CO09.1} and later in the Sloan Digital Sky Survey (SDSS; \citealt{WA09.1, LI10.1, SM10.1}) )."
" Although: a double-peaked line profile is expected for dual AGN, it could also be produced by eas kinematics in a single AGN: follow-up observations are necessary to distinguish between the two scenarios."," Although a double-peaked line profile is expected for dual AGN, it could also be produced by gas kinematics in a single AGN; follow-up observations are necessary to distinguish between the two scenarios."
 Follow-up optical lonegslt spectroscopy. near-infrared miaeme. aud adaptive optics unaeing have added eirciuustautial evidence that many double-peaked svstenus may iudeed be dual ACN (Fuetal.2010:Liuetal. 2011).. but direct resolution of two separate AGN is required for direct evidence of dual ACN.," Follow-up optical longslit spectroscopy, near-infrared imaging, and adaptive optics imaging have added circumstantial evidence that many double-peaked systems may indeed be dual AGN \citep{FU10.1, LI10.2, SH10.1, GR11.1, RO11.1}, but direct resolution of two separate AGN is required for direct evidence of dual AGN."
 Tere we present such evidence for dual ACN in the formu of Chandra/ACIS observations of two ACN sources separated by 1.9 hey kpc in the 2=0.1569 ealaxy600835., Here we present such evidence for dual AGN in the form of /ACIS observations of two AGN sources separated by 1.9 $h^{-1}_{70}$ kpc in the $z=0.1569$ galaxy.
7.. This ealaxy was first identified as a dual AGN caucidate by its double-peaked Type 2 AGN spectrin in SDSS. where the peaks are separated by 350 kaa | (Linet[Oal.HL]20105:Smithetal. 2010).," This galaxy was first identified as a dual AGN candidate by its double-peaked Type 2 AGN spectrum in SDSS, where the peaks are separated by 350 km $^{-1}$ \citep{LI10.1, SM10.1}."
. Our follow-up Lick/Ixast lougslit spectra of the galaxy show two cinission components separated on the skyby 1.9 hey kpc. or 60δ. an angular separation resolvable by Chandra.," Our follow-up Lick/Kast longslit spectra of the galaxy show two emission components separated on the skyby 1.9 $h^{-1}_{70}$ kpc, or $\farcs$ 68, an angular separation resolvable by ."
 The follow-up Chandra observatious reveal double X-ray. sources with the same spatial separation aud orientation ou the sky, The follow-up observations reveal double X-ray sources with the same spatial separation and orientation on the sky
 107 (Arp1997.1998.1999:Bell2002a.b.c.d.2004.2006;andMeDiarmid2006.andGutiérrez.2006).," $10^{8}$ \citep{arp97,arp98,arp99,bel02a,bel02b,bel02c,bel02d,bel04,bel06,bel06a,bel07,bur99,gal05,lop06}."
". z; (Dell2004).. ~500 (Tillt1996.1997:BellanclComeanDell.ComeanandRussell2004).. IL, | | "," $_{i}$ \citep{bel04}, $\sim500$ \citep{tif96,tif97,bel03,bel04a}, $_{o}$ $^{-1}$ $^{-1}$ "
and pressure ionization effects are described in Kawka&Vennes (2006).. but further improvements to the models will be described in a forthcoming publication.,"and pressure ionization effects are described in \citet{kaw2006}, but further improvements to the models will be described in a forthcoming publication."
 All relevant species (H. . Hs. H;. Hy) are included in the statistical equilibrium equation.," All relevant species (H, $^+$ , $_2$, $_2^+$, $_3^+$ ) are included in the statistical equilibrium equation."
 We employed the Hy partition function of Neale&Tennyson (1905)., We employed the $_3^+$ partition function of \citet{nea1995}.
 Electrons contributed by identifiable trace elements (e.g.. calcium) are also included in the charge conservation equation. although the ionization of hydrogen atoms and molecules dominate the electron budget.," Electrons contributed by identifiable trace elements (e.g., calcium) are also included in the charge conservation equation, although the ionization of hydrogen atoms and molecules dominate the electron budget."
" The model atmospheres include opacities caused by H bound-bound. bound-free and free-free transitions. H bound-free and free-free transitions. H5—H» collision-induced absorptio (CIA.Borysowetal..1997,2001).. and the H»-H ana H-H collision-induced absorptions in the far Ένα wing (seeKowalski&Saumon.2006) using opacity tables from Rohrmannetal.(2011)."," The model atmospheres include opacities caused by H bound-bound, bound-free and free-free transitions, $^-$ bound-free and free-free transitions, $_2-$ $_2$ collision-induced absorption \citep[CIA,][]{bor1997,bor2001}, and the $_2$ -H and H-H collision-induced absorptions in the far $\alpha$ wing \citep[see][]{kow2006} using opacity tables from \citet{roh2011}."
.. Finally. the H» and H Rayleigh scattering are included along with electron scattering.," Finally, the $_2$ and H Rayleigh scattering are included along with electron scattering."
 Synthetic colours as well as detailed hydrogen and heavy element line profiles are computed using the model structures., Synthetic colours as well as detailed hydrogen and heavy element line profiles are computed using the model structures.
 Table 3. lists some photometric properties of the cool models., Table \ref{tbl-model} lists some photometric properties of the cool models.
 The colour indices at shorter wavelengths are effected by the Lye collision-induced absorptions., The colour indices at shorter wavelengths are effected by the $\alpha$ collision-induced absorptions.
 Allardetal.(2008) calculated new He line profiles including self-broadening effects and found a half-width at half maximum (HWHM) z40% larger than estimated by Ali&Griem(1965.1966) and comparable to the calculations of Barklemetal.(2000a).," \citet{all2008} calculated new $\alpha$ line profiles including self-broadening effects and found a half-width at half maximum (HWHM) $\approx40$ larger than estimated by \citet{ali1965,ali1966} and comparable to the calculations of \citet{bar2000a}."
. We adopted the Ha HWHM from Allardetal.(2008) and the Hf. and Hy cross-sections and velocity parameters. converted into HWHM. from Barklemetal.(2000a).," We adopted the $\alpha$ HWHM from \citet{all2008} and the $\beta$, and $\gamma$ cross-sections and velocity parameters, converted into HWHM, from \citet{bar2000a}."
. We describe the heavy element line profiles in the following section., We describe the heavy element line profiles in the following section.
 The dominant broadening mechanism ts collision with hydrogen atoms., The dominant broadening mechanism is collision with hydrogen atoms.
 We employed the coefhcients of(2000b).. where the full-width at half-maximum (FWHM) of the Lorentzian profiles is given by where log=—7.562 for Cavt4226 and -7.76 for H&KK at 7=10000 K. and aw=0.238 for Cart4226 and 0.223 for H&KK. Although the adopted broadening parameters do not include the etfect of hydrogen molecules. these provide 25066 of the gas pressure in some layers.," We employed the coefficients of, where the full-width at half-maximum (FWHM) of the Lorentzian profiles is given by where $\log{\Gamma} = -7.562$ for $\lambda4226$ and $-7.76$ for K at $T=10\,000$ K, and $\alpha = 0.238$ for $\lambda4226$ and $0.223$ for K. Although the adopted broadening parameters do not include the effect of hydrogen molecules, these provide $>$ of the gas pressure in some layers."
 Following the approximate treatment of Kurucz&Avrett(1981).. the broadening parameter Γοng0.85n(H»). and. therefore. hydrogen molecules may contribute to the total line width.," Following the approximate treatment of \citet{kur1981}, the broadening parameter $\Gamma\propto n_{\rm H}+0.85\,n({\rm H}_2)$, and, therefore, hydrogen molecules may contribute to the total line width."
 We found that neither or equivalent widths are significantly affected in models at Tay=5400 KK. but abundances inferred from Cart4226 may be underestimated by a factor of ~2 in models at KK. Fundamentals of stellar line formation in the presence of a magnetic field are described by Unno(1956) and Martin&Wickramasinghe (1981)... who also describe significant magneto-optical effects.," We found that neither or equivalent widths are significantly affected in models at $T_{\rm eff}=5\,400$ K, but abundances inferred from $\lambda4226$ may be underestimated by a factor of $\sim 2$ in models at K. Fundamentals of stellar line formation in the presence of a magnetic field are described by \citet{unn1956} and \citet{mar1981}, who also describe significant magneto-optical effects."
 In particular. Unno(1956) showed that the effect of field inclination with respect to to the line-of-sight on the relative intensity of c and π components reaches a maximum at an angle of 55°.," In particular, \citet{unn1956} showed that the effect of field inclination with respect to to the line-of-sight on the relative intensity of $\sigma$ and $\pi$ components reaches a maximum at an angle of $^\circ$."
 Moreover. Martin&Wickra-masinghe(1981) showed that taking into account magneto- effects may enhance the depth of the 7 components upon certain conditions.," Moreover, \citet{mar1981} showed that taking into account magneto-optical effects may enhance the depth of the $\pi$ components upon certain conditions."
 Kemic(1975) studied the quadratic Zeeman effect for the H and K lines and showed that the linear Zeeman effect at fields of 15 MG is still dominant., \citet{kem1975} studied the quadratic Zeeman effect for the H and K lines and showed that the linear Zeeman effect at fields of 15 MG is still dominant.
 The H and K lines are the result of transitions between the ground state with AL.S=5.0.1. and the excited states withALS=1.1.1 and 3.1.1. respectively. where J is the total angular momentum. L is the orbital angular momentum and S is the spin angular momentum.," The H and K lines are the result of transitions between the ground state with $J,L,S = \frac{1}{2},0,\frac{1}{2}$, and the excited states with$J,L,S = \frac{1}{2},1,\frac{1}{2}$ and $\frac{3}{2},1,\frac{1}{2}$, respectively, where $J$ is the total angular momentum, $L$ is the orbital angular momentum and $S$ is the spin angular momentum."
 The Lt4226 line ts the result of transitions between the ground state with and the excited state with J.L.S=1.1.0.," The $\lambda4226$ line is the result of transitions between the ground state with $J,L,S = 0,0,0$ and the excited state with $J,L,S = 1,1,0$."
" The levels are split by a magnetic field into 2/+| components defined by the magnetic quantum number oo uu where tis the wavelength inA.. B is the magnetic field inMG. e ts the electron charge. 7, is the electron rest mass and c 1s the speed of light."," The levels are split by a magnetic field into $2J+1$ components defined by the magnetic quantum number $m=-J,...,J$ : where $\lambda$ is the wavelength in, $B$ is the magnetic field inMG, $e$ is the electron charge, $m_e$ is the electron rest mass and $c$ is the speed of light."
" The Landé factor and the magnetic quantum number of the upper and lower levels are given by gy.H1, and ej. my. respectively."," The Landé factor and the magnetic quantum number of the upper and lower levels are given by $g_u, m_u$ and $g_l, m_l$ , respectively."
" Landé factors for caleium and other elements except iron were calculated assuming LS coupling: Forthe 4S,;» stateof+1) Call. the experimentally determined g¢ factor of 2.00225664 (Tommaseoetal.2003) agrees with the theoretically calculated factor based on the LS"," Landé factors for calcium and other elements except iron were calculated assuming LS coupling: Forthe $_{1/2}$ stateof , the experimentally determined $g$ factor of 2.00225664 \citep{tom2003} agrees with the theoretically calculated factor based on the LS"
velocities here agree with the positions aud velocities of the masers detected in Η. Iu II two mascrs detected in archival data were included for completeness. even though they were not detected in the observations: ll and #119.,"velocities here agree with the positions and velocities of the masers detected in I. In I two masers detected in archival data were included for completeness, even though they were not detected in the observations: 4 and 19."
 These two masers were uot re-detected in our new observations and we conclude they are extinct., These two masers were not re-detected in our new observations and we conclude they are extinct.
 Overall our results indicate stable conditions over at least vear-long terii for the masers., Overall our results indicate stable conditions over at least year-long term for the masers.
 For the relatively short timescales aud sparse sapling of maser amplitudes reported on here. it is dificult to find patterus such as periodicity. a steady flux increase or decrease.," For the relatively short timescales and sparse sampling of maser amplitudes reported on here, it is difficult to find patterns such as periodicity, a steady flux increase or decrease."
 Future more densely sampled survevs of the maser amplitudes over timescales of vears will be needed to address the question of whether variability patterus exist., Future more densely sampled surveys of the maser amplitudes over timescales of years will be needed to address the question of whether variability patterns exist.
 In this paper we therefore discuss these masers iu terms of showing variability or not., In this paper we therefore discuss these masers in terms of showing variability or not.
 By a visual iuspectiou of vefplots and the VI in Table 2.. we find that 10 out of 26 mmasers show variability with V721.0.," By a visual inspection of \\ref{plots} and the $VI$ in Table \ref{fluxes}, we find that 10 out of 26 masers show variability with $VI>1.0$."
" No maser is showing extreme variability. with the highest observed VE=2.56,"," No maser is showing extreme variability, with the highest observed $VI=2.56$."
 Alavbe except for the weakest sources in II (4111. #220 aud #222) we do not see a correlation between high V7 aud weak flux. where measured fiux errors mi be a large fraction of the measured flux.," Maybe except for the weakest sources in I 11, 20 and 22) we do not see a correlation between high $VI$ and weak flux, where measured flux errors may be a large fraction of the measured flux."
 We also do not see a correlation between flux aud array configuration. where larger svuthesized beams (at later dates) would be able to pick up more fiux if the mascrs were aneularly exteuded.," We also do not see a correlation between flux and array configuration, where larger synthesized beams (at later dates) would be able to pick up more flux if the masers were angularly extended."
 There may be a bias in the Hux if the measurements at individual epochs have used different uuubers of channels., There may be a bias in the flux if the measurements at individual epochs have used different numbers of channels.
 The effect is not eutirelv clear. aud one could argue that varving laser eniüssion. apart from varving in flux. may also vary in width of he spectral feature.," The effect is not entirely clear, and one could argue that varying maser emission, apart from varying in flux, may also vary in width of the spectral feature."
 The results obtained below secun o indicate that this poteutial dependence would oly cause a minor adjustinent to the VZ umumbers. but not he treud.," The results obtained below seem to indicate that this potential dependence would only cause a minor adjustment to the $VI$ numbers, but not the trend."
 The inost significant ποια is that the region of SNR/ISM mascrs to the northeast has a much larger raction of variable iuasers than in other reeious., The most significant trend is that the region of SNR/ISM masers to the northeast has a much larger fraction of variable masers than in other regions.
 Iu he northeast. two out of 10 imasers (s0%)) have VIo>LO with a median of V7=1.77.," In the northeast, two out of 10 masers ) have $VI>1.0$ with a median of $VI=1.77$."
 In coutrast he southeru SNR/ISM interaction. region. the three scattered nortlivestern iiasers. aud the CND masers have wo out of 15 masers (1354)) that are variable with a nedian VE=1.9L while the majority of 13 imascrs )) are nou-variable with a median VZ of0.10.," In contrast the southern SNR/ISM interaction region, the three scattered northwestern masers and the CND masers have two out of 15 masers ) that are variable with a median $VI=1.94$, while the majority of 13 masers ) are non-variable with a median $VI$ of 0.40."
 Fieure 3 ddsplavs the position of cach mascr feature. overlaid ou the 1.2GIIz coutiuuuu fux density image.," Figure \ref{maserpos}
 displays the position of each maser feature, overlaid on the GHz continuum flux density image."
 With this spatial Gdistribution iu mind. the origi of the maser variabilitv is discussed in rofdiscussion..," With this spatial distribution in mind, the origin of the maser variability is discussed in \\ref{discussion}."
 During the 195 davs over which our study took place the most significant result is that the uortheastern mascrs associated with the SNR/ISM interaction show a higher deeree of variabilitv than the rest of the mascrs., During the 195 days over which our study took place the most significant result is that the northeastern masers associated with the SNR/ISM interaction show a higher degree of variability than the rest of the masers.
 Over, Over
(Masonetal.2001).,\citep{mason01}.
. These observations achieved: tvpical limiting magnitudes for a 5e detection of 19.5 (V). 20.7 (D). 20.1 (UJ. 20.1 (UVWI) and 19.4 (UVM2).," These observations achieved typical limiting magnitudes for a $5\sigma$ detection of 19.5 (V), 20.7 (B), 20.1 (U), 20.1 (UVW1) and 19.4 (UVM2)."
 Regions of the images near to NGC 6810 are marrecl by elliptical stray light. features caused by internal rellections within the OM tclescope., Regions of the images near to NGC 6810 are marred by elliptical stray light features caused by internal reflections within the OM telescope.
 Copies of the H-band ancl continuumesubtracted images of NGC 6810. presented. in. Llamecd&Dev-(1999) were obtained [rom the NASA Extragalactic Database., Copies of the R-band and continuum-subtracted images of NGC 6810 presented in \citet{hameed99} were obtained from the NASA Extragalactic Database.
" Astrometric solutions for these images were calculated based. on the V-band OAL images. which match up with Digital Sky Survey images to within X1""."," Astrometric solutions for these images were calculated based on the V-band OM images, which match up with Digital Sky Survey images to within $\la 1 \arcsec$."
" Pointing olfsets during the OM observations were typically also 1"".", Pointing offsets during the OM observations were typically also $\la 1 \arcsec$.
 NGC 6810 is not bright enough an X-ray source to provide useful X-ray spectra with the grating spectrometers on (tbe. ROS)., NGC 6810 is not bright enough an X-ray source to provide useful X-ray spectra with the grating spectrometers on (the RGS).
 The ROS data.are. not discussed further., The RGS dataare not discussed further.
 Images in=O0.3 Ld. ld 28 and 2.8 . 8.0 keV energy bands were created [rom the EPIC PN and MOS data.," Images in $E=0.3$ – 1.1, 1.1 – 2.8 and 2.8 – 8.0 keV energy bands were created from the EPIC PN and MOS data."
 Phese enerey bands were chosen based on the X-ray. spectrum of NGC 6810. discussed. below., These energy bands were chosen based on the X-ray spectrum of NGC 6810 discussed below.
 The images from cach MOS detector. were combined. for the purposes of further image analysis., The images from each MOS detector were combined for the purposes of further image analysis.
" In each energy band the background intensity was estimated as the mean surface brightness on the same ος1) chip as NGC 6810. after excluding the data within a radius of 30""J of any detected point-like N-ray source or within 90"" of NGC 6810."," In each energy band the background intensity was estimated as the mean surface brightness on the same CCD chip as NGC 6810, after excluding the data within a radius of $30\arcsec$ of any detected point-like X-ray source or within $90\arcsec$ of NGC 6810."
 A simple uniform-level background. subtraction is appropriate given the relative compactness and brightness of the X-ray emission from NGC 6810., A simple uniform-level background subtraction is appropriate given the relative compactness and brightness of the X-ray emission from NGC 6810.
 Figure., Figure.
" 1 shows adaptively smoothed PN and combined ΑΟ images of NGC 6810 overlaid with surface brightness iso-contours [rom images smoothed with a £M44A=15"" Gaussian mask.", \ref{fig:ximages} shows adaptively smoothed PN and combined MOS images of NGC 6810 overlaid with surface brightness iso-contours from images smoothed with a $FWHM=15\arcsec$ Gaussian mask.
 Displaving the PN and combined MOS images separately is helpful when studying faint. extended emission in that it allows us to ascertain whether certain features of interest are genuine and which may be statistical or smoothing artifacts., Displaying the PN and combined MOS images separately is helpful when studying faint extended emission in that it allows us to ascertain whether certain features of interest are genuine and which may be statistical or smoothing artifacts.
" Phe lowest contour levels displayed in these images are between 3. 4.70 above the background. given the AWAD=15"" uniform smoothing (seeLardcas-tle 2000)."," The lowest contour levels displayed in these images are between 3 – $\sigma$ above the background, given the $FWHM=15\arcsec$ uniform smoothing \citep[see][]{hardcastle00}."
" In these smoothed images the soft extended X-ray emission can be traced out to a projected angular separation ols~60 T0"" (—8S. Okpe) from the plane of NGC 6810.", In these smoothed images the soft extended X-ray emission can be traced out to a projected angular separation of $z \sim 60$ – $70\arcsec$ $\sim 8$ – 9 kpc) from the plane of NGC 6810.
 The emission appears to arise from within the optical disk interior to a radius of ~50” (6.5 kpe) from the center of the ealaxy., The emission appears to arise from within the optical disk interior to a radius of $\sim 50\arcsec$ (6.5 kpc) from the center of the galaxy.
 The centroid of the soft. X-ray emission is slightly olfset to the north west of the center of the galaxy., The centroid of the soft X-ray emission is slightly offset to the north west of the center of the galaxy.
 In the hard X-ray band (47=2.8) 8.0 keV) a single point-like region of emission. roughly coincident with the center of NGC 6810. is visible.," In the hard X-ray band $E=2.8$ – 8.0 keV) a single point-like region of emission, roughly coincident with the center of NGC 6810, is visible."
 Smoothine will artificially increase the apparent extent of the X-ray emission., Smoothing will artificially increase the apparent extent of the X-ray emission.
 To better quantify the projected. size of the soft N-ray emission 1-dimoensional profiles of the X-ray surface brightness along the major and minor axes of NGC GSIO were created (see Fig. 2)).," To better quantify the projected size of the soft X-ray emission 1-dimensional profiles of the X-ray surface brightness along the major and minor axes of NGC 6810 were created (see Fig. \ref{fig:profiles}) ),"
 excluding the emission fron X-ray sources not obviously. associated with NGC 6810., excluding the emission from X-ray sources not obviously associated with NGC 6810.
 The profiles clisplay the total surface soft) X-ray xightness. the surface brightness of the expected vackeround (both the genuine X-ray and particle vackeround) and the background-subtracted. emission associated with NGC 6810.," The profiles display the total surface soft X-ray brightness, the surface brightness of the expected background (both the genuine X-ray and particle background) and the background-subtracted emission associated with NGC 6810."
 X version of the SAS processing askEDETECT_CUAIN. altered to correctly treat the presence of dilfuse emission in and around NGC 6810. was used Oo detect X-ray sources. and create maps of the total rvackeround ancl exposure time over the entire field of view of cach of the MOSI. MOS2 and PN instruments.," A version of the SAS processing task, altered to correctly treat the presence of diffuse emission in and around NGC 6810, was used to detect X-ray sources and create maps of the total background and exposure time over the entire field of view of each of the MOS1, MOS2 and PN instruments."
 The surface brightness profiles were creating usine this ocation-specilic background and exposure time information. excluding statistically significant X-ray sources (other than GC 0510 itself). and. deteetor-specific chip-gaps. bad CCD pixels and bad. columns.," The surface brightness profiles were creating using this location-specific background and exposure time information, excluding statistically significant X-ray sources (other than NGC 6810 itself) and detector-specific chip-gaps, bad CCD pixels and bad columns."
 The values of the MOS ancl PN backerouncs produced: by this method. is within 1% of the average local background: values used for the images displaved in Fig., The values of the MOS and PN backgrounds produced by this method is within $\sim 1$ of the average local background values used for the images displayed in Fig.
 1 ancl the X-ray spectra discussed in 34..., \ref{fig:ximages} and the X-ray spectra discussed in \ref{sec:results:spectral}.
" Each bin in the surface brightness profiles was initially the sum of the data in a region of width 275 parallel o the axis of interest and. total length. 76"". perpendicular ο that axis.", Each bin in the surface brightness profiles was initially the sum of the data in a region of width $2\farcs5$ parallel to the axis of interest and total length $76\arcsec$ perpendicular to that axis.
 Subsequently. consecutive bins were rebinned with the aim of achieving a minimum of 20 counts per final in., Subsequently consecutive bins were rebinned with the aim of achieving a minimum of 20 counts per final bin.
 1n an ellort to maximize the signal-to-noise for the aintest emission the data from the. PN ancl both MOS detectors was combined. and it is these combined. profiles hat are shown in the top panels of Fig. 2..," In an effort to maximize the signal-to-noise for the faintest emission the data from the PN and both MOS detectors was combined, and it is these combined profiles that are shown in the top panels of Fig. \ref{fig:profiles}."
 In general the surface brightness profiles produced from the PN. MOSI and MOS2 detectors individually were very similar to each other (with the PN ancl MOS2 profiles being the most similar to each other). justifving their combination.," In general the surface brightness profiles produced from the PN, MOS1 and MOS2 detectors individually were very similar to each other (with the PN and MOS2 profiles being the most similar to each other), justifying their combination."
 The profiles from he MOSI detector were the most dillerent. and are shown in he lower panels of Fig. 2..," The profiles from the MOS1 detector were the most different, and are shown in the lower panels of Fig. \ref{fig:profiles},"
 but still display the same general catures., but still display the same general features.
 The soft X-ray minor and major axis surface brightness xoliles for NGC 6810. are similar in form to the dilluse emission. profiles found. in. X-ray observations of jüghlv-inclined. starburst galaxies (Stricklandetal.2004)... although note that in theNAZALNew/on data we can not robustly remove X-ray emission. from the point. sources in NGC 6810.," The soft X-ray minor and major axis surface brightness profiles for NGC 6810 are similar in form to the diffuse emission profiles found in X-ray observations of highly-inclined starburst galaxies \citep{strickland04a}, although note that in the data we can not robustly remove X-ray emission from the point sources in NGC 6810."
 “Phe surface brightness crops rapidly with increasing distance from the nucleus of the galaxy along both minor and major axes., The surface brightness drops rapidly with increasing distance from the nucleus of the galaxy along both minor and major axes.
 The minor axis profiles also clearly show the asymmetry in the soft extended emission between the cast and west sides of the disk. presumably caused. by absorption within the disk of NGC 0510.," The minor axis profiles also clearly show the asymmetry in the soft extended emission between the east and west sides of the disk, presumably caused by absorption within the disk of NGC 6810."
 No obvious end or edge to the this extended soft. X-ray emission can be discerned [rom this data., No obvious end or edge to the this extended soft X-ray emission can be discerned from this data.
" At low signal-o-nolse the emission appears to extends out to projectec distances of [z| or |r| 10 15 kpe (~τὸ 115"") from he center of the galaxy.", At low signal-to-noise the emission appears to extends out to projected distances of $|z|$ or $|r| \sim $ 10 – 15 kpc $\sim 75$ – $115\arcsec$ ) from the center of the galaxy.
 However it is possible that. rather han genuine emission from these locations. we are insteac secing emission from the bright central regions of NGC 6810 in the very extended. but low surface brightness. wings of he PSE.," However it is possible that, rather than genuine emission from these locations, we are instead seeing emission from the bright central regions of NGC 6810 in the very extended, but low surface brightness, wings of the PSF."
" However. at intermediate clistances ron the nucleus the extended emission must predominantly »' eenuine. as the central core of the PSE is too compac (ENLIML ~ 6%. Half Encrey Width ~ 12"") to be responsible or the ~1 diameter region of bright soft. X-ray emission Fig."," However, at intermediate distances from the nucleus the extended emission must predominantly be genuine, as the central core of the PSF is too compact (FWHM $\sim 6\arcsec$ , Half Energy Width $\sim 12 \arcsec$ ) to be responsible for the $\sim 1\arcmin$ diameter region of bright soft X-ray emission Fig."
 1. or Fig. 2))., \ref{fig:ximages} or Fig. \ref{fig:profiles}) ).
 With no clear edge and the ellect. of the PSE it is, With no clear edge and the effect of the PSF it is
The second condition is hat the resolution must be sufficient to resolve the location of t1e Stagnation point on the binary axis.,The second condition is that the resolution must be sufficient to resolve the location of the stagnation point on the binary axis.
 This is increasingly demanding as 7) decreases. but the increase in compuational cost is steeyer when working in the 2D setup (see refanalytical)).," This is increasingly demanding as $\eta$ decreases, but the increase in computational cost is steeper when working in the 2D setup (see \\ref{analytical}) )."
 The last conditions relates directly to. the instabilities., The last conditions relates directly to the instabilities.
" For ij=1/320.03125. in a Se simulation box. we found that a simulation with n,=125 needs 7 levels of retinemen in order to avoid numerical damping of the instabilities."," For $\eta=1/32=0.03125$, in a $8a$ simulation box, we found that a simulation with $n_x=128$ needs 7 levels of refinement in order to avoid numerical damping of the instabilities."
 At lower resolutions we see the initial development of the TAI far from the binary but it is quickly advected out of the simulation box withou being maintained., At lower resolutions we see the initial development of the TAI far from the binary but it is quickly advected out of the simulation box without being maintained.
 The NTSI is not triggered and the tinal resul is stable., The NTSI is not triggered and the final result is stable.
 We find that the shell needs to be resolved by at least 4 computational cells on the binary axis in order to develop the NTSI., We find that the shell needs to be resolved by at least 4 computational cells on the binary axis in order to develop the NTSI.
 Resolving the shell shock structure is the stringiest constrain on the numerical resolution., Resolving the shell shock structure is the stringiest constraint on the numerical resolution.
 The thickness of the shell for the 2D adiabatic simulations given in Fig., The thickness of the shell for the 2D adiabatic simulations given in Fig.
 2. (upper left panel) can be used to estimate the numerical resolution required to achieve this for a given i., \ref{fig:shock_pos} (upper left panel) can be used to estimate the numerical resolution required to achieve this for a given $\eta$ .
 It drastically decreases for low values of à (slightly less so in 3D. which show thicker structures when η«1/32=0.03125. see ref3dstudy).," It drastically decreases for low values of $\eta$ (slightly less so in 3D, which show thicker structures when $\eta\leq 1/32=0.03125$, see \\ref{3dstudy}) )."
 The shell width is thinner in the isothermal case so the values derived from Fig., The shell width is thinner in the isothermal case so the values derived from Fig.
 2 are strict lower limits for the required resolution., \ref{fig:shock_pos} are strict lower limits for the required resolution.
 Large scale simulation of a system with low 7 and isothermal winds require high resolutions for the instabilities to develop., Large scale simulation of a system with low $\eta$ and isothermal winds require high resolutions for the instabilities to develop.
 The NTSI develops at slightly lower resolutions when the KHI is present and acts as the initial seed perturbation., The NTSI develops at slightly lower resolutions when the KHI is present and acts as the initial seed perturbation.
 For instance. with η=1/32. isothermal winds and ry.=2ex the NTSI develops with 6 levels of refinement instead of 7 in the case of equal winds.," For instance, with $\eta=1/32$, isothermal winds and $v_{1\infty}=2v_{\infty}$ the NTSI develops with 6 levels of refinement instead of 7 in the case of equal winds."
 However. it seems that the effect decreases with lower values of 7).," However, it seems that the effect decreases with lower values of $\eta$."
 The shell always needs to be resolved. if only minimally. because the NTSI involves an imbalance of momentum the thin shock layer.," The shell always needs to be resolved, if only minimally, because the NTSI involves an imbalance of momentum the thin shock layer."
 The Kelvin-Helmholtz instability in adiabatic winds is easier to model., The Kelvin-Helmholtz instability in adiabatic winds is easier to model.
 It develops even for low resolution simulations when the velocity difference between both winds is large enough., It develops even for low resolution simulations when the velocity difference between both winds is large enough.
 For jj=1/32. adiabatic winds and ος=2ex the instability develops for 4 levels of refinement.," For $\eta=1/32$, adiabatic winds and $v_{1\infty}=2v_{\infty}$ the instability develops for 4 levels of refinement."
 The study of the large scale 3D evolution of unstable colliding winds remains a tremendous computational challenge., The study of the large scale 3D evolution of unstable colliding winds remains a tremendous computational challenge.
 We have studied the morphology and the instability of colliding wind regions using numerical simulations., We have studied the morphology and the instability of colliding wind regions using numerical simulations.
 Compared to previous works. our study extends to much lower values of the wind momentum ratio. larger simulation domain and higher spatia resolution thanks to adaptive mesh refinement.," Compared to previous works, our study extends to much lower values of the wind momentum ratio, larger simulation domain and higher spatial resolution thanks to adaptive mesh refinement."
 We investigate the applicability of semi-analytical estimates or the contac discontinuity. finding that the solution of ? is the bes approximation to the asymptotic opening angle or small η.," We investigate the applicability of semi-analytical estimates for the contact discontinuity, finding that the solution of \citet{Stevens:1992on} is the best approximation to the asymptotic opening angle for small $\eta$."
 We find that the weaker wind can be entirely confined to a smal region instead of expanding freely up to infinity over some solid angle when low + colliding winds are considered in both the isothermal and adiabatic limits., We find that the weaker wind can be entirely confined to a small region instead of expanding freely up to infinity over some solid angle when low $\eta$ colliding winds are considered in both the isothermal and adiabatic limits.
 Instabilities in the colliding wine region are important because of the mixing and variability they induce., Instabilities in the colliding wind region are important because of the mixing and variability they induce.
 Resolving the shock structure isrequired to follow the development of instabilities. which imposes increasingly stringen minimal numerical requirements for smaller jj. ," Resolving the shock structure isrequired to follow the development of instabilities, which imposes increasingly stringent minimal numerical requirements for smaller $\eta$ "
notoriously unreliable X D relationship.,notoriously unreliable $\Sigma$ $D$ relationship.
 A more receut X D study places €12.810.6 at. 1023 kpe (Case Bhattacharva 1998)., A more recent $\Sigma$ $D$ study places G42.8+0.6 at $10\pm3$ kpc \nocite{cb98} (Case Bhattacharya 1998).
 Coven the considerable uncertainties associated with this technique. it would obviously be premature to rule out an association between €12.8|0.6 and either of the two neutron stars imm question.," Given the considerable uncertainties associated with this technique, it would obviously be premature to rule out an association between G42.8+0.6 and either of the two neutron stars in question."
 In this reeard. we note that the recent discovery by Vrba et al. (," In this regard, we note that the recent discovery by Vrba et al. ("
2000) ο à massive star cluster oulv 12 arcsec from SCR 1900]11 sugeests that it may have been formed iu this cluster rather than C12.8|0.6.,2000) \nocite{vhl+00} of a massive star cluster only 12 arcsec from SGR 1900+14 suggests that it may have been formed in this cluster rather than G42.8+0.6.
 Vrba et al., Vrba et al.
 estimate the star cluster to lie at 115 kpe. (, estimate the star cluster to lie at 14.5 kpc. (
b) Age estimates: The ages of auy uon-historical Supernova renmnants are strouglv coupled with their distances since absolute remnant sizes. along with assunptious about the expausion velocities are required to constrain the ages.,"b) Age estimates: The ages of any non-historical supernova remnants are strongly coupled with their distances since absolute remnant sizes, along with assumptions about the expansion velocities are required to constrain the ages."
 Hence the age of 612.8|0.6 is also subject to considerable uncertainty., Hence the age of G42.8+0.6 is also subject to considerable uncertainty.
 Vasisht et al. (, Vasisht et al. (
1991) quote au age of 104 vy but. eiven the above rauge of distance estimates. this could easily be uncertain by factors of a few.,"1994) quote an age of $10^4$ yr but, given the above range of distance estimates, this could easily be uncertain by factors of a few."
 For SCR 19001LL. the traditional asstuiptious about dipolar spin-down are thought not to apply aud the age is quite ucertain with current estimates of ~104 vr (see e.c. Ixouveliotou ct al.," For SGR 1900+14, the traditional assumptions about dipolar spin-down are thought not to apply and the age is quite uncertain with current estimates of $\sim10^4$ yr (see e.g. Kouveliotou et al."
 1999)., 1999).
 For PSR J1907|09185. he 38-kvr characteristic age is probably indicative of its true age.," For PSR J1907+0918, the 38-kyr characteristic age is probably indicative of its true age."
 This is somewhat model depeucdent since he age would be reduced if e.g. the birth spin period of PSR JL907|O9LS8 was close to its current value or even increased if the neutron star braking is less than hat expected from pure maguectic dipole braking (see e.g. Manchester Tavlor 1977)., This is somewhat model dependent since the age would be reduced if e.g. the birth spin period of PSR J1907+0918 was close to its current value or even increased if the neutron star braking is less than that expected from pure magnetic dipole braking (see e.g. Manchester Taylor 1977).
 Tn sumnnuuw. based ou cumreuth-available evidence we conclude that both jeutron stars appear to be voung enough to be considered as plausible candidates for an association with G12.8|0.6. (," In summary, based on currently-available evidence we conclude that both neutron stars appear to be young enough to be considered as plausible candidates for an association with G42.8+0.6. ("
c) Transverse speed estimates: Both neutron stars lic about 20 arcnmuün from the ceuter of G12.810.6 which iuplies a trausverse velocity of 1000Dz/f; lan s| to cary either of them to their present position.,"c) Transverse speed estimates: Both neutron stars lie about 20 arcmin from the center of G42.8+0.6 which implies a transverse velocity of $4000\,\,D_7/t_4$ km $^{-1}$ to carry either of them to their present position."
" Here Ds is the distance in units of 7 kpc aud f, is the age in units of 10! vr.", Here $D_7$ is the distance in units of 7 kpc and $t_4$ is the age in units of $10^4$ yr.
" Although the exact values of Dz aud f, are hnighlv uncertain. it is unlikely that they are such that the required velocity estimate is below 1000 ki κLl."," Although the exact values of $D_7$ and $t_4$ are highly uncertain, it is unlikely that they are such that the required velocity estimate is below 1000 km $^{-1}$."
 For either of the neutrou stars. then. the nuplied transverse velocities would place them at the far extremes of the preseutlv-observed distribution — (ILurison. Lxue Anderson 1993).," For either of the neutron stars, then, the implied transverse velocities would place them at the far extremes of the presently-observed distribution \nocite{hla93} (Harrison, Lyne Anderson 1993)."
 To ultimately test for an association between PSR J1907|0918 and €12.5|0.6. and coustrain the age of the pulsar. a proper motion measurement is required.," To ultimately test for an association between PSR J1907+0918 and G42.8+0.6, and constrain the age of the pulsar, a proper motion measurement is required."
" The predicted pulsar proper motion is 120/f, mas Í.", The predicted pulsar proper motion is $120/t_4$ mas $^{-1}$.
 Future VLBI proper motion neasureineuts of PSR J1907|0918. perhaps using Arecibo-Effelsbere-CGDT are hiehly desirable.," Future VLBI proper motion measurements of PSR J1907+0918, perhaps using Arecibo-Effelsberg-GBT are highly desirable."
" To summarize, based on the cmrenuth-available information. we couclude that the proposed association between €12.5|0.6 aud SCR 1900]LL is. contrary to frequent clans in the literature. far from) secure since there Is no reason agaiust arguing equallv stronely in favor of PSR J1907]|0918 as being the neutron star produced iu the supernova explosion rather than SCR 1900;11."," To summarize, based on the currently-available information, we conclude that the proposed association between G42.8+0.6 and SGR 1900+14 is, contrary to frequent claims in the literature, far from secure since there is no reason against arguing equally strongly in favor of PSR J1907+0918 as being the neutron star produced in the supernova explosion rather than SGR 1900+14."
 Indeed. the additional possibility that ucither of these voung neutron stars is associated with CI2.8[0.6 rendus attractive at this stage!," Indeed, the additional possibility that neither of these young neutron stars is associated with G42.8+0.6 remains attractive at this stage!"
 As noted by Cacnsler (2000). large positional offsets between neutron stars and supernova renmants are more likely a result of random Hine-of-ieht aligmments rather than eecuuimoelv associated high-velocity neutron stars.," As noted by Gaensler \nocite{gae00} (2000), large positional offsets between neutron stars and supernova remnants are more likely a result of random line-of-sight alignments rather than genuinely associated high-velocity neutron stars."
 This may well be the case here aud further observations (6.9. deeper iultióvaveleugth maps of the region aud propor-mnotion measurements) are clearly desirable to help resolve this most perplexing situation., This may well be the case here and further observations (e.g. deeper multi-wavelength maps of the region and proper-motion measurements) are clearly desirable to help resolve this most perplexing situation.
 We are grateful to Clivssa Ixouveliotou for sugeesting the radio search of SCR 19001E11 aud for releasing the source coordinates in advance of publication., We are grateful to Chryssa Kouveliotou for suggesting the radio search of SGR 1900+14 and for releasing the source coordinates in advance of publication.
 We thank Alex Wolszczan for making the PSPM available to us and Don Backer for providing access to the ABPP., We thank Alex Wolszczan for making the PSPM available to us and Don Backer for providing access to the ABPP.
 Without either of these instrmucuts. the observatious presented here would not have been possible.," Without either of these instruments, the observations presented here would not have been possible."
 We would also like to thank Feruaudo Camilo aud Inerid Stairs for helpful discussions concerning Arecibo timing observations. as well as Alaura McLaughlin. Jim Cordes. Fernando Calo. Chis Salter. Bryan Caeusler. Clirvssa IKouveliotou. Pete Woods. and the referee. Vicky [aspi for a umber of useful conuueuts on earlier versions of this mnauuscript.," We would also like to thank Fernando Camilo and Ingrid Stairs for helpful discussions concerning Arecibo timing observations, as well as Maura McLaughlin, Jim Cordes, Fernando Camilo, Chris Salter, Bryan Gaensler, Chryssa Kouveliotou, Pete Woods, and the referee, Vicky Kaspi, for a number of useful comments on earlier versions of this manuscript."
 Arecibo Observatory is run bv Cornell University under contract with the National Scicuce Foundation., Arecibo Observatory is run by Cornell University under contract with the National Science Foundation.
their signatures and a regularity in their occurrence.,their signatures and a regularity in their occurrence.
 The regular occurrence of similar slow elitches produces a periodic sawvtooth-like modulation of (he rotation frequency with a period οἱ GOO davs., The regular occurrence of similar slow glitches produces a periodic sawtooth-like modulation of the rotation frequency with a period of 600 days.
 As discussed above. the timing residuals do not show strictly. periodic evelieal changes because this pulsar possesses a high level of timing noise and the observed amplitude and the shape of the eveles depend on a polynomial model fitted.," As discussed above, the timing residuals do not show strictly periodic cyclical changes because this pulsar possesses a high level of timing noise and the observed amplitude and the shape of the cycles depend on a polynomial model fitted."
 A derived sawtooth-like modulation of the rotation frequency. shown in Figure 2((¢) bv the curve D. is a result of reconstruction of evclical changes that could take place in (he pulsars rotation lrequency if the pulsar rotation is modeled by a simple spin down model.," A derived sawtooth-like modulation of the rotation frequency, shown in Figure \ref{resid1}( (c) by the curve B, is a result of reconstruction of cyclical changes that could take place in the pulsar's rotation frequency if the pulsar rotation is modeled by a simple spin down model."
 The curve D reflects the properties of an actual process that generated slow elitches during ~ 30 vears., The curve B reflects the properties of an actual process that generated slow glitches during $\sim$ 30 years.
 This process was interrupted bv a large eliteh that occurred in 2009 November., This process was interrupted by a large glitch that occurred in 2009 November.
 The pulsar D0919--06. as well as D1822—09. clearly exhibits that the pulsars rotation frequency underwent (wo (vpes ol discontinuities that can be classified as normal and slow glitehes.," The pulsar B0919+06, as well as $-$ 09, clearly exhibits that the pulsar's rotation frequency underwent two types of discontinuities that can be classified as normal and slow glitches."
 Comparison of the rotation parameters for the pulsars showing the slow elitches is presented in Table 6 where (he pulsars are listed in order of decreasing of their age., Comparison of the rotation parameters for the pulsars showing the slow glitches is presented in Table \ref{3psr} where the pulsars are listed in order of decreasing of their age.
 The table gives (he pulsars D1950 name. the rotation parameters p. P. and P. characteristic age T=P/2P. and surface magnetic field B=3.2x10(PP)!? G. It is seen that all these pulsars are relatively old pulsars with ages greater than e10? vears.," The table gives the pulsar's B1950 name, the rotation parameters $\nu$, $\dot\nu$, and $\ddot\nu$, characteristic age $\tau = {P}/{2{\dot{P}}}$, and surface magnetic field $B=3.2\times10^{19}({P}{\dot{P}})^{1/2}$ G. It is seen that all these pulsars are relatively old pulsars with ages greater than $\sim 10^{5}$ years."
 The oldest pulsar D1642—023 has onlv the slow elitches., The oldest pulsar $-$ 03 has only the slow glitches.
 Their auplitude ancl the (ime interval between elitches strictly obey a certain law. that is. the glitch sequence in this pulsar possesses Che predicted. steady-state properties.," Their amplitude and the time interval between glitches strictly obey a certain law, that is, the glitch sequence in this pulsar possesses the predicted, steady-state properties."
 The two others. D09194-06 and D1322—09. are substantially vounger. have the hieher Irequency derivatives 7 ancl P. and the stronger magnetic fields.," The two others, B0919+06 and $-$ 09, are substantially younger, have the higher frequency derivatives $\dot\nu$ and $\ddot\nu$, and the stronger magnetic fields."
 They experience large elitches of normal signature that followed (he oscillatory process in the rotation lrequency identified with slow elitches., They experience large glitches of normal signature that followed the oscillatory process in the rotation frequency identified with slow glitches.
 A comparison of the pulsar parameters in this table indicates that a tendency (o have a sequence of slow glitehes with steady-state properties is correlated with the characteristic age of (he pulsar., A comparison of the pulsar parameters in this table indicates that a tendency to have a sequence of slow glitches with steady-state properties is correlated with the characteristic age of the pulsar.
 Probably. the number of pulsars with slow elitches in (heir rotation frequency. will considerably increase in the future.," Probably, the number of pulsars with slow glitches in their rotation frequency will considerably increase in the future."
 The candidates can be pulsars that exhibit evelical changes in the (Gming residuals., The candidates can be pulsars that exhibit cyclical changes in the timing residuals.
 Long sequences of the timing residuals were recently published for 366 pulsars observed αἱ the Jodrell Bank between 1968 and 2006 (Hobbsοἱal.2010)., Long sequences of the timing residuals were recently published for 366 pulsars observed at the Jodrell Bank between 1968 and 2006 \citep{hob10}.
.. A detailed analvsis of these data shows that the (timing residuals of some pulsars have clear evelical changes during all the period of observations., A detailed analysis of these data shows that the timing residuals of some pulsars have clear cyclical changes during all the period of observations.
 According to Hobbsetal. (2010).. the quasi-periodic structure in (he timing resiluals is clearly visible lor six pulsars: D1540—06. D1642—03. D1813—04. Dl1s26—17T. D1828—11. and B21484+63.," According to \citet{hob10}, , the quasi-periodic structure in the timing residuals is clearly visible for six pulsars: $-$ 06, $-$ 03, $-$ 04, $-$ 17, $-$ 11, and B2148+63."
" It should benoted Chat all these pulsars are old pulsars with ages varving from 1x10? to 3.5xLO"" vears.", It should benoted that all these pulsars are old pulsars with ages varying from $1\times 10^{5}$ to $3.5\times 10^{7}$ years.
" Among, these pulsars there are two pulsars D1642—03 and BI828—11 that were investigated earlier.", Among these pulsars there are two pulsars $-$ 03 and $-$ 11 that were investigated earlier.
 The results of spectral analvsis by Hobbsetal.(2010) for these two pulsars agree. with the resultsof previous papers (Stairsetal.2000:Shabanova 2001)..," The results of spectral analysis by \citet{hob10} for these two pulsars agree with the resultsof previous papers \citep{sta00,sha01}. ."
 It is known that clear periodic structure, It is known that clear periodic structure
in favor of local irracliatiou sceiarios (Mclxeegauοἱal.2000:Ciounelleet2001).. but it is uuclear wheher the radioactivity was »oduced in the solar nebula. or in an earlier phase of evolution. such asdi expaucdiug supernova envelopes. (Came‘on200110) or in the winds ejected [rom H-depletect Wol-Rayet (WR) stars (Arnouldetal.2000.,"in favor of local irradiation scenarios \citep{mck00,gou01}, but it is unclear whether the radioactivity was produced in the solar nebula, or in an earlier phase of evolution, such as in expanding supernova envelopes \citep{cam01b} or in the winds ejected from H-depleted Wolf-Rayet (WR) stars \citep{arn00}."
. These considerations lead to the conclusion that a condination of stellar nucleosvynthiesis coupled. with stellar aud/or local irradiation appears to be lie jest explanation for the known isotopic a1013alies., These considerations lead to the conclusion that a combination of stellar nucleosynthesis coupled with stellar and/or local irradiation appears to be the best explanation for the known isotopic anomalies.
 Iuceed. t1ere Is acculinuating evidence tliat lie yresence of the short-lived radioactivities in tle early solar system may rectire a fairly involvect ustory for their complete explanation (Meverall|Clayton2000).," Indeed, there is accumulating evidence that the presence of the short-lived radioactivities in the early solar system may require a fairly involved history for their complete explanation \citep{mey00}."
. The theory of the triggered origin of the solar system therefcye reimallis al altractive sceuarlo o explain the presence of at least some of the shor-lived radioactivities in the οιurly solar system., The theory of the triggered origin of the solar system therefore remains an attractive scenario to explain the presence of at least some of the short-lived radioactivities in the early solar system.
 Iu he last few years. the viability of the proposal has beeu investigated through utinerical simulations wine several different simulation methods (Boss1995:FosteraudBoss1996.1997:and 2000):: for reviews. see Vanhala(1998) aud Bossaud.Vauhala(200().," In the last few years, the viability of the proposal has been investigated through numerical simulations using several different simulation methods \citep{bos95,fos96,fos97,bos97,bos98,cam97,van98,van00}; for reviews, see \citet{vab98} and \citet{bos00}."
. The stinulatious suggest that molecular clou cores cali be triggeredMD into collapse by noderatey slow (~10-15 kin 1) Shock waves (Boss1995:al.1997:Valhalaand.Caineron 1905).. and the tiine scale of the process. 10? vr (Fosterand.Boss1996:VanhalaCameron1998).. is sulliciently short [or the racdioactivities to have survived in tle Least'ed amounts.," The simulations suggest that molecular cloud cores can be triggered into collapse by moderately slow $\sim$ 10-45 km $^{-1}$ ) shock waves \citep{bos95,fos96,cam97,van98}, and the time scale of the process, $\sim$ $^5$ yr \citep{fos96,van98}, is sufficiently short for the radioactivities to have survived in the measured amounts."
 Calculations studyiug tle lining of radioactivities ino the forming solar system have shown that shock wave material can )e Injectec into the collapsitig system when the postsliock gas cools rapidly. resulting in (nearly isotherma shocks (FosteratdBoss1997:VanhalaaudCamerou1995:Vauhala2000).," Calculations studying the mixing of radioactivities into the forming solar system have shown that shock wave material can be injected into the collapsing system when the postshock gas cools rapidly, resulting in (nearly) isothermal shocks \citep{fos97,van98,van00}."
. In tIs Case. shock wave material is iujected iuto the collapsing molecular cloud core through Rayeiel-Taylor (RT) fingers. with au elliciency of and over a time period of 700.000. vears (FosterandBoss1997:Vanhalaaud2000).," In this case, shock wave material is injected into the collapsing molecular cloud core through Rayleigh-Taylor (RT) fingers, with an efficiency of and over a time period of 700,000 years \citep{fos97,van00}."
". The ""adioactivities can be injeced into the collapsing system even if they are far behiid the leacine edge of the shock wave (BossandFoster1998).", The radioactivities can be injected into the collapsing system even if they are far behind the leading edge of the shock wave \citep{bos98}.
". It is also possible that shock wave material ca be iujeted iuto tlie system iu iou-isotherimal shock waves where the postshock gas remains hot beind he shock οί. bi itlese aspects of the prollei have not )eeu investigated further because they appear to be beyo xl«""urrent computational power (VaulhaaandCameron1908)."," It is also possible that shock wave material can be injected into the system in non-isothermal shock waves where the postshock gas remains hot behind the shock front, but these aspects of the problem have not been investigated further because they appear to be beyond current computational power \citep{van98}."
. Consequently. he cliscussion of tlje 1jection process las conceutratecd onu the isothermal case.," Consequently, the discussion of the injection process has concentrated on the isothermal case."
 Previous calculatious |ave suggested that tlie abundatces of the injected material may ave experienced spatial and tenj»oral variation in the early solar system., Previous calculations have suggested that the abundances of the injected material may have experienced spatial and temporal variation in the early solar system.
 This is based on two discoveries of tlje Interaction between he shock wave aud the molecuar cloud core., This is based on two discoveries of the interaction between the shock wave and the molecular cloud core.
 First. there appears be a lag ofa lew x 10.000 vr between the time the center of tle compressed molecular cloud core is pushed into collapse aud the time the radioactivities carried by the shock wave arrive deep iu the system. (Vaulialaaid.Boss2000).," First, there appears be a lag of a few $\times$ 10,000 yr between the time the center of the compressed molecular cloud core is pushed into collapse and the time the radioactivities carried by the shock wave arrive deep in the system \citep{van00}."
. This is due to the fact tlat while the core can be pushed into collapse by the compressional wave transmitted through the core as the shock wave is deceleratec by the outer layers of the core. injection becomes efficient ouv alter the RT-lingers have developed fully at the surface of the core.," This is due to the fact that while the core can be pushed into collapse by the compressional wave transmitted through the core as the shock wave is decelerated by the outer layers of the core, injection becomes efficient only after the RT-fingers have developed fully at the surface of the core."
 This makes it possible for he amount of radioactivities to have, This makes it possible for the amount of radioactivities to have
"and Afax to be the pulse duration (this is actually what is termed the ""effective width” of the pulse).",and ${\Delta}t_{\rm SN}$ to be the pulse duration (this is actually what is termed the “effective width” of the pulse).
 We thenfiud Δίων tobe ~[5 d iu Band ~138 d (about three times longer) in 2., We thenfind ${\Delta}t_{\rm SN}$ to be $\sim 45$ d in $B$ and $\sim 138$ d (about three times longer) in $R$.
 We can estimate the II. columu cdeusitv from the extinction toward the echo., We can estimate the H column density from the extinction toward the echo.
 We assmue that. since the echo lies quite close to the line-of-sight to SN 2003ed. the echo suffers the same amount of extinction as docs the SN.," We assume that, since the echo lies quite close to the line-of-sight to SN 2003gd, the echo suffers the same amount of extinction as does the SN."
 For the SN we estimate a totalreddening L(BV)—0.13£0.03 mae (Van Dyk ct al., For the SN we estimate a total reddening $E(B-V)=0.13 \pm 0.03$ mag (Van Dyk et al.
 2003: Smartt et al., 2003; Smartt et al.
 2001 derive a consistent estimate of the SN reddening. but with a larecr uncertainty. E(BVj=0.11£0.16 mag).," 2004 derive a consistent estimate of the SN reddening, but with a larger uncertainty, $E(B-V)=0.11 \pm 0.16$ mag)."
 Next. we inst subtract the Calactic reddening contriübution. ΓιοWV)=0.07 imag (Schlegel. Finkbeimer. Davis 1998).," Next, we must subtract the Galactic reddening contribution, $E(B-V)=0.07$ mag (Schlegel, Finkbeiner, Davis 1998)."
" Asstunine the ratio of total-to-selective extinction Ry=3.1 (eg... Cardelli. Clayton. Mathis 1989). we find Ay=0.19 mag internal to the host galaxy,"," Assuming the ratio of total-to-selective extinction $R_V = 3.1$ (e.g., Cardelli, Clayton, Mathis 1989), we find $A_V = 0.19$ mag internal to the host galaxy."
 Bohlin. Savage. Drake (1978) found a fairly coustaut enipirical relation over the diffuse interstellar muediunu iu the Galaxy. Ny=5c10E(QBOW). which provides a normalization for the extinction curve ely/Nyp=5.3«102? em (Weingartuer Draine 2001).," Bohlin, Savage, Drake (1978) found a fairly constant empirical relation over the diffuse interstellar medium in the Galaxy, $N_H=5.8
\times 10^{21} E(B-V)$, which provides a normalization for the extinction curve $A_V/N_H = 5.3 \times 10^{-22}$ $^2$ (Weingartner Draine 2001)."
" From this relation. aud including our estimated muacertaimty iu the reddening. we derive Nj,=3.5(41.7)«10?"" cu? in the dust sheet."," From this relation, and including our estimated uncertainty in the reddening, we derive $N_H=3.5 (\pm 1.7) \times 10^{20}$ $^2$ in the dust sheet."
" Tn Table 2 we present the fluxes F4, iun each baud for several echo models. for which we have varied the composition of the dust grains."," In Table 2 we present the fluxes $F_{\rm echo}$ in each band for several echo models, for which we have varied the composition of the dust grains."
 We have calculated the set of models for both our distance assuuption aud the ILleudry et al. (, We have calculated the set of models for both our distance assumption and the Hendry et al. (
2005) distance estimate.,2005) distance estimate.
 The first model iu the set is the ciffuse Galactic dust model from Woeimgartuer Draine (2001) aud Draine (2003): it assunies solar abundances with Ry=3.1 aud the total C abuudauce per II uucleon be=56 ppii. with comparable contributions of carbonaceous and silicate dust with radii in the rauge 5.0 2.0 you. We also consider models with Ry=3.1 which are cither pure silicates or pure carbonaceous eraius.," The first model in the set is the diffuse Galactic dust model from Weingartner Draine (2001) and Draine (2003); it assumes solar abundances with $R_V=3.1$ and the total C abundance per H nucleon $b_C=56$ ppm, with comparable contributions of carbonaceous and silicate dust with radii in the range 5.0 – 2.0 $\mu$ m. We also consider models with $R_V=3.1$ which are either pure silicates or pure carbonaceous grains."
 Finally. we cousider a 11odel with comparable silicate aud carbonaceous erain composition. but assmnuiue Ry=L0.," Finally, we consider a model with comparable silicate and carbonaceous grain composition, but assuming $R_V=4.0$."
 The results are also shown eraphlically in Figure 3., The results are also shown graphically in Figure 3.
 The overall agreement of the models aud the observations is remarkably eood., The overall agreement of the models and the observations is remarkably good.
 With the various assunied model inputs. we are reproducing the observed echo reasonably well. aud this further nuplies that the echo likely arises from the ciffuse iuterstellar dust near the SN.," With the various assumed model inputs, we are reproducing the observed echo reasonably well, and this further implies that the echo likely arises from the diffuse interstellar dust near the SN."
" The uucertaiuties in the model fuxes (arising mostly frou the uncertainties iu the echo geometrical mecasurcuicuts and in our reddening estimate) are rather lareec. but ο to the measurement wncertaintics in the 6served. fluxes,"," The uncertainties in the model fluxes (arising mostly from the uncertainties in the echo geometrical measurements and in our reddening estimate) are rather large, but are comparable to the measurement uncertainties in the observed fluxes."
 What we notice is that the C|Si model agrees quite well with the observations., What we notice is that the C+Si model agrees quite well with the observations.
 The value of Ay: (3.0 or L1) has little bearing on this agreement., The value of $R_V$ (3.0 or 4.1) has little bearing on this agreement.
 The pure Si-vich dust model is also consistent with the observations: the pure C-vich dust model. less so (although it agrees to within the uucertainties for d=7.2 Alpe).," The pure Si-rich dust model is also consistent with the observations; the pure C-rich dust model, less so (although it agrees to within the uncertainties for $d=7.2$ Mpc)."
 Iu fact. for the larger assumed SN distauce (d=9.3 Mpc). the pure carbonaceous dust model is no longer cousisteut with the observations at either band and can be ruled out.," In fact, for the larger assumed SN distance $d=9.3$ Mpc), the pure carbonaceous dust model is no longer consistent with the observations at either band and can be ruled out."
 We note that the remaining models calculated for the larger SN distance generally tend to underestimate the flux. althoueh takiug iuto account the large uncertainties in both the observed and model fluxes. it is impossible to rule them out eutirely.," We note that the remaining models calculated for the larger SN distance generally tend to underestimate the flux, although taking into account the large uncertainties in both the observed and model fluxes, it is impossible to rule them out entirely."
 However. we teutatively sugecst that the observed echo may indicate that the actual SN distance is closer to the simaller value we assumed in Van Dyk et al. (," However, we tentatively suggest that the observed echo may indicate that the actual SN distance is closer to the smaller value we assumed in Van Dyk et al. ("
2003) than the larger oue determined by Heudry ct al. (,2003) than the larger one determined by Hendry et al. (
2005: and the similarly larecr distance asstuned by Simartt et al.,2005; and the similarly larger distance assumed by Smartt et al.
 2001)., 2004).
 This. along with the value of Ay. has iuplicatious for the absolute magnitude. and therefore the initial mass. of the SN progenitor.," This, along with the value of $R_V$, has implications for the absolute magnitude, and therefore the initial mass, of the SN progenitor."
 A lugher Ay would inuplw that the progenitor was. at most. (kl mag more huninous than what we estimated iu Van Dyk et al.," A higher $R_V$ would imply that the progenitor was, at most, $\sim 0.1$ mag more luminous than what we estimated in Van Dyk et al."
 ITosvever. the larger distance would require the star to be 0.6 mag iore luminous. which would lucrease the mass estimate by ~1AL. (des it would nuply that the initial mass was closer to ~10AZ.).," However, the larger distance would require the star to be $\sim 0.6$ mag more luminous, which would increase the mass estimate by $\sim 1\ M_{\odot}$ (i.e., it would imply that the initial mass was closer to $\sim 10\ M_{\odot}$ )."
 The relative agreement between the observed echo aud the echo models based on the shorter distance reassures ux that our low progenitor mass estimate (~8 9 ALL). although uncomfortably near the theoretical Iit for core collapse (Woosley Weaver 1986). is realistic.," The relative agreement between the observed echo and the echo models based on the shorter distance reassures us that our low progenitor mass estimate $\sim 8$ –9 $M_{\odot}$ ), although uncomfortably near the theoretical limit for core collapse (Woosley Weaver 1986), is realistic."
 We lave confined the presence of a scattered light echo around the nearby Type ILplateaun ον 2003ed i MIFL., We have confirmed the presence of a scattered light echo around the nearby Type II-plateau SN 2003gd in M74.
 This discovery could only have been made iu inages produced with the superior augular resolution of theHST ACS/TIRC at sufficiently late times for the SN., This discovery could only have been made in images produced with the superior angular resolution of the ACS/HRC at sufficiently late times for the SN.
" We conclude that the echo arises from dust in the -iterstellay SN environment. and our modeling (withiu ορje large uncertainties in the observations. which further —xopaeate into the models) sueeests that this dust. both Be1 conipositiou and in grain size distribution. is not unlike EAust in the diffuse Galactic interstellar οςτα, although is also possible the dust could be more silicate-vich thui Arbon-oich."," We conclude that the echo arises from dust in the interstellar SN environment, and our modeling (within the large uncertainties in the observations, which further propagate into the models) suggests that this dust, both in composition and in grain size distribution, is not unlike dust in the diffuse Galactic interstellar medium, although it is also possible the dust could be more silicate-rich than carbon-rich."
 Iu fact. our echo models teud to disfavor Aust in the SN euvironnment which is more abundant in carbonaceous eras than silicates. (," In fact, our echo models tend to disfavor dust in the SN environment which is more abundant in carbonaceous grains than silicates. ("
We note that Sueenuan 2005 found that the echo may arise from siuall ο eraius.),We note that Sugerman 2005 found that the echo may arise from small carbon-rich grains.)
 The models are not particularly scusitive to the value of Ry (but we did not compute models with Ry>L)., The models are not particularly sensitive to the value of $R_V$ (but we did not compute models with $R_V > 4$ ).
 However. miodels based on the shorter distance to the SN hat we assunied in Van Dyk et al. (," However, models based on the shorter distance to the SN that we assumed in Van Dyk et al. ("
2003: 7.2 Mpc) appear o be somewhat more cousisteut with the observed echo hau those for the longer distance assunued by Simartt et al (,2003; 7.2 Mpc) appear to be somewhat more consistent with the observed echo than those for the longer distance assumed by Smartt et al. (
2001: 9.1 Mpc) aud DHeudry et al. (,2004; 9.1 Mpc) and Hendry et al. (
2005: 9.5 Apc). though the uncertainties are large.,"2005; 9.3 Mpc), though the uncertainties are large."
 These latter two actors slightly iuerease our confidence in the relatively ow estimate (8 9 M.) for the initial mass of the SN xogenitor we derived in Vau Dxk et al., These latter two factors slightly increase our confidence in the relatively low estimate $\sim 8$ –9 $M_{\odot}$ ) for the initial mass of the SN progenitor we derived in Van Dyk et al.
 Frou Nyy aud asuniug a path leneth £=AJz35 pC or the dust sheet. the IF uuuber density would be ngz1 cii.," From $N_H$ and assuming a path length $L = \Delta l \approx 35$ pc for the dust sheet, the H number density would be $n_{\rm H} \approx 7$ $^{-3}$."
 Combined with the extinction to the SN. this is consistent with the expectation that light echoes likely emerge from regions with myc10 / and dezd mae (Suecriman 2003).," Combined with the extinction to the SN, this is consistent with the expectation that light echoes likely emerge from regions with $n_{\rm H} \approx 10$ $^{-3}$ and $A_V \lesssim 1$ mag (Sugerman 2003)."
 This echo should be further monitored with LST. including use of additional bands. particularly in the UV. to far better coustrain the nature of the scattering dust aud the echo geometry. and to reveal further new or evolving structures in the echo.," This echo should be further monitored with , including use of additional bands, particularly in the UV, to far better constrain the nature of the scattering dust and the echo geometry, and to reveal further new or evolving structures in the echo."
 The work of A.V.FJs group at UC Berkeley is supported, The work of A.V.F.'s group at UC Berkeley is supported
" Here, we will not delve into the intricacies of fitting the light curves, but focus on the determination of the scale length of the exponential disk."," Here, we will not delve into the intricacies of fitting the light curves, but focus on the determination of the scale length of the exponential disk."
" Despite the long tradition (see references in section 1), the important and comprehensive study by Knapen&vanderKruit(1991) showed that the errors in these data are still significantly large (~25%), especially if they were obtained from photographic plates."," Despite the long tradition (see references in section 1), the important and comprehensive study by \citet{KvK91} showed that the errors in these data are still significantly large $\approx 25\%$ ), especially if they were obtained from photographic plates."
" The uncertainties depend on image depth, image sky coverage, data reduction, disk region fitting, the order in which bulges, bars, or other components are fitted."," The uncertainties depend on image depth, image sky coverage, data reduction, disk region fitting, the order in which bulges, bars, or other components are fitted."
" These matters become more complicated when analysing with SDSS images which are relatively shallow, and even more so when automatically fitting thousands of galaxies which cover a wide range of brightness and morphologies."," These matters become more complicated when analysing with SDSS images which are relatively shallow, and even more so when automatically fitting thousands of galaxies which cover a wide range of brightness and morphologies."
" To avoid complications that are not related to the nature of our analysis (e.g.,Fathi&Peletier2003),, we have decided to derive the disk scale length simply by fitting an exponential profile to a pre-defined disk region of each galaxy, i.e., the region where we assume the light to be dominated by the exponential profile."," To avoid complications that are not related to the nature of our analysis \citep[e.g., ][]{Fathi03}, we have decided to derive the disk scale length simply by fitting an exponential profile to a pre-defined disk region of each galaxy, i.e., the region where we assume the light to be dominated by the exponential profile."
 This means that we are simply cutting out the central regions of the galaxies where bulges and strong bars are expected., This means that we are simply cutting out the central regions of the galaxies where bulges and strong bars are expected.
 We determine the disk region by empirically fitting the equation (2)) to a set of ranges where we expect the disk to dominate the derived surface brightness profiles., We determine the disk region by empirically fitting the equation \ref{eq:exponential}) ) to a set of ranges where we expect the disk to dominate the derived surface brightness profiles.
" We use the isoA parameter to estimate this range, and randomly select 800 galaxies, to which we apply this test both in r and i-band."," We use the $isoA$ parameter to estimate this range, and randomly select 800 galaxies, to which we apply this test both in $r$ and $i$ -band."
" In Fig. 7,,"," In Fig. \ref{fig:diskrange},"
" we show the resulting disk scale lengths when fitting the regions presented in Table 2, and when normalised to our nominal of the isoA radius, we find that the derived scale lengths change by less than for a wide range of assumed disk ranges (seen as the unshaded region in the rightmost panels in Fig. 4))."," we show the resulting disk scale lengths when fitting the regions presented in Table \ref{tab:diskrange} , and when normalised to our nominal of the $isoA$ radius, we find that the derived scale lengths change by less than for a wide range of assumed disk ranges (seen as the unshaded region in the rightmost panels in Fig. \ref{fig:thecode}) )."
" We further note that the distribution of the data points for each test, normalised to the isoA range, is well represented by a Gaussian, and 2596—the 10096error bars in Fig."," We further note that the distribution of the data points for each test, normalised to the $isoA$ range, is well represented by a Gaussian, and the error bars in Fig."
 7 are indeed symmetric., \ref{fig:diskrange} are indeed symmetric.
" For a similar test, we assume sky regions at different distances from each galaxy centre and assess the effect of the sky subtraction on the derived scale lengths."," For a similar test, we assume sky regions at different distances from each galaxy centre and assess the effect of the sky subtraction on the derived scale lengths."
" Applied to the same randomly selected 800 galaxies, we found that, assuming that the sky is represented by the 2.0+0.25xisoA region, robust scale length and surface brightness measurements are delivered."," Applied to the same randomly selected 800 galaxies, we found that, assuming that the sky is represented by the $2.0\pm 0.25 \times isoA$ region, robust scale length and surface brightness measurements are delivered."
" Although the SDSS is one of the most influential and ambitious astronomical surveys, the depth of its images in all bands are not equal."," Although the SDSS is one of the most influential and ambitious astronomical surveys, the depth of its images in all bands are not equal."
" Here we have chosen to analyse only the galaxies for which SDSS provides spectroscopic redshifts (in order to investigate the redshift evolution the parameters we derived), where SDSS is complete for r-band magnitude «17.7."," Here we have chosen to analyse only the galaxies for which SDSS provides spectroscopic redshifts (in order to investigate the redshift evolution the parameters we derived), where SDSS is complete for $r$ -band magnitude $<17.7$."
" The images in other bands are not equally deep and/or complete to this magnitude limit, partly due to the significantly different transmission curves for the different filters."," The images in other bands are not equally deep and/or complete to this magnitude limit, partly due to the significantly different transmission curves for the different filters."
" Including atmospheric extinction and detector efficiency, the peak quantum efficiency of the system in u and z-bands are z1096, g and i-band 35%, and r-band =50%."," Including atmospheric extinction and detector efficiency, the peak quantum efficiency of the system in $u$ and $z$ -bands are $\approx 10\%$, $g$ and $i$ -band $\approx 35\%$ , and $r$ -band $\approx 50\%$."
" Thus it is necessary to apply a magnitude cut which varies depending on the band, fainter than which we are not able to derive reliable scale lengths."," Thus it is necessary to apply a magnitude cut which varies depending on the band, fainter than which we are not able to derive reliable scale lengths."
" For each pair of SDSS filters, we expect that the scale"," For each pair of SDSS filters, we expect that the scale"
for 2«0.10 (hevoud 2=0.10 we do uot see practically any galaxy). for auv galaxy type (obtained frou the fit of data from Mannuucci ct al.,"for $z<0.40$ (beyond $z=0.40$ we do not see practically any galaxy), for any galaxy type (obtained from the fit of data from Mannucci et al."
 2001)., 2001).
" The redshift is :(74,44) Corresponds to the velocity ¢(rpays)=1007,54IxMpe) lans Cinear DImbble law taken as an approximation for low redshift ealaxies).", The redshift is $z(r_{phys})$ corresponds to the velocity $v(r_{phys})=100r_{phys}(h^{-1}{\rm Mpc})$ km/s (linear Hubble law taken as an approximation for low redshift galaxies).
 The huuinositv function d is aken from Iochauck et al. (, The luminosity function $\Phi $ is taken from Kochanek et al. (
2001).,2001).
 We neelect the evolution of the huninosity fuuction. which is nuld (iexO44. Pozzetti et al.," We neglect the evolution of the luminosity function, which is mild $\frac{dM_K}{dz}\approx 0.5$, Pozzetti et al."
 2003) aud very siall at the mean redshift of our galaxies (62)z (L083): a shift of ~0.01 magnitudes will produce a variation of ~MA in both (6N) and CN) in the same direction: since we calculate the rate ANAL. this sniall variations mostly compensate. aud only second order variations remain. which are uceligible.," 2003) and very small at the mean redshift of our galaxies $\langle z
\rangle \approx 0.083$ ): a shift of $\sim 0.04$ magnitudes will produce a variation of $\sim 2$ in both $\langle \delta N\rangle$ and $\langle N\rangle $ in the same direction; since we calculate the rate $\frac{\delta N}{N}$, this small variations mostly compensate, and only second order variations remain, which are negligible."
 The errors associated with the uncertainty iu the Iuninositv itself do maily affect the normalization. which is also irrelevant for the calculation of ON/N since normalization factor caucels out.," The errors associated with the uncertainty in the luminosity itself do mainly affect the normalization, which is also irrelevant for the calculation of $\delta N/N$ since normalization factor cancels out."
 Finally. the two-point correlation function is taken to he (for colmoving coordinates): where ΕΤ stands for Fourier transform. and + is eiven iu units 7 !Mpe;," Finally, the two-point correlation function is taken to be (for comoving coordinates): where $FT[...]$ stands for Fourier transform, and $r$ is given in units $h^{-1}$ Mpc."
 The truucatiou of tle two-point correlation fiction is sot at 10h spe. but the exact value does not matter. our results do not depend on this nüuiunn scale: it is introduced to avoid computer algorithii problems at very πα scales.," The truncation of the two-point correlation function is set at 10 $h^{-1}$ kpc, but the exact value does not matter, our results do not depend on this minimum scale; it is introduced to avoid computer algorithm problems at very small scales."
 More exactly. the change of regie should occur at the matching poiut," More exactly, the change of regime should occur at the matching point"
The past decade has seen a rapid increase in the rate of discovery and classification of variable stars. mainly as a result of time-domain photometric surveys whose primary goals were to search for other phenomena. such as microlensing or planetary transits tthe OGLE survey. Udalskietal. 2008)).,"The past decade has seen a rapid increase in the rate of discovery and classification of variable stars, mainly as a result of time-domain photometric surveys whose primary goals were to search for other phenomena, such as microlensing or planetary transits the OGLE survey, \citealt{uda08}) )."
 Sources of stellar variability are wide ranging. from the potentially large-amplitude signatures of eclipses. star spots and pulsations. down to the sub-millimagnitude changes induced by granulation.," Sources of stellar variability are wide ranging, from the potentially large-amplitude signatures of eclipses, star spots and pulsations, down to the sub-millimagnitude changes induced by granulation."
 The typical precision and time sampling of ground-based surveys confines the associated variability studies to ‘classical’ variable stars. with amplitudes of a percent or more.," The typical precision and time sampling of ground-based surveys confines the associated variability studies to `classical' variable stars, with amplitudes of a percent or more."
" In these surveys. the Sun. whose total output never varies by more than around peak-to-peak2011). even at the maximum of its activity cycle. would appear as a ""quiet or ""constant! star."," In these surveys, the Sun, whose total output never varies by more than around peak-to-peak, even at the maximum of its activity cycle, would appear as a `quiet' or `constant' star."
 However. space-based transit surveys such as CoRoT (Baglin2003) and (Boruckietal.2010) are sensitive to micro-variability? down to anc well below the solar level. on timescales ranging from minutes to months and. over the entire lifetime ofKepler. years.," However, space-based transit surveys such as CoRoT \citep{bag03} and \citep{bor10} are sensitive to `micro-variability' down to and well below the solar level, on timescales ranging from minutes to months and, over the entire lifetime of, years."
 Measuring the basic characteristics of the variability (amplitude. periodicity. ete...)) across large samples of stars. and comparing them to stellar parameters such as age. mass and composition. Is a first step towards a better understanding of the underlying phenomena.," Measuring the basic characteristics of the variability (amplitude, periodicity, ) across large samples of stars, and comparing them to stellar parameters such as age, mass and composition, is a first step towards a better understanding of the underlying phenomena."
 Many of the latter are ill-understood. because they are related to convection and magnetism which are challenging to model.," Many of the latter are ill-understood, because they are related to convection and magnetism which are challenging to model."
 Variability statistics also have a crucial impact on exoplanet studies. particularly for radial-velocity searches or radial-velocity confirmation of transiting planet candidates (seee.g.Pontetal.2011).," Variability statistics also have a crucial impact on exoplanet studies, particularly for radial-velocity searches or radial-velocity confirmation of transiting planet candidates \citep[see e.g.][]{pon10}."
. Data from the nission is particularly amenable to statistical variability studies because of the instrument's unprecedented and vast field of viewdeg)., Data from the mission is particularly amenable to statistical variability studies because of the instrument's unprecedented and vast field of view.
 The Quarter 1 (QI) data. which was made public in June 2010. has already been studied by Basrietal.(2010.hereafter BIO)... who show that somewhat less than half of the dwarf stars surveyed by are more variable thar the Sun on timescales of up to a month. with the fractior increasing from earlier to later spectral types.," The Quarter 1 (Q1) data, which was made public in June 2010, has already been studied by \citet[][hereafter B10]{bas10}, who show that somewhat less than half of the dwarf stars surveyed by are more variable than the Sun on timescales of up to a month, with the fraction increasing from earlier to later spectral types."
 Basrietal.(2011.hereafterBIT) went on to demonstrate that periodic variable stars have significantly larger amplitudes. as a sample. thar aperiodic variables.," \citet[][hereafter B11]{bas11} went on to demonstrate that periodic variable stars have significantly larger amplitudes, as a sample, than aperiodic variables."
 Finally. Ciardietal.(2011.hereafterC11) performed a complementary study of the same sample using dispersion rather than amplitude as a variability statistic. anc studying likely dwarfs and giants separately using the stellar parameters provided in the Input Catalog (KIC. Browetal.2011:Batalha 2010)).," Finally, \citet[][hereafter C11]{cia10} performed a complementary study of the same sample using dispersion rather than amplitude as a variability statistic, and studying likely dwarfs and giants separately using the stellar parameters provided in the Input Catalog (KIC, \citealt{bro11,bat10}) )."
 Variability statistics have also been determined using the 10 days of commissioning data (QO). with the aim of developing methods to characterise anc select specific types of variable (Walkowiez&Basri2010).," Variability statistics have also been determined using the 10 days of commissioning data (Q0), with the aim of developing methods to characterise and select specific types of variable \citep{wal10}."
. In C11. the data were corrected for systematics using the team’s Pre-Search Data Conditioning (PDC) method (Jenkinsetal.2010).," In C11, the data were corrected for systematics using the team's Pre-Search Data Conditioning (PDC) method \citep{jen10}."
. This inspired us to investigate further the apparent bimodality in the variability of dwarf stars observed byKepler.. with particular attention to the effect of different systematics correction methods.," This inspired us to investigate further the apparent bimodality in the variability of dwarf stars observed by, with particular attention to the effect of different systematics correction methods."
 The goal of the present paper 1s to revisit the work of C11 and B10.11 following the application of a new astrophysically robust de-trending method. designed to preserve intrinsic variability signals. and remove as fully as possible. the systematics.," The goal of the present paper is to revisit the work of C11 and B10,11 following the application of a new astrophysically robust de-trending method, designed to preserve intrinsic variability signals and remove as fully as possible the systematics."
 We quantify trends. previously identified and new. between variability characteristics and stellar properties. and investigate the nature of the variability in more detail. 1n order to gain further insight into the underlying mechanisms.," We quantify trends, previously identified and new, between variability characteristics and stellar properties, and investigate the nature of the variability in more detail, in order to gain further insight into the underlying mechanisms."
 We describe the data and our systematics correction process in Section 2.., We describe the data and our systematics correction process in Section \ref{sec:sysindat}.
 Ir Section 3 we discuss our choice of variability statistics and examinethe fractions and physical properties of the low anc high variability samples. focussing on their periodic and stochastic nature in Section 4..," In Section \ref{sec:stats_sect} we discuss our choice of variability statistics and examinethe fractions and physical properties of the low and high variability samples, focussing on their periodic and stochastic nature in Section \ref{sec:perstoch}."
 We present our conclusions and discuss the implications of our findings in Section 5.., We present our conclusions and discuss the implications of our findings in Section \ref{sec:disc}. .
" The Quarter 1 (QI) observations took place over ~33.5 days between May 13!"" and June 15! 2009.", The Quarter 1 (Q1) observations took place over $\sim$ 33.5 days between May $^{\rm th}$ and June $^{\rm th}$ 2009.
 156.097 targets," 156,097 targets"
Deep images were obtained on 2002 February 2 using the (STIS) CCD onHST. both with open filter (SOCCD) and long-pass filter (F28X50LP).,"Deep images were obtained on 2002 February 2 using the (STIS) CCD on, both with open filter (50CCD) and long-pass filter (F28X50LP)."
 In each filter. a total exposure time of 10.400 s was obtained in a standard parallelogram dither pattern with four exposures at each point to facilitate the rejection of cosmic rays.," In each filter, a total exposure time of 10,400 s was obtained in a standard parallelogram dither pattern with four exposures at } each point to facilitate the rejection of cosmic rays."
 The data were processed using the standard STIS pipeline. registered. and combined according to the dither pattern.," The data were processed using the standard STIS pipeline, registered, and combined according to the dither pattern."
 Figure 4 shows a portion of the resulting open filter image centered onJ1$36., Figure 4 shows a portion of the resulting open filter image centered on.
2+5925.. Aperture photometry was used to obtain STIS magnitudes from the PHOTLAM and PHOTZPT keywords in the image headers., Aperture photometry was used to obtain STIS magnitudes from the PHOTLAM and PHOTZPT keywords in the image headers.
 We then adopted the transformation equations derived by Rejkuba et al. (, We then adopted the transformation equations derived by Rejkuba et al. (
2000) from observations of stars in the irregular galaxy WLM to convert STIS magnitudes in the open and long-pass filters to the V and I Kron-Cousins system.,2000) from observations of stars in the irregular galaxy WLM to convert STIS magnitudes in the open and long-pass filters to the $V$ and $I$ Kron-Cousins system.
 We also examined a set of objects in the STIS images that we had previously calibrated 1n ground-based photometry. and the zero-points agree.," We also examined a set of objects in the STIS images that we had previously calibrated in ground-based photometry, and the zero-points agree."
 Thus. we find that the limiting magnitudes for point-source detection at the 3o level are V=28.8 from the 50CCD image and and /=26.5 from the F28X50LP.," Thus, we find that the limiting magnitudes for point-source detection at the $3\sigma$ level are $V = 28.8$ from the 50CCD image and and $I = 26.5$ from the F28X50LP."
 The main uncertainty 1n this calibration is due to the broad passbands of the STIS filters. and is estimated as 0.2 mag.," The main uncertainty in this calibration is due to the broad passbands of the STIS filters, and is estimated as 0.2 mag."
 In order to tie the andChandra images to the same astrometric reference frame. we used the positions of 10 objects that are present on the both the STIS and ground-based CCD images to transfer the USNO-A2.0 reference frame to the STIS image.," In order to tie the and images to the same astrometric reference frame, we used the positions of 10 objects that are present on the both the STIS and ground-based CCD images to transfer the USNO-A2.0 reference frame to the STIS image."
 These objects include stars and compact galaxies., These objects include stars and compact galaxies.
 The dispersion among these 10 secondary astrometric standards from the fit to their positions is 07065. or slightly larger than | STIS pixel.," The dispersion among these 10 secondary astrometric standards from the fit to their positions is $0\farcs065$, or slightly larger than 1 STIS pixel."
 Thus. the combined uncertainty in andHST astrometry is less than 072: we use a conservative error circle of radius 072 in Figure 4 to indicate the location of the X-ray source on the STIS image.," Thus, the combined uncertainty in and astrometry is less than $0\farcs2$ ; we use a conservative error circle of radius $0\farcs2$ in Figure 4 to indicate the location of the X-ray source on the STIS image."
 Since the and observations were made only one month apart. we are confident that any proper motion. already limited to «1 vi. is of no importance.," Since the and observations were made only one month apart, we are confident that any proper motion, already limited to $< 1^{\prime\prime}$ $^{-1}$, is of no importance."
 The error circle excludes all of the optical objects within theROSAT error circle that were detected by Totani. Kawasaki. Kawai (2002) in their B-band image obtained on the Subaru telescope.," The error circle excludes all of the optical objects within the error circle that were detected by Totani, Kawasaki, Kawai (2002) in their $B$ -band image obtained on the Subaru telescope."
 At the northeast edge of the error circle there is only a marginal source of Vz29.0+0.4 in the STIS 50CCD image. but it is not present in the F28X50LP image.," At the northeast edge of the error circle there is only a marginal source of $V \approx 29.0 \pm 0.4$ in the STIS 50CCD image, but it is not present in the F28X50LP image."
 Since this is not even a 3o detection. we consider that the X-ray source ts formally undetected optically. with upper limits of V>28.5 and />26.5. and that it must therefore be a neutron star with fy/>6000.," Since this is not even a $3\sigma$ detection, we consider that the X-ray source is formally undetected optically, with upper limits of $V > 28.5$ and $I > 26.5$, and that it must therefore be a neutron star with $f_X/f_V > 6000$."
 The absence of an optical detection places additional constraints on the distance and temperature of the neutron starJ1836., The absence of an optical detection places additional constraints on the distance and temperature of the neutron star.
24+5925.. A magnitude limit of V>28.5 corresponds to a flux <0.014pJy at a wavelength of 5500 A.. whereas the Rayleigh-Jeans flux from a neutron star at that wavelength would be 0.03075(R46diooYid y. where 15 is T4. in units of 10° Κ. Ryo is the radius in units of 10 km. and dioo is the distance in units of 100 pc.," A magnitude limit of $V > 28.5$ corresponds to a flux $< 0.014\,\mu$ Jy at a wavelength of 5500 , whereas the Rayleigh-Jeans flux from a neutron star at that wavelength would be $0.030\,T_5\,(R_{10}/d_{100})^2\,\mu$ Jy, where $T_5$ is $T_{\infty}$ in units of $10^5$ K, $R_{10}$ is the radius in units of 10 km, and $d_{100}$ is the distance in units of 100 pc."
 Therefore. the nominal X-ray fitted temperature of 3«10? K. if coming from the full surface of the neutron star. would require d7250 pc. ," Therefore, the nominal X-ray fitted temperature of $3 \times 10^5$ K, if coming from the full surface of the neutron star, would require $d > 250$ pc. }"
The absence of a “strong.” re. ~0.5 mJy radio source at the position of wwas already known from Paper |. In order to make a more sensitive search for radio pulsations. we obtained six observations on different days in 2001 February and March for 2.3 hr each with the Jodrell Bank 76 m Lovell telescope at a frequency of 1400 MHz.," The absence of a “strong,” i.e, $\sim 0.5$ mJy radio source at the position of was already known from Paper I. In order to make a more sensitive search for radio pulsations, we obtained six observations on different days in 2001 February and March for 2.3 hr each with the Jodrell Bank 76 m Lovell telescope at a frequency of 1400 MHz."
 A 64« 1MMHz filter bank and a sampling time of | ms were used., A $64\times 1$ MHz filter bank and a sampling time of 1 ms were used.
 Since the maximum expected dispersion measure is small at this high-latitude location. DM <17 pe οπ for d<1 kpe according to the Taylor Cordes (1993) model. the limiting flux density for pulsations is insensitive to distance or period in the range P>100 ms. being 0.1 mJy for a duty cycle of10%.," Since the maximum expected dispersion measure is small at this high-latitude location, DM $\leq
17$ pc $^{-3}$ for $d \leq 1$ kpc according to the Taylor Cordes (1993) model, the limiting flux density for pulsations is insensitive to distance or period in the range $P > 100$ ms, being 0.1 mJy for a duty cycle of."
 A period in this range is expected for an older.Geminga-like pulsar.," A period in this range is expected for an older,Geminga-like pulsar."
 For P<100 ms. the sensitivity diminishes rapidly. to zz0.25 mJy at 10ms.," For $P < 100$ ms, the sensitivity diminishes rapidly, to $\approx 0.25$ mJy at 10ms."
 Scintillation might be expected for these DM values and observing parameters. but our six observations," Scintillation might be expected for these DM values and observing parameters, but our six observations"
Before normalising the composite spectra. we visually compare the continuum slopes of the three composites.,"Before normalising the composite spectra, we visually compare the continuum slopes of the three composites."
 We tind no significant trend for reddening with FIR luminosity in the continuum slopes of the composite spectra., We find no significant trend for reddening with FIR luminosity in the continuum slopes of the composite spectra.
 Differences between the composite slopes are negligible compared to the dispersion of slopes among the individual quasars making up each composite., Differences between the composite slopes are negligible compared to the dispersion of slopes among the individual quasars making up each composite.
 We fit a power-lw continuum to the FIR-intermediate composite spectrum using wavelength regions devoid of absorption or emission. 25 wide and centred at. A11460. and. A11770. followingWarneretal.(2003).," We fit a power-law continuum to the FIR-intermediate composite spectrum using wavelength regions devoid of absorption or emission, 25 wide and centred at 1460, and 1770, following\citet{Warner03}."
.. Unfortunately. the FIR-faint and FIR-bright composites have incomplete spectral coverage past ~A 11700Α.. so we fit the continuum using the available wavelength region around 11460.À.. plus the region from A11335-1355Α.. which is also— used by Juarezetal. (20091. and," Unfortunately, the FIR-faint and FIR-bright composites have incomplete spectral coverage past $\sim$ 1700, so we fit the continuum using the available wavelength region around 1460, plus the region from 1335-1355, which is also used by \citet{Juarez09}, , and"
"have access to the surface and are capable of eventual evaporation, or we can assume that only a layer near the surface can evaporate.","have access to the surface and are capable of eventual evaporation, or we can assume that only a layer near the surface can evaporate."
" The final colors of Kuiper belt objects depend only on the chemistry in a very small layer near the surface, so both types of objects would appear identical on the surface, but the amount of material that needs to be evaporated for an object to appear depleted differs greatly in these two scenarios."," The final colors of Kuiper belt objects depend only on the chemistry in a very small layer near the surface, so both types of objects would appear identical on the surface, but the amount of material that needs to be evaporated for an object to appear depleted differs greatly in these two scenarios."
" Even very small short period comets still appear to have volatiles such as CO, which would quickly evaporate in the surface layers of almost any object in the primordial Kuiper belt (Bockelée-Morvanetal.2004)."," Even very small short period comets still appear to have volatiles such as CO, which would quickly evaporate in the surface layers of almost any object in the primordial Kuiper belt \citep{2004come.book..391B}."
. We thus assume that evaporation is only a surficial process and we set the depth required to deplete the surface to 100 m. Changing this depth by orders of magnitude in either direction does not qualitatively change the conclusions below., We thus assume that evaporation is only a surficial process and we set the depth required to deplete the surface to 100 m. Changing this depth by orders of magnitude in either direction does not qualitatively change the conclusions below.
 The second important assumption is the length of time that passes from when the disk dissipates and the objects are exposed to sunlight to when the objects are scattered to their more distant locations., The second important assumption is the length of time that passes from when the disk dissipates and the objects are exposed to sunlight to when the objects are scattered to their more distant locations.
" In the canonical Nice model, the scattering by the planetary instability is the cause of the Late Heavy Bombardment 650 Myr after the formation of the solar system (Gomesetal. so the primordial objects spend a long period of 2005),,time closer to the sun."," In the canonical Nice model, the scattering by the planetary instability is the cause of the Late Heavy Bombardment 650 Myr after the formation of the solar system \citep{2005Natur.435..466G}, so the primordial objects spend a long period of time closer to the sun."
" Such a long period is not a requirement of an instability model, however, so the actual time remains unconstrained."," Such a long period is not a requirement of an instability model, however, so the actual time remains unconstrained."
" Regardless of the exposure time, an irradiated crust can develop much more quickly (Hudsonetal.2008),, so additional exposure time can no longer affect colors."," Regardless of the exposure time, an irradiated crust can develop much more quickly \citep{2008ssbn.book..507H}, so additional exposure time can no longer affect colors."
" We set our exposure time to be 10 Myr, but, again, values differing by an order of magnitude in either direction do not qualitatively change the conclusions below."," We set our exposure time to be 10 Myr, but, again, values differing by an order of magnitude in either direction do not qualitatively change the conclusions below."
" For a wide range of assumptions, evaporation in the early Kuiper belt behaves as shown in Figure 1."," For a wide range of assumptions, evaporation in the early Kuiper belt behaves as shown in Figure 1."
 Water ice is involatile at all distances and COs and H25 are involatile throughout the Kuiper belt., Water ice is involatile at all distances and $_2$ and $_2$ S are involatile throughout the Kuiper belt.
" Objects residing in the outer parts of the Kuiper belt retain CH3OH, then C2H2, C3He, and HCN over a small range of distances."," Objects residing in the outer parts of the Kuiper belt retain $_3$ OH, then $_2$ $_2$, $_2$ $_6$, and HCN over a small range of distances."
" NHs is retained only near the distance of the current cold classical Kuiper belt, and CH4, No, and CO are depleted on surface layers throughout the Kuiper belt except for the largest objects."," $_3$ is retained only near the distance of the current cold classical Kuiper belt, and $_4$, $_2$, and CO are depleted on surface layers throughout the Kuiper belt except for the largest objects."
" For the specific parameters chosen, the evaporation line of CH3OH appears near 20 AU, which would be in the middle of the primordial disk of KBOs."," For the specific parameters chosen, the evaporation line of $_3$ OH appears near 20 AU, which would be in the middle of the primordial disk of KBOs."
" Experiments on ice irradiation in the outer solar system have primarily focused on specific chemical pathways and species rather than less precise coloration (Hudsonetal.2008),, nonetheless some trends appear clear."," Experiments on ice irradiation in the outer solar system have primarily focused on specific chemical pathways and species rather than less precise coloration \citep{2008ssbn.book..507H}, nonetheless some trends appear clear."
" Objects that form between the inner edge of the primordial Kuiper belt and approximately 20 AU will have temperatures sufficiently high that the only major ices that remain on the surface are H20, COs, and (with a much smaller abundance) H2S. No specific experiments have been done on the coloration or albedo of such a mixture after irradiation, but irradiation of H3O, 00» mixtures is known to produce carbonic acid and more complex hydrocarbons (Mooreetal.1991;Delitsky&Lane 1998).."," Objects that form between the inner edge of the primordial Kuiper belt and approximately 20 AU will have temperatures sufficiently high that the only major ices that remain on the surface are $_2$ O, $_2$, and (with a much smaller abundance) $_2$ S. No specific experiments have been done on the coloration or albedo of such a mixture after irradiation, but irradiation of $_2$ O, $_2$ mixtures is known to produce carbonic acid and more complex hydrocarbons \citep{1991JGR....9617541M, 1998JGR...10331391D}."
" Irradiation of such hydrocarbons then leads to the loss of hydrogen, the production of larger carbon chains, and the eventual carbonization of the surface."," Irradiation of such hydrocarbons then leads to the loss of hydrogen, the production of larger carbon chains, and the eventual carbonization of the surface."
 The final product is a dark neutrally colored spectrally bland surface (Andronicoetal.1987;Moroz2004;Palumboetal. 2004)..," The final product is a dark neutrally colored spectrally bland surface \citep{1987A&A...184..333A, 2004Icar..170..214M,2004AdSpR..33...49P}."
" Irradiation of a H30, 00ο mix has been speculated to be a cause of the very dark crust of Callisto (McCordetal.1998).."," Irradiation of a $_2$ O, $_2$ mix has been speculated to be a cause of the very dark crust of Callisto \citep{1998JGR...103.8603M}."
 Such surfaces describe the dark neutrally colored centaurs well., Such surfaces describe the dark neutrally colored centaurs well.
" In addition, with the exception of some of the larger water ice rich KBOs, the neutral KBOs have low albedos similar to those of the centaurs (Stansberryetal. 2006)."," In addition, with the exception of some of the larger water ice rich KBOs, the neutral KBOs have low albedos similar to those of the centaurs \citep{2006ApJ...643..556S}."
 We propose that the objects that are now the neutral-colored KBOs were formed in the inner part of the primordial disk and scattered into the current Kuiper belt., We propose that the objects that are now the neutral-colored KBOs were formed in the inner part of the primordial disk and scattered into the current Kuiper belt.
" When they scatter inward to become centaurs their temperatures do not increase markedly beyond those that were experiences at formation, so surfaces do not evolve significantly."," When they scatter inward to become centaurs their temperatures do not increase markedly beyond those that were experiences at formation, so surfaces do not evolve significantly."
" Beyond approximately 20 AU, several major hydrocarbon species can remain on the surface."," Beyond approximately 20 AU, several major hydrocarbon species can remain on the surface."
 The most abundant of these is methanol, The most abundant of these is methanol.
 A Raman study of residues remaining after methanol irradiation showed that the methanol residue is surprisingly lacking in the signature of amorphous carbon (Ferinietal., A Raman study of residues remaining after methanol irradiation showed that the methanol residue is surprisingly lacking in the signature of amorphous carbon \citep{2004A&A...414..757F}.
" Indeed, Brunettoetal. have shown that methanol,2004).. when irradiated to dosages(2006) expected for solar system aged KBOs (Cooperetal.2003),, does not turn dark and neutral but instead retains higher albedos and redder colors."," Indeed, \citet{2006ApJ...644..646B}
 have shown that methanol, when irradiated to dosages expected for solar system aged KBOs \citep{2003EM&P...92..261C}, does not turn dark and neutral but instead retains higher albedos and redder colors."
 The colors and albedos are similar to those seen in the red centaurs and likely also the medium-size red KBOS., The colors and albedos are similar to those seen in the red centaurs and likely also the medium-size red KBOS.
" Methanol, intriguingly, is also the only involatile molecule other than water identified either on a centaur (Cruikshanketal.1998) or a KBO (Barucci 2006).."," Methanol, intriguingly, is also the only involatile molecule other than water identified either on a centaur \citep{1998Icar..135..389C} or a KBO \citep{2006A&A...455..725B}."
 We propose that the presence of methanol on the primordial surface of a KBO allows that KBO to maintain a higher albedo redder irradiation crust., We propose that the presence of methanol on the primordial surface of a KBO allows that KBO to maintain a higher albedo redder irradiation crust.
" It is possible — indeed likely — that that the C2Hg and the C2H2 and HCN evaporation lines, which are just beyond"," It is possible – indeed likely – that that the $_2$ $_6$ and the $_2$ $_2$ and HCN evaporation lines, which are just beyond"
more consistent with Phase III evolution during this period from the behaviour of the X-ray flux.,more consistent with Phase III evolution during this period from the behaviour of the X-ray flux.
" The usable energy range of the PCA (2-25 keV) means however that their observations were less sensitive to Ny, aud (they did not detect the emergence of the soft component.", The usable energy range of the PCA (2-25 keV) means however that their observations were less sensitive to $N_H$ and they did not detect the emergence of the soft component.
 Ab =13.6 davs. our fits to the data imply οςc1700 knis t.," At $t = 13.6$ days, our fits to the data imply $v_s \simeq 1700$ km $^{-1}$."
" We note that we could obtain complete consistency between the VLBA imagery and simple derivation of v, from the model fits iLa —2.7xLOM em 7 and a larger distance of d=2.4 kpe. if the remnant were behaving precisely as one would expect for the Primakoff (Phase II) solution throughout."," We note that we could obtain complete consistency between the VLBA imagery and simple derivation of $v_s$ from the model fits if $a = 2.7 \times 10^{10}$ cm $^{-2/3}$ and a larger distance of $d = 2.4$ kpc, if the remnant were behaving precisely as one would expect for the Primakoff (Phase II) solution throughout."
 However. the time dependence of ὃν. μι ond. Εγω. are more like what one would expect in Phase HI. but in this case b—3.7xLOM cem |? and for consistency with ihe VLBA results. the distance would have to be increased to 3 kpc.," However, the time dependence of $v_s$, $[N_H]_W$ and $F_{unabs}$ are more like what one would expect in Phase III, but in this case $b = 3.7 \times 10^{11}$ cm $^{-1/2}$ and for consistency with the VLBA results, the distance would have to be increased to 3 kpc."
 Such high values of (he distance were accepted prior to 1985 (Bode1987)., Such high values of the distance were accepted prior to 1985 \citep{bods87}.
. On the other hand. Snijders(1957) notes that the failure to detect material at the tvpieal velocities of the Carina Arm inZUE speclra places an upper limit on the distance to RS Oph of about 2 kpc.," On the other hand, \citet{sni87} notes that the failure to detect material at the typical velocities of the Carina Arm in spectra places an upper limit on the distance to RS Oph of about 2 kpc."
 In addition of course. we have shown that there is a transition at around /=6 davs between remnant phases.," In addition of course, we have shown that there is a transition at around $t = 6$ days between remnant phases."
 Thus although the simple aualvsis is supportive of the «qualitative correctness of the current models. as outlined below. detailed physical modelling is required to build a fully self-consistent picture.," Thus although the simple analysis is supportive of the qualitative correctness of the current models, as outlined below, detailed physical modelling is required to build a fully self-consistent picture."
 Our X-ray observations of (he very earliest phases of the 2006 outburst of RS Oph are broadly consistent with the basic model of remnant evolution proposed by Bode and further explored by O'Brienetal.(1992)., Our X-ray observations of the very earliest phases of the 2006 outburst of RS Oph are broadly consistent with the basic model of remnant evolution proposed by \citet{bod85} and further explored by \citet{obr92}.
. In particular. it appears that Phase I mav have ended by |~ 6d as predicted.," In particular, it appears that Phase I may have ended by $t \sim 6$ d as predicted."
 ILowever. our first-order analysis suggests that the remnant moved rapidly into Phase HI in this outburst.," However, our first-order analysis suggests that the remnant moved rapidly into Phase III in this outburst."
 Fitstothe XRT data have also been attempted using vinekal where (he abundances have been allowed to vary., Fits to the XRT data have also been attempted using vmekal where the abundances have been allowed to vary.
 For the single temperature fit. (his resulted in a trend of decreasing elemental abundance enhancement with (ime. as would be expected as emission lrom the enriched nova ejecta at early times becomes increasingly less significant than that [rom the eenerally less enriched red giant wind.," For the single temperature fit, this resulted in a trend of decreasing elemental abundance enhancement with time, as would be expected as emission from the enriched nova ejecta at early times becomes increasingly less significant than that from the generally less enriched red giant wind."
 A series of fits using multi-temperature mekals was also performed producing improved fits to the data., A series of fits using multi-temperature mekals was also performed producing improved fits to the data.
 However. detailed models involving full hvdrodynamical simulations are ultimatelv required (ο [it the X-ray data in a physically meaningful wav.," However, detailed models involving full hydrodynamical simulations are ultimately required to fit the X-ray data in a physically meaningful way."
 ‘These simulations produce an evolving range of temperatures and densities with radius from the central source (see O'Brienetal. (1992)))., These simulations produce an evolving range of temperatures and densities with radius from the central source (see \citet{obr92}) ).
 Resonant scattering of in the red giant wind may also be occurring. and therefore some of the X-ray. emission lineflix may be due to photoionization of this gas by the X-ray emitting shock.," Resonant scattering of X-rays in the red giant wind may also be occurring, and therefore some of the X-ray emission lineflux may be due to photoionization of this gas by the X-ray emitting shock."
 In addition.," In addition,"
also included in the fit to model the iron line features.,also included in the fit to model the iron line features.
 Further to this an edge was added at 0.7393 keV. The fit parameters are shown in Table 5.., Further to this an edge was added at 0.7393 keV. The fit parameters are shown in Table \ref{J1616_spectral_fits_parameters}.
 There is no candidate period in the power spectrum that indicates a coherent modulation associated with a potential spin period and the 585 s (~148 cyeles day~!) QPO seen by ? is not seen here either., There is no candidate period in the power spectrum that indicates a coherent modulation associated with a potential spin period and the 585 s $\sim148$ cycles $^{-1}$ ) QPO seen by \citet{pretorius09} is not seen here either.
 The 10.2 hr peak that we see is at roughly twice the orbital period reported by ?.. however the uncertainty here ts relatively large. so this may be coincidence.," The 10.2 hr peak that we see is at roughly twice the orbital period reported by \citet{pretorius09}, however the uncertainty here is relatively large, so this may be coincidence."
 The spectrum ts fairly typical of IPs and suggests a power law with a partial covering absorber with tron line features., The spectrum is fairly typical of IPs and suggests a power law with a partial covering absorber with iron line features.
 Further to this. the detection of the edge implies the existence of a warm absorber - likely the pre-shock flow being photoionized.," Further to this, the detection of the edge implies the existence of a warm absorber - likely the pre-shock flow being photoionized."
 This has been seen in two other [Ps (V1223 Ser (?) and IRXS J173021.5-055933 (?))) and further adds to the case for IP classification., This has been seen in two other IPs (V1223 Sgr \citep{mukai01} and 1RXS J173021.5-055933 \citep{demartino08}) ) and further adds to the case for IP classification.
 It should be noted that two other models gave only a slightly worse fit - two partial covering absorbers on an Model and a power law distribution of covering fractions (PWAB) on aMKCFLOW. ¥7= 1.07 and 1.06 respectively., It should be noted that two other models gave only a slightly worse fit - two partial covering absorbers on an model and a power law distribution of covering fractions ) on a. $\chi_r^2$ = 1.07 and 1.06 respectively.
 The optical spectra taken by ? and ? along with the X-ray spectrum seen here point towards an IP classification of J1616.," The optical spectra taken by \citet{masetti06} and \citet{pretorius09}
 along with the X-ray spectrum seen here point towards an IP classification of J1616."
 However. the lack of any feasible spin period candidate in either the optical or X-ray means that the classification cannot be confirmed.," However, the lack of any feasible spin period candidate in either the optical or X-ray means that the classification cannot be confirmed."
 Perhaps the geometry of the accretion column is such that the magnetic and spin axis are aligned (or very nearly aligned). so no spin modulation would be seen. as discussed by 2..," Perhaps the geometry of the accretion column is such that the magnetic and spin axis are aligned (or very nearly aligned), so no spin modulation would be seen, as discussed by \citet{ramsay08}."
 The IP classification of J1616 is thus unproven. though it remains as a candidate system.," The IP classification of J1616 is thus unproven, though it remains as a candidate system."
 V2487 Oph was discovered in the optical as a possible nova in 1998 at magnitude 9.5 (2).., V2487 Oph was discovered in the optical as a possible nova in 1998 at magnitude 9.5 \citep{nakano98}.
 ? showed a plot of the rapid decline in visual magnitude (courtesy of the AAVSO). indicating V2487 Oph was a very fast nova.," \citet{lynch00} showed a plot of the rapid decline in visual magnitude (courtesy of the AAVSO), indicating V2487 Oph was a very fast nova."
" They also presented NIR spectra, showing an overabundance in carbon soon after the outburst."," They also presented NIR spectra, showing an overabundance in carbon soon after the outburst."
 ? modelled the optical light curve and concluded it has a WD mass of 1.35+40.0144... and the mass transfer rate indicated a recurrence period of about 40 years.," \citet{hachisu02} modelled the optical light curve and concluded it has a WD mass of $\pm$ $M_{\odot}$, and the mass transfer rate indicated a recurrence period of about 40 years."
 ? reported data of V2487 Oph taken 2.7 years after its discovery., \citet{hernanz02} reported data of V2487 Oph taken 2.7 years after its discovery.
" They fit their spectrum with a two-temperature plasma model (Tiny = 0.2 keV and T5;,> 48 keV). which suggests a shocked gas (due an accretion flow) was present."," They fit their spectrum with a two-temperature plasma model $T_{low}$ = 0.2 keV and $T_{high} \geq$ 48 keV), which suggests a shocked gas (due an accretion flow) was present."
 They also find an iron line at 6.4 keV. The error circles of V2487 Oph and IRXS J173200.0-191349 are coincident. and the and fluxes are similar. indicating they may be the same source.," They also find an iron line at 6.4 keV. The error circles of V2487 Oph and 1RXS J173200.0–191349 are coincident, and the and fluxes are similar, indicating they may be the same source."
 V2487 Oph was then found in the survey (?).., V2487 Oph was then found in the survey \citep{barlow06}.
 ? searched archival data for previous eruptions and found one occurred in 1900. confirming V2487 Oph as a recurrent nova.," \citet{pagnotta09} searched archival data for previous eruptions and found one occurred in 1900, confirming V2487 Oph as a recurrent nova."
 V2487 Oph was observed by over two consecutive days (see Table 1))., V2487 Oph was observed by over two consecutive days (see Table \ref{observing_log}) ).
 The total good time was 3312 s in PCU2., The total good time was 312 s in PCU2.
 The 2-10 keV energy band raw count rate varied between 2.8-6.3 et s!. and the background count rate (generated from the calibration files) 072.8 et s!.," The 2–10 keV energy band raw count rate varied between 2.8–6.3 ct $^{-1}$ , and the background count rate (generated from the calibration files) 0–2.8 ct $^{-1}$."
 The average background subtracted count rate was 1.1 et s., The average background subtracted count rate was 1.1 ct $^{-1}$ .
 There is another X-ray source in the field of view that we would expect to contribute a comparable count rate to V2487 Oph., There is another X-ray source in the field of view that we would expect to contribute a comparable count rate to V2487 Oph.
 Fig., Fig.
 15 shows the cLEANed power spectrum of the background subtracted 2-10 keV light curve., \ref{V2487_cleaned} shows the ed power spectrum of the background subtracted 2–10 keV light curve.
 The largest peak is at approximately 367 cycles day”! (235.2+0.1 s). the second largest peak is at 18.950.1 s - very close (but not within errors) of beinga potential first harmonie of the235.2 peak.," The largest peak is at approximately 367 cycles $^{-1}$ $\pm$ 0.1 s), the second largest peak is at $\pm$ 0.1 s - very close (but not within errors) of beinga potential first harmonic of the235.2 peak."
 The folded data (in all bands) shows only marginally coherent modulation.," The folded data (in all bands) shows only marginally coherent modulation,"
GC stripping. so that the decrement of 5x depends strongly on thatof Noe. (,"GC stripping, so that the decrement of $S_{\rm N}$ depends strongly on thatof $N_{\rm GC}$. ("
"ui) Final Spx depends on (4; such that models with smaller e, have larger final Six because of its less ellicient tidal stripping of GC's.",iii) Final $S_{\rm N}$ depends on $a_{\rm gc}$ such that models with smaller $a_{\rm gc}$ have larger final $S_{\rm N}$ because of its less efficient tidal stripping of GCs.
 Comparison of our results ancl the observations suggests that in order to explain. both the observed. Nec and. Sy.large. (," Comparison of our results and the observations suggests that in order to explain both the observed $N_{\rm GC}$ and $S_{\rm N}$,. ("
Either Model 2 or 3 can. best. reproduce. the observed properties of NGC 1404 in the present study). (,Either Model 2 or 3 can best reproduce the observed properties of NGC 1404 in the present study). (
iv) Our models predict that Sp correlates with the distance of GCs host galaxy. from the centre of Fornax cluster (equivalently. from NCC 1399) in such a way tha a galaxy with the larger distance shows larger Sp.,"iv) Our models predict that $S_{\rm N}$ correlates with the distance of GCs' host galaxy from the centre of Fornax cluster (equivalently, from NGC 1399) in such a way that a galaxy with the larger distance shows larger $S_{\rm N}$."
 In Fig 7 we plot the observed SN of Fornax ellipticals as a function of eluster-centric radius., In Fig 7 we plot the observed SN of Fornax ellipticals as a function of cluster-centric radius.
 Phese data suggest that galaxies a larger racii have higher SN., These data suggest that galaxies at larger radii have higher SN.
 The location of NGC 1404 GC's in Fig., The location of NGC 1404 GCs in Fig.
 6 also suggests that if the observed. lower Sy anc Noo in NGC 1404 are due to tidal stripping by the cluster tidal field. ὃν should be as large as (or larger than) 0.5 ane (ao ds 0 2. (," 6 also suggests that if the observed lower $S_{\rm N}$ and $N_{\rm GC}$ in NGC 1404 are due to tidal stripping by the cluster tidal field, $e_{\rm p}$ should be as large as (or larger than) 0.5 and $a_{\rm gc}$ is $\sim$ 2. ("
"v) Projected. surface number. density profiles ane velocity dispersion of LOGCs (around NGC 1399) are likely to be shallower and smaller. respectively. in the model with smaller ey. though we could investigate these only for the models with a, = 2.0 because of the very smaller number of LOGCs in other models.","v) Projected surface number density profiles and velocity dispersion of ICGCs (around NGC 1399) are likely to be shallower and smaller, respectively, in the model with smaller $e_{\rm p}$, though we could investigate these only for the models with $a_{\rm gc}$ = 2.0 because of the very smaller number of ICGCs in other models."
 These results indicate that physical properties of LOGCs in. a cluster depends on the orbital populations (e.e.. the mean orbital eccentricitv) of galaxies in the cluster.," These results indicate that physical properties of ICGCs in a cluster depends on the orbital populations (e.g., the mean orbital eccentricity) of galaxies in the cluster."
 So far. we have focused on just one elliptical galaxy (NGC 1404) in Fornax cluster and the origin of its relatively low Sx.," So far, we have focused on just one elliptical galaxy (NGC 1404) in Fornax cluster and the origin of its relatively low $S_{\rm
N}$."
 Based on these results. we propose that we can assess the importance of tidal stripping by strong cluster gravitational fields in the formation of low Sx elliptical galaxies generat by. checking the following three observable physical properties of GC svstems of cluster cllipticals.," Based on these results, we propose that we can assess the importance of tidal stripping by strong cluster gravitational fields in the formation of low $S_{\rm N}$ elliptical galaxies by checking the following three observable physical properties of GC systems of cluster ellipticals."
 Phe foremost is, The foremost is
" GRBlljorthetal.2003.. Bloometal.2004:GRB Cobbetal.2004.. Malesanietal.2004.. ThomsenetMal.2004:, Campanael2006.. Ferreroetal.2006.. Mirabaletal.2006.. Cobbetal.2006a.andreferencestherein: DellaValleetal.2006a:: Wooslev&Bloom2006)). M(GRDMM redshift2=0.0889 at2=0.125: €""el DellaValle--etds2006b.. Fvnboetal.2006.. Gal-Yainetal.2006.. Nuetal.».2009)). apparent”(Toy=r.lss. Trojanaietal.2009)). GRB (Gehrelsetal.2004) (UT Trojaelal.2009)). 2.05+0.07 (—0.4pape AL7.5 9.02€0.3xLO© 7 ουeldes2009).."," \citealt{Hjorth+03}, \citealt{Stanek+03}, \citealt{Bloom+04}; \citealt{Cobb+04}, \citealt{Gal-Yam+04},\citealt{Malesani+04}, \citealt{Thomsen+04}; \citealt{Campana+06}, \citealt{Ferrero+06}, \citealt{Mirabal+06}, \citealt[][and references therein]{Cobb+06a}; \citealt{DellaValle+06a}; \citealt{Woosley+Bloom06}) $z=0.0889$ $z=0.125$ \citealt{Cobb+06b}, \citealt{DellaValle+06b}, \citealt{Fynbo+06}, \citealt{Gal-Yam+06}, \citealt{Xu+09}) $T_{90} = 7.1$ \citealt{gcn10191}) \citep{Gehrels+04} \citealt{gcn10191}) $2.05\pm0.07$ $-0.4$ $7.5$ $9.0\pm0.3\times10^{-6}$ $^{-2}$ \citep{gcn10197}."
"""a 1.98üun9.5(he7...> (Evansetal.ne2009).. (Smithetal.2009)... Romingetal.2005))", $1.98^{+0.15}_{-0.14}$$9.8^{+3.3}_{-3.1}\times10^{20}$$^{-2}$ \citep{gcn10201}. \citep{gcn10192}. \citealt{Roming+05})
at birth of order 10 is. a second branch is mace up of the anomalous XN-rav pulsars (ANPs: Mereghetti Stella 1998 aud refs.,"at birth of order 10 ms, a second branch is made up of the anomalous X-ray pulsars (AXPs; Mereghetti Stella 1998 and refs."
 therem: Duncan Thompson 1996) and the soft x-ray repeaters (SCRs: Cline et al., therein; Duncan Thompson 1996) and the soft $\gamma$ -ray repeaters (SGRs; Cline et al.
" 1982: Ίνκα Frail 1993). likely ""maguetars. with magnetic fields in the ranee 1014.102? ο (Vasisht Cottliclf 1997: Ixouveliotou et al."," 1982; Kulkarni Frail 1993), likely “magnetars”, with magnetic fields in the range $10^{14} - 10^{15}$ G (Vasisht Gotthelf 1997; Kouveliotou et al."
 1998)., 1998).
 The maguetars typically ive long spin periods. aud their steady cussion las ouly heen observed in the N-rav band.," The magnetars typically have long spin periods, and their steady emission has only been observed in the X-ray band."
" The new pulsar lies ina --- rege with B,25«&1055 G. Few regular pulsars have implied magnetic fields strictly above the quantum critical field. B,,~LbsLot? Ce in act. the only such pulsar kuown thus far is the radio osa PSR 17114 which has an interred field of Bi,=~5.5 (Pivavoroff. Ikaspi Camilo 2000: Caruilo et al"," The new pulsar lies in a transitional regime with $B_p \simeq 5 \times 10^{13}$ G. Few regular pulsars have implied magnetic fields strictly above the quantum critical field, $B_{cr} \simeq 4.4 \times 10^{13}$ G; in fact, the only such pulsar known thus far is the radio pulsar PSR $-$ 1744 which has an inferred field of $B_{13} \simeq 5.5$ (Pivavoroff, Kaspi Camilo 2000; Camilo et al."
 2000). just above this limit.," 2000), just above this limit."
 Free electrons. evrate relativistically in Boo>ιν with radi less than the electron Compton wavelength. ο.," Free electrons gyrate relativistically in $B > B_{cr}$ with radii less than the electron Compton wavelength, $\hbar/m_ec$."
 Tn regular pulsars such as the Crab aud Vela the purely quantum process of single photou pai-production +>ele is invoked as a source of particle acceleration (Sturrock 1971)., In regular pulsars such as the Crab and Vela the purely quantum process of single photon pair-production $\gamma \rightarrow e^{+}e^{-}$ is invoked as a source of particle acceleration (Sturrock 1971).
 It has been sueeested that for magnetars with Do>D... the quantum clectrodyvnamical process of photon splitting (5340). may compete with pair production and act as a quenching mechanisin for clectrous. auc suppress radio cluission (Baring Ibudiug 1998).," It has been suggested that for magnetars with $B > B_{cr}$, the quantum electrodynamical process of photon splitting $\gamma \rightarrow
\gamma\gamma$ ), may compete with pair production and act as a quenching mechanism for electrons, and suppress radio emission (Baring Harding 1998)."
" Note that the large pulse duty cvcle in ssueeests that unlike the Crab. which has a sharp ligh-energv pulse. the X-ray producing particles in this pulsar are largely in the outer maguctosphere,"," Note that the large pulse duty cycle in suggests that unlike the Crab, which has a sharp high-energy pulse, the X-ray producing particles in this pulsar are largely in the outer magnetosphere."
 Also. it remains to be seen if ls detectable as a radio pulsar.," Also, it remains to be seen if is detectable as a radio pulsar."
 The shape of the radio pulse profile. vis-a-vis the N-rav pulse would also be of considerable iuterest.," The shape of the radio pulse profile, vis-a-vis the X-ray pulse would also be of considerable interest."
Y5O.,YSO.
 As illustrated bv the simulation described in the next paragraph. the most massive objects erow most rapidly. doubling (their mass and migrating to the center of the cloud within fewx10? vears.," As illustrated by the simulation described in the next paragraph, the most massive objects grow most rapidly, doubling their mass and migrating to the center of the cloud within $few \times 10^5$ years."
 Fragmentation (Melee Tan 2002: 2003: IXxrunholz 2006) or the capture of sibling stus (Moeckel Bally 2006: 2007a.b) may produce multiple svstenis.," Fragmentation (McKee Tan 2002; 2003; Krumholz 2006) or the capture of sibling stars (Moeckel Bally 2006; 2007a,b) may produce multiple systems."
 These processes could have resulted in the assembly of a sub-cluster of massive stars and binaries in the center of the core., These processes could have resulted in the assembly of a sub-cluster of massive stars and binaries in the center of the core.
 Figure 4 shows a sample result of a simplified numerical model illustrating the combined effects of Boncli-Llovle accretion. onto. protostellar seeds. orbit decay. formation of a hierarchical svstem of massive stars. ancl dvnamical ejection.," Figure 4 shows a sample result of a simplified numerical model illustrating the combined effects of Bondi-Hoyle accretion onto protostellar seeds, orbit decay, formation of a non-hierarchical system of massive stars, and dynamical ejection."
E This simulation starts with a 50 M. clump of gas in à Phunmer potential given by e(r)=—GM/y(r?4a) where A/ is the total mass and @ = 0.025 pe is the core radius. (,Ê This simulation starts with a 50 $_{\odot}$ clump of gas in a Plummer potential given by $\Phi (r) = -GM / \sqrt{(r^2 + a^2)}$ where $M$ is the total mass and $a$ = 0.025 pc is the core radius. (
This potential is used to avoid the singularitv at (he center of an isothermal sphere).,This potential is used to avoid the singularity at the center of an isothermal sphere).
 The clump is non-rotating with an internal velocily dispersion (effective sound speed) given by the local Virial velocity so that. (ο first order. the clamp is stable to global gravitational collapse.," The clump is non-rotating with an internal velocity dispersion (effective sound speed) given by the local Virial velocity so that, to first order, the clump is stable to global gravitational collapse."
 At the start of the simulation. 20 protostellar seeds. each having an initial mass of 0.5 M... are distributed randomly in the chuup.," At the start of the simulation, 20 protostellar seeds, each having an initial mass of 0.5 $_{\odot}$, are distributed randomly in the clump."
 Their velocity distribution is virialized in (he combined gravitational potential of the stars and gas and their velocity vector orientations are random., Their velocity distribution is virialized in the combined gravitational potential of the stars and gas and their velocity vector orientations are random.
" The simulation was run 100 limes. each with a different randomly. chosen set of initial locations and velocities for the seeds,"," The simulation was run 100 times, each with a different randomly chosen set of initial locations and velocities for the seeds."
 Stellar motions are caleulated from the gravitational forces exerted by other stars aud the static gas cloud. modified bv the effects of BIL accretion.," Stellar motions are calculated from the gravitational forces exerted by other stars and the static gas cloud, modified by the effects of BH accretion."
 The inter-stellar forces are directly caleulated. and the orbits are integrated with a global but variable (ime-step using a (7.8) order lRunge-Nutta pair (Prince ancl Dormaucl 1981).," The inter-stellar forces are directly calculated, and the orbits are integrated with a global but variable time-step using a (7,8) order Runge-Kutta pair (Prince and Dormand 1981)."
 The code is a standard integrator with the addition of a static potential. identified with the natal gas. that can be accreted onto protostellar seeds.," The code is a standard n-body integrator with the addition of a static potential, identified with the natal gas, that can be accreted onto protostellar seeds."
 The code follows the stellar mass growth ancl orbit evolution in the combined eravitational potential of the clump plus embedded stars., The code follows the stellar mass growth and orbit evolution in the combined gravitational potential of the clump plus embedded stars.
 Similar approaches have been used by Moeckel and Clarke (2010) ancl Daumezrdt and Ixlessen (2010)., Similar approaches have been used by Moeckel and Clarke (2010) and Baumgardt and Klessen (2010).
 The protostellar seeds experience DII accretion Irom the clump and grow in mass with the accretion radius set to the minimum of the DII radius or one-third the distance to the closest neiehboring star., The protostellar seeds experience BH accretion from the clump and grow in mass with the accretion radius set to the minimum of the BH radius or one-third the distance to the closest neighboring star.
 To keep the computations simple. the response of the gas to the passage of the stars is not modeled.," To keep the computations simple, the response of the gas to the passage of the stars is not modeled."
" This simplification is reasonable because the zone influenced by the passage of a protostellar seed is restricted to ils gravitational radius. rg;=G,MÍV? (S 200 AU for AZ, = 1 AL. and V; = 2 km !)."," This simplification is reasonable because the zone influenced by the passage of a protostellar seed is restricted to its gravitational radius, $r_G = G_* M / V_*^2$ $ \sim $ 200 AU for $M_*$ = 1 $_{\odot}$ and $V_*$ = 2 km $^{-1}$ )."
" Random motions in the chump will tend to fill-in the cavity formed by the passage of a star on a lime scale 7cος,> 10* vears but much less than the 10"" vear evolution time-scale of the protostars ancl clump.", Random motions in the clump will tend to fill-in the cavity formed by the passage of a star on a time scale $\tau \sim r_G / C_s \ge$ $10^3$ years but much less than the $10^5$ year evolution time-scale of the protostars and clump.
 Thus. 1t is assunied (hat on average. the envelope remains spherical auc smooth.," Thus, it is assumed that on average, the envelope remains spherical and smooth."
 However. as gas is accreted onto the stars the mass of the gas cloud. and thus the contribution of the gas to," However, as gas is accreted onto the stars the mass of the gas cloud, and thus the contribution of the gas to"
>40 GeV flux of ~(0.7—L4)x101 photons 7? |. ie. of the same order as the predictions of the leptonic jet models.,"$> 40$ GeV flux of $\sim (0.7 -
1.4) \times 10^{-10}$ photons $^{-2}$ $^{-1}$, i.e. of the same order as the predictions of the leptonic jet models."
 However. in contrast to the leptonic models. the hieh-energv emission is expected to extend bevond 1 TeV at a [lux level of photons 7? !.," However, in contrast to the leptonic models, the high-energy emission is expected to extend beyond 1 TeV at a flux level of $\Phi_{> 1 \, {\rm TeV}} \sim (3 - 18) 
\times 10^{-14}$ photons $^{-2}$ $^{-1}$."
 Such a flux level is well within the reach of future high-sensitivity instruments like VERITAS. (, Such a flux level is well within the reach of future high-sensitivity instruments like VERITAS. (
7) SPD models consistent with the March 1998 EGRET spectrum under-predict the Mav 1998 BeppoSAX PDS hard X-rav spectrum.,7) SPB models consistent with the March 1998 EGRET spectrum under-predict the May 1998 BeppoSAX PDS hard X-ray spectrum.
 They result in higher >40 GeV flixes of ~(5—13)10M photons em7s !. but weaker TeV [Inxes of !! photons 7s !.," They result in higher $> 40$ GeV fluxes of $\sim (5 - 13) \times 10^{-10}$ photons $^{-2}$ $^{-1}$, but weaker TeV fluxes of $\Phi_{> 1 \, {\rm TeV}} \sim 
(0.7 - 9) \times 10^{-14}$ photons $^{-2}$ $^{-1}$."
 In this case. STACEE and CELESTE may be able to get a weak detection of the source. and it would still be a promising candidate for detection by the Πίντο VERITAS array.," In this case, STACEE and CELESTE may be able to get a weak detection of the source, and it would still be a promising candidate for detection by the future VERITAS array."
 In conclusion. leptonic ancl hacronic jet model fits to W Comae make drastically different predietions with respect to the expected verv-high enerey emission bevond ~100 GeV. A detection of W Comae at those photon energies with future. high-sensitivity air Ceerenkov detector arravs woukl pose a serious challenge to leptonic jet models. and mieht favor hadronic models instead.," In conclusion, leptonic and hadronic jet model fits to W Comae make drastically different predictions with respect to the expected very-high energy emission beyond $\sim 100$ GeV. A detection of W Comae at those photon energies with future, high-sensitivity air Čeerenkov detector arrays would pose a serious challenge to leptonic jet models, and might favor hadronic models instead."
 We thank D. Smith for inspiring discussions. and G. Ghisellini for making the data on the May. 1998 SED of Wo Comae available to us.," We thank D. Smith for inspiring discussions, and G. Ghisellini for making the data on the May 1998 SED of W Comae available to us."
 We are also grateful to J. Chiang lor careful reading of (he manuscript and helpful comments., We are also grateful to J. Chiang for careful reading of the manuscript and helpful comments.
 The work of MD was supported by NASA through Chandra Postdoctoral Fellowship grant PF 9-10007 awarded by the Chandra X-ray Center. which is operated by the Smithsonian Astrophysical Observatory for NASA t," The work of MB was supported by NASA through Chandra Postdoctoral Fellowship grant PF 9-10007 awarded by the Chandra X-ray Center, which is operated by the Smithsonian Astrophysical Observatory for NASA under contract NAS 8-39073."
, RM acknowledges support from NSF grant PHY-9983836.
hanks the Bundesministerium lur Bildung und Forschung for financial support through DESY grant. Verbundforschung 05611166., AR thanks the Bundesministerium fürr Bildung und Forschung for financial support through DESY grant Verbundforschung 05CH1PCA6.
llz prohibit anv detection of harmonic content higher (han the fundamental.,Hz prohibit any detection of harmonic content higher than the fundamental.
 We therefore make (he assumption that the pulsations are purely sinusoid. in which case the power expectedars Isjo 0.25or Nor21]ως where Nyused 18 Che number of pulsed photons and ο is the total number of photons.," We therefore make the assumption that the pulsations are purely sinusoid, in which case the power expected is 0.25 $N_{pulsed}^2/N_{tot}$, where $N_{pulsed}$ is the number of pulsed photons and $N_{tot}$ is the total number of photons."
 By setting5 the power to be the observed value of 38 ancl takine5 into account the background. we find that the pulsed fraction is 1622354.," By setting the power to be the observed value of 38 and taking into account the background, we find that the pulsed fraction is $\pm$."
.. As a further attempt to detect pulsations from CNOU JOLL0043.1-721134 we analyzed data from the observation (rp600195a00) which has the highest sensitivity (ο a pulsed signal., As a further attempt to detect pulsations from CXOU J0110043.1-721134 we analyzed data from the observation (rp600195a00) which has the highest sensitivity to a pulsed signal.
 It has the second longest exposure (16 ks) of anv of the PSPC observations aud has bv far the shortest duration (85 ks)., It has the second longest exposure (16 ks) of any of the PSPC observations and has by far the shortest duration (85 ks).
 Other comparable exposure PSPC and IRI observations have durations more than 15 times longer aud (heir sensitivitv (o a periodic signal is diluted by (he search range needed (o cover potential Irequeney. variation over (his length of time., Other comparable exposure PSPC and HRI observations have durations more than 15 times longer and their sensitivity to a periodic signal is diluted by the search range needed to cover potential frequency variation over this length of time.
 The duration of rp600195a00 is such that no phase slippage is expected for Ps in the range of the AXPs., The duration of rp600195a00 is such that no phase slippage is expected for $\dot P$ s in the range of the AXPs.
 The source for this observation is within 12.7. of the center of the PSPC field of view where the angular resolution is excellent. (hus a selection circle of 80 pixels (407) could be used.," The source for this observation is within 12.7' of the center of the PSPC field of view where the angular resolution is excellent, thus a selection circle of 80 pixels (40”) could be used."
 With this selection 392 photons were retrieved. of which an estimated 37.641 are background.," With this selection 392 photons were retrieved, of which an estimated $\pm1$ are background."
 À comparison of the pulse height spectrum for the source and the background. showed that sensitivity to a pulsed signal could be enhanced by eliminating photons with energies less (han 0.4 keV. With (his restriction 358 photons remained.," A comparison of the pulse height spectrum for the source and the background, showed that sensitivity to a pulsed signal could be enhanced by eliminating photons with energies less than 0.4 keV. With this restriction 358 photons remained."
 The (mes of arrival of these photons were adjusted to the barvcenter of the solar system and an fft was performed., The times of arrival of these photons were adjusted to the barycenter of the solar system and an fft was performed.
 We restrictecl our search region to be near one of the two possible frequencies seen in the CHANDRA observation. either 0.1247 or 0.1838 (Iz).," We restricted our search region to be near one of the two possible frequencies seen in the CHANDRA observation, either 0.1247 or 0.1838 (Hz)."
 To account for a lrequency change over (he 9.6 vears between (he (vo observations we assume (hat the object has characteristics of the known AXPs., To account for a frequency change over the 9.6 years between the two observations we assume that the object has characteristics of the known AXPs.
 This implies that the pulsar is spinning down. and that," This implies that the pulsar is spinning down, and that"
"To illustrate the efficiency of the above electron shock surfing mechanism. let us discuss the c component of the equation of motion. where P.=m,U,, is the electron momentum.","To illustrate the efficiency of the above electron shock surfing mechanism, let us discuss the $x$ component of the equation of motion, where $P_x = m_e U_{ex}$ is the electron momentum."
 In our interest. 7. on the right-hand side can be replaced by the electrostatic solitary wave (ESW) produced in the shock transition region.," In our interest, $E_x$ on the right-hand side can be replaced by the electrostatic solitary wave (ESW) produced in the shock transition region."
" If the electric force of €£a is larger than the Lorentz force of (6/c)c,Do. the electrons are trapped and gain their energy."," If the electric force of $e E_{\rm esw}$ is larger than the Lorentz force of $(e/c) v_y B_0$, the electrons are trapped and gain their energy."
 During the non-adiabatic acceleration phase. the velocity vy Increases.," During the non-adiabatic acceleration phase, the velocity $v_y$ increases."
" If the electron satisfies the condition of elo.<(e/e)e,Bo. il can escape Irom ESW and will be convected downstream."," If the electron satisfies the condition of $e E_{\rm esw} < (e/c) v_y B_0$, it can escape from ESW and will be convected downstream."
 The amplitude of ESW max be estimated by equating the wave energy density to the drift energv densitv of the incoming electron (Ishiharaοἱal.1981).. and we obtain. where V;~2ry is the relative velocity between the reflected ions and the incoming electrons.," The amplitude of ESW may be estimated by equating the wave energy density to the drift energy density of the incoming electron \citep{Ish81}, and we obtain, where $V_d \sim 2 v_0$ is the relative velocity between the reflected ions and the incoming electrons."
 The οποιον conversion factor a is of order of OCT)., The energy conversion factor $\alpha$ is of order of $O(1)$.
 Although the nonlinear saturation process of DI still remains controversial. we adapt here (he conversion [actor a discussed by Ishiharaetal.(1981):Dieckmann(2000b).. and we use. Belore discussing the efficiency of electron acceleration. it is better to check whether or not the above estimation is reasonable.," Although the nonlinear saturation process of BI still remains controversial, we adapt here the conversion factor $\alpha$ discussed by \citet{Ish81,Die00b}, and we use, Before discussing the efficiency of electron acceleration, it is better to check whether or not the above estimation is reasonable."
 We first analvze (he magnetospheric observation data for a high Mach number shock., We first analyze the magnetospheric observation data for a high Mach number shock.
" From the satellite observations in the Earth's bow shock. we know thal Z4.=50~200mV/m (Matsumotoetal.1997:Dale1993).. while the solar wind. motional electric field £j is a few mV/m. Then. the ratio ofE,/Eg is of order of 10*~10?."," From the satellite observations in the Earth's bow shock, we know that $E_{\rm esw} = 50 \sim 200~{\rm mV/m}$ \citep{Mat97,Bal98}, while the solar wind, motional electric field $E_0$ is a few mV/m. Then, the ratio of$E_{\rm esw}/E_0$ is of order of $10^1 \sim 10^2$."
" On the other hand. by virtue of Eq.(2)). the ratio of Ei to the motional electric field £i,=vpBo/e in upstream can be obtained. Since the Alfvénn velocity V4 is about 300 km/s in the solar wind. the theoretical estimation of £L;£y becomes 13~ 23."," On the other hand, by virtue of \ref{eq:eswamp}) ), the ratio of $E_{\rm esw}$ to the motional electric field $E_0 = v_0 B_0/c$ in upstream can be obtained, Since the Alfvénn velocity $V_A$ is about 300 km/s in the solar wind, the theoretical estimation of $E_{\rm esw}/E_0$ becomes $13 \sim 23$ ."
 This value is remarkably close to the observation., This value is remarkably close to the observation.
"values of the free parameters of the jet model chosen, =25 and e,=0.029 (Figure 1)), e.=0.01 (Figure 2)) rj;/rs;reproduce the data well for Lj,=0.15L (Figure 1)), Ly=0.2 2)), and weak magnetic field.","values of the free parameters of the jet model chosen, $r_j/r_s = 25$ and $\epsilon_e = 0.029$ (Figure \ref{fig:1}) ), $\epsilon_e = 0.01$ (Figure \ref{fig:2}) ) reproduce the data well for $L_k = 0.15 L_E$ (Figure \ref{fig:1}) ), $L_k = 0.2$ (Figure \ref{fig:2}) ), and weak magnetic field."
" Here we chose B(Figure=10? G, which implies matter dominated outflow - up/uxyc3x1075."," Here we chose $B =
10^3$ G, which implies matter dominated outflow - $u_B/u_{k} \simeq 3
\times 10^{-8}$."
" As explained above, the exact value of the magnetic field is unimportant as long as the constraints in equation 3.2 are fulfilled."," As explained above, the exact value of the magnetic field is unimportant as long as the constraints in equation \ref{eq:constraints2} are fulfilled."
" In this scenario as well, for £jj=1 a power law index p=2.0 results in very good fit to the data."," In this scenario as well, for $\xi_{pl} = 1$ a power law index $p=2.0$ results in very good fit to the data."
" In this paper, we present a new jet-dominated model that is able to reproduce the main spectral properties seen at the X-rays of many XRBs in the hard state."," In this paper, we present a new jet-dominated model that is able to reproduce the main spectral properties seen at the X-rays of many XRBs in the hard state."
" Our key motivation is the spectral break which is often seen at ~ few - few tens,, and (within these sources) the nearly universal spectral slope, Fyοςv1/? observed below this break."," Our key motivation is the spectral break which is often seen at $\sim$ few - few tens, and (within these sources) the nearly universal spectral slope, $F_\nu \propto \nu
^{1/2}$ observed below this break."
" This spectral slope is a natural outcome of emission from electrons whose energy distribution is determined by rapid radiative cooling, following acceleration at the jet base."," This spectral slope is a natural outcome of emission from electrons whose energy distribution is determined by rapid radiative cooling, following acceleration at the jet base."
" The rapid cooling can result from either synchrotron emission or Comptonization of the disk photons, or as likely, a contribution from both processes."," The rapid cooling can result from either synchrotron emission or Comptonization of the disk photons, or as likely, a contribution from both processes."
" We derive in equation 3.1, 9.2. the required constraints on the free model parameters for both these scenarios, and demonstrate the resulting spectrum in Figures 1 and 2.."," We derive in equation \ref{eq:constraints1}, \ref{eq:constraints2} the required constraints on the free model parameters for both these scenarios, and demonstrate the resulting spectrum in Figures \ref{fig:1} and \ref{fig:2}."
" 'The key difference between our model and earlier jet models is that here we self-consistently consider the temporal variation of the particle distribution due to the radiative cooling, and consider those variations in the spectral calculations."," The key difference between our model and earlier jet models is that here we self-consistently consider the temporal variation of the particle distribution due to the radiative cooling, and consider those variations in the spectral calculations."
" Thus, we are able to make 8 clear separation between the acceleration process and the cooling processes; namely, we do not assume that the energy lost by the electrons is necessarily fully replenished by any heating source."," Thus, we are able to make a clear separation between the acceleration process and the cooling processes; namely, we do not assume that the energy lost by the electrons is necessarily fully replenished by any heating source."
" Indeed, in recent years it has become clear that this decoupling between acceleration and cooling may hold the key to understanding the spectral properties of XRBs."," Indeed, in recent years it has become clear that this decoupling between acceleration and cooling may hold the key to understanding the spectral properties of XRBs."
" As a result, in recent years several numerical models which consider this separation have been constructed in the study of XRBs spectra (e.g.,Belmontetal.2008;Maitraetal. 2009)."," As a result, in recent years several numerical models which consider this separation have been constructed in the study of XRBs spectra \citep[e.g.,][]{BM08, MMF09}."
". In essence, these models are very similar to the numerical model used here."," In essence, these models are very similar to the numerical model used here."
" A key advantage of our model is its ability of solving the kinetic equations over many orders of magnitude in time and energy scales, which makes it ideal for producing broadband spectra."," A key advantage of our model is its ability of solving the kinetic equations over many orders of magnitude in time and energy scales, which makes it ideal for producing broadband spectra."
" As indicated in Figures 1 and 2,, both synchrotron- and Compton-dominated scenarios can in principle reproduce the X-ray spectrum of a typical source showing the predicted break."," As indicated in Figures \ref{fig:1} and \ref{fig:2}, both synchrotron- and Compton-dominated scenarios can in principle reproduce the X-ray spectrum of a typical source showing the predicted break."
" The values of the free model parameters are, however, significantly different in the two scenarios."," The values of the free model parameters are, however, significantly different in the two scenarios."
" Whilethe Compton dominated scenario requires r/r,<1015, LycLg and a relatively weak magnetic field, the synchrotron dominated scenario better fits the data with r/r,= 103, LjXi107?Lg and a strong magnetic field, B>109 G. In our opinion, at least for XTE J1118+480, the data favor a synchrotron-dominated scenario, similar to the conclusions of Markoffetal.(2001);Maitra(2009)."," Whilethe Compton dominated scenario requires $r/r_s \lesssim 10^{1.5}$, $L_k \simeq L_E$ and a relatively weak magnetic field, the synchrotron dominated scenario better fits the data with $r/r_s \gtrsim 10^2$ , $L_k \lesssim 10^{-2} L_E$ and a strong magnetic field, $B\gtrsim 10^6$ G. In our opinion, at least for XTE J1118+480, the data favor a synchrotron-dominated scenario, similar to the conclusions of \cite{MFF01,Maitra+09}."
". First, the characteristic scaling is similar to the inner (truncated) disk radius, rin."," First, the characteristic scaling $r_j$ is similar to the inner (truncated) disk radius, $r_{in}$."
" r;Thus, obviously, this radius marks a physical transition in the properties of the flow, which can result from development and launching of jets."," Thus, obviously, this radius marks a physical transition in the properties of the flow, which can result from development and launching of jets."
" Second, the non-thermal radio spectra observed in the hard states provide clear indication for jet-synchrotron emission."," Second, the non-thermal radio spectra observed in the hard states provide clear indication for jet-synchrotron emission."
 Emission at the radio band is expected once the plasma reaches a scale much larger than r; (and with corresponding much weaker magnetic field)., Emission at the radio band is expected once the plasma reaches a scale much larger than $r_j$ (and with corresponding much weaker magnetic field).
" As the magnetic field is expected to decay along the jet (due, e.g., to Poynting-flux conservation), unless there is a generation of strong magnetic fields at some scale r> rj, astrong magnetic field at the jet base decays along the jet, and thus the same field may be the source of the radio"," As the magnetic field is expected to decay along the jet (due, e.g., to Poynting-flux conservation), unless there is a generation of strong magnetic fields at some scale $r\gg r_j$ , astrong magnetic field at the jet base decays along the jet, and thus the same field may be the source of the radio"
2 the brightest sources are always clustered relatively nearby.,2 the brightest sources are always clustered relatively nearby.
 Η scenario 2 is adopted. a lack of events from a nearby collection of matter can be used to argue against a model with relatively high source densities (e.g.. see Gorbunovetal.(2003))). Wasian(1995) d (Allardetal.2007)))," If scenario 2 is adopted, a lack of events from a nearby collection of matter can be used to argue against a model with relatively high source densities (e.g., see \citet{Gorbunov}) \citet{Waxman} \ref{fig:Qbar} \citep{Allard})"
 Η scenario 2 is adopted. a lack of events from a nearby collection of matter can be used to argue against a model with relatively high source densities (e.g.. see Gorbunovetal.(2003))). Wasian(1995) d (Allardetal.2007))).," If scenario 2 is adopted, a lack of events from a nearby collection of matter can be used to argue against a model with relatively high source densities (e.g., see \citet{Gorbunov}) \citet{Waxman} \ref{fig:Qbar} \citep{Allard})"
General relativity predicts that the wavelength οἱ electromagnetic radiation ds sensitive το gravitationa potentials. an effect which is calledredshift.,"General relativity predicts that the wavelength of electromagnetic radiation is sensitive to gravitational potentials, an effect which is called."
 Photous ravelliug from the surface of ast scattering wil necessarily travel through the eravitational potential of Large Scale Structure (LSS) on their wav to the observer: these wil be bluc-shitted as they cuter the potential wel and red-shifted as they exit the potential., Photons travelling from the surface of last scattering will necessarily travel through the gravitational potential of Large Scale Structure (LSS) on their way to the observer; these will be blue-shifted as they enter the potential well and red-shifted as they exit the potential.
 These shifts will accumulate along the line of sight of the observer., These shifts will accumulate along the line of sight of the observer.
 The tota shift in wavoleugth will translate iuto a chanec in the measured temiperature-teniperature anisotropy of the CAIB. and can be calculated by: where T is the temperature ofthe CAIB. iis the couformal iue. defined by dj=alt)dr and iy auc ap represeut the conforma times today aud at the surface of las scattering respectively.," The total shift in wavelength will translate into a change in the measured temperature-temperature anisotropy of the CMB, and can be calculated by: where $T$ is the temperature of the CMB, $\eta$ is the conformal time, defined by $\rd \eta = \frac{\rd t}{a(t)}$ and $\eta_0$ and $\eta_L$ represent the conformal times today and at the surface of last scattering respectively."
 The uit vector n is along the Lue of sieht aud he gravitational poteutial B(x.η) epends on position aud time.," The unit vector $\hat{\bf n}$ is along the line of sight and the gravitational potential $\Phi({\bf x}, \eta)$ depends on position and time."
 The iuteeral depends ou the rate of change of he potential ®!2d@/dy., The integral depends on the rate of change of the potential $\Phi'=\rd \Phi / \rd\eta$.
 Iu a universe with no dark οπσον or curvature. the cosunic (linear) eravitational potential docs no varv with nue. so that such a blue- aud τοςπμ will always cancel out. because $=0 aud there will be no net effect on the wavelength of the photon.," In a universe with no dark energy or curvature, the cosmic (linear) gravitational potential does not vary with time, so that such a blue- and red-shift will always cancel out, because $\Phi' = 0$ and there will be no net effect on the wavelength of the photon."
 Ilowever. in the presence of dark enerev or curvature νε the right haud side of Equation 1. will be as the cosmic potential will chauge with time7). resulting im a secondary auisotropy in the CMD temperature field.," However, in the presence of dark energy or curvature , the right hand side of Equation \ref{sec:theory:eq:isw} will be non-null as the cosmic potential will change with time, resulting in a secondary anisotropy in the CMB temperature field."
 The ISW effect leads to a linear scale secondary anisotropy iu the temperature field of the CAIB. and will thus affect the CAIB temperature power spectruui at laree scales.," The ISW effect leads to a linear scale secondary anisotropy in the temperature field of the CMB, and will thus affect the CMB temperature power spectrum at large scales."
 Due to the primordial anisotropies aud cosmic variance on large scales. the ISW signal is dificult to detect directly in the temperature map of he CAIB. but showed it could be detected through cross-correlation of the CMD with a local tracer of mass.," Due to the primordial anisotropies and cosmic variance on large scales, the ISW signal is difficult to detect directly in the temperature map of the CMB, but showed it could be detected through cross-correlation of the CMB with a local tracer of mass."
 The first attempt to detect the ISW effec involved correlating the Cosmic Microwave Backeround explorer data with NRB aud NVSS data(?)., The first attempt to detect the ISW effect involved correlating the Cosmic Microwave Background explorer data with XRB and NVSS data.
. This analvsis did not find a significant correlation between the local tracers of mass aud the CAIB., This analysis did not find a significant correlation between the local tracers of mass and the CMB.
 Since the release of data from the Wilkinson Microwave Anisotropy Probe over 20 studies (see Table 1)) have investigated cross-correlations between the different. vears of WALAP data and local tracers selected. using various wavelengths: N-ravsurvey): opticalealaxies)..QSOx)..LRCis).. APAD):: near infrared2\TASS):: radioNVSS).," Since the release of data from the Wilkinson Microwave Anisotropy Probe over $20$ studies (see Table \ref{sec:theory:tab:detections}) ) have investigated cross-correlations between the different years of WMAP data and local tracers selected using various wavelengths: X-ray; optical, ; near infrared; radio."
. The full sky WALAP data have sufficient resolution on large scales that the measure of the ISW sieual is cosmic variance limited., The full sky WMAP data have sufficient resolution on large scales that the measure of the ISW signal is cosmic variance limited.
 The best LSS probe of the ISW effect should include maxiuun sky coverage and full redshift coverage of the dark energy dominated era(?)., The best LSS probe of the ISW effect should include maximum sky coverage and full redshift coverage of the dark energy dominated era.
. No such survey exists vet. so there is room for duprovement on the ISW cletecueon as larger and larger LSS surveys arise;," No such survey exists yet, so there is room for improvement on the ISW detection as larger and larger LSS surveys arise."
 For lis reason. when we review the current {ον detections. we classify thei according to their tracer of LSS. aud not he CMD map used.," For this reason, when we review the current ISW detections, we classify them according to their tracer of LSS, and not the CMB map used."
 The iicasure of the ISW signal can be done iu various statistical spaces: we classify detections in Table 1. iuto liree measurement ‘domains’: D1 correspouds to spherical i;unmonic space: D2 to configuration space and D3 to wavelet space. (, The measure of the ISW signal can be done in various statistical spaces; we classify detections in Table \ref{sec:theory:tab:detections} into three measurement `domains': D1 corresponds to spherical harmonic space; D2 to configuration space and D3 to wavelet space. (
Iu. Section ??.. we review the different ucthods for quantitving the statistical significance of cach neasurelucut1.,"In Section \ref{sec:method}, we review the different methods for quantifying the statistical significance of each measurement)."
 There are only two analyses which use CODE as CMD data?).. and both report uull detections. which can reasonably be due to the low angular resolution of CODE even at large scales.," There are only two analyses which use COBE as CMB data, and both report null detections, which can reasonably be due to the low angular resolution of COBE even at large scales."
 The rest are done correlating WALAP data from vears 1. 3 and 5 (respecively“WI. CWO and W5 in table 1)).," The rest are done correlating WMAP data from years 1, 3 and 5 (respectively`W1', `W3' and `W5' in table \ref{sec:theory:tab:detections}) )."
 Most ISW detections reported in Table 1. are relatively weak (< 30) and this is expected from theory for a concordance cosmology., Most ISW detections reported in Table \ref{sec:theory:tab:detections} are relatively `weak' $<3\sigma$ ) and this is expected from theory for a concordance cosmology.
 Higher detections are reported or the NVSS survey(??7).. though weal: aud. mareinal detections using NWSS data are also reported(??).," Higher detections are reported for the NVSS survey, though weak and marginal detections using NVSS data are also reported."
. Theh detections are often made using a wavelet analysis).. hough a simular study by the same authors usine the sale data but a different analysis method finds a weaker signalOo(?7).," High detections are often made using a wavelet analysis, though a similar study by the same authors using the same data but a different analysis method finds a weaker signal."
. The highest detection is reported usingC» a olographic combination of all smrveys?).. as expected eiven the larger redshift coverage of tle analysis.," The highest detection is reported using a tomographic combination of all surveys, as expected given the larger redshift coverage of the analysis."
 Several analyses have heen revisited το. seek confirmation of previous detections., Several analyses have been revisited to seek confirmation of previous detections.
 ιν some cases results are very snadlu(???).. for SDSS LRGs: for SDSS Quasars: ??).. for 2ALASS). but in some cases they are controversially differeut (for e.g. aud?).. for NVSS or and2).. foy 2\LASS).," In some cases, results are very similar, for SDSS LRGs; for SDSS Quasars; , for 2MASS), but in some cases they are controversially different (for e.g. and, for NVSS or and, for 2MASS)."
 We also notice that as certain surveys are revisited. there is a trend for the statistical siguificance to be reduced: for c.g.. detections from 2MASS decrease frou a 2.59 detection(2).. to 2a(?).. to (.5o to weak(?).," We also notice that as certain surveys are revisited, there is a trend for the statistical significance to be reduced: for e.g., detections from 2MASS decrease from a $2.5\sigma$ detection, to $2\sigma$, to $0.5\sigma$ to `weak'."
. Detections using SDSS LRGs decrease from 2.56(?). to 2.2.20(?7).. to aunareiual. (2).," Detections using SDSS LRGs decrease from $2.5\sigma$, to $2-2.2\sigma$, to `marginal' ."
. Furthermore. there tends to be a ‘sociological bias in the interpretation of the confidence on the signal detection.," Furthermore, there tends to be a `sociological bias' in the interpretation of the confidence on the signal detection."
 The first detections interpret à 230 detection as ‘tentative’ ).. while further studies with similar detectionlevel report independent evidence of dark energy!(??)..," The first detections interpret a $2-3 \sigma$ detection as `tentative' , while further studies with similar detectionlevel report `independent evidence of dark energy'."
when the components merge.,when the components merge.
 On the other hand. if the sdB star is truly a core-Lle burning ELIB star. its mass is expected to be close to 0.5Ανν and with q derived. from. the light curve solution constraining the mass of the unseen white cart. the total mass may be considerably less than Mou.," On the other hand, if the sdB star is truly a core-He burning EHB star, its mass is expected to be close to $0.5\,M_{\odot}$, and with $q$ derived from the light curve solution constraining the mass of the unseen white dwarf, the total mass may be considerably less than $M_{\rm Ch}$."
 Improved light curves ancl radial velocity cata will help to address this question., Improved light curves and radial velocity data will help to address this question.
 In the meantime. it is well to bear in mind the discussion in Jellerv Simon (1997) concerning whether the canonical 0.5AL. for ELLB stars is really empirically svell-cdetermined.," In the meantime, it is well to bear in mind the discussion in Jeffery Simon (1997) concerning whether the canonical $0.5\,M_\odot$ for EHB stars is really empirically well-determined."
" ""Third. short-perioc binary systems of the same type as KPD 0422|5421. namely sdD. |. white chart. are rare. and WD 0422|5421 arguably possesses the shortest period of them all. "," Third, short-period binary systems of the same type as KPD 0422+5421, namely sdB + white dwarf, are rare, and KPD 0422+5421 arguably possesses the shortest period of them all. ["
One double BM system has a shorter »eriod: WD 0957-666. he22-=0.060993 davs. (Moran. Marsh. DBragaglia 1997),"One double degenerate system has a shorter period: WD 0957-666, $P = 0.060993$ days, (Moran, Marsh, Bragaglia 1997)."
 mass is estimated to be only 934.. however]," The total mass is estimated to be only $0.69\,M_{\odot}$, however.]"
 Phe catalog of Ritter Ixolb. (1998) ists three possible sd|7white cwarl” pairs ssdB stars known to be in close binaries with undetected companions). and the sample of Saller. Livio. Yungelson (1998) contains seven sdD binaries with undetected. companions (including wo from Ritter Ixolb's catalog). of which five have known orbital periods.," The catalog of Ritter Kolb (1998) lists three possible sdB+“white dwarf” pairs sdB stars known to be in close binaries with undetected companions), and the sample of Saffer, Livio, Yungelson (1998) contains seven sdB binaries with undetected companions (including two from Ritter Kolb's catalog), of which five have known orbital periods."
 The six objects with known orbital periods (in order ofincreasing peare PG 1432|159(P=5.39 jours). PG. 2345|318 (PP?rio)=5.78 hours). Feige 36(2=8.5 jours). PG as0101]039 (2=13.7 hours). LIZ22 (ZUX CVn. P=13.77 bis and Ton 245 (2=2.5 days).," The six objects with known orbital periods (in order of increasing period)are PG 1432+159$P=5.39$ hours), PG 2345+318 $P=5.78$ hours), Feige 36$P=8.5$ hours), PG 0101+039 $P=13.7$ hours), HZ 22 (=UX CVn, $P=13.77$ hours), and Ton 245 $P=2.5$ days)."
 A related sdO | deal, A related sdO + white dwarf(?)
 system is LID 49798 with 2=1.55 days (Phackeray 1970)., system is HD 49798 with $P=1.55$ days (Thackeray 1970).
 The best known member of the κα | white ciwarf class is LIZ 22 (Llumason Zwicky 1947: Young. Nelson. Miclbrecht 1072: Greenstein 1973)," The best known member of the sdB + white dwarf class is HZ 22 (Humason Zwicky 1947; Young, Nelson, Mielbrecht 1972; Greenstein 1973)."
 CGreenstein. (L973) anc Young Wentworth (1982) argue that the unseen companion of LIZ 22 is a white dwarf., Greenstein (1973) and Young Wentworth (1982) argue that the unseen companion of HZ 22 is a white dwarf.
 However. it is also possible that LIZ 22 contains a low mass main sequence star Companion. so the nature of LZ 22 is still an open question.," However, it is also possible that HZ 22 contains a low mass main sequence star companion, so the nature of HZ 22 is still an open question."
 Another recent sdB|white cdwarl” candidate is V46 in the elobular cluster MA (Ixaluzny. Thompson WkKrezeminsky 1997). which is an sdB star that shows an apparently. sinusoidal light curve with a period of 1.045 hours.," Another recent sdB+“white dwarf” candidate is V46 in the globular cluster M4 (Kaluzny, Thompson Krzeminsky 1997), which is an sdB star that shows an apparently sinusoidal light curve with a period of 1.045 hours."
 This period is much longer than those seen in the pulsating sdB stars (eC14026 stars. Wilkenny et al.," This period is much longer than those seen in the pulsating sdB stars (EC14026 stars, Kilkenny et al."
 1997a)., 1997a).
 However. the nature of the Companion star is not known and it is not clear if the photometric period is the orbital period.," However, the nature of the companion star is not known and it is not clear if the photometric period is the orbital period."
 Since it has the shortest. confirmed. orbital period. KWPD 042215421. will evolve. more quickly gravitational wave radiation of angular momentum than the other πο white dwarl systems.," Since it has the shortest confirmed orbital period, KPD 0422+5421 will evolve more quickly gravitational wave radiation of angular momentum than the other sdB + white dwarf systems."
 Ritter (1986) gives the time required for a detached binary to reach the semi-detachecl state: where the masses are in solar units. the periods are in days. and where which assumes that the radius of the πα] star stays ixed.," Ritter (1986) gives the time required for a detached binary to reach the semi-detached state: where the masses are in solar units, the periods are in days, and where which assumes that the radius of the sdB star stays fixed."
 Using 2.)/2=0.57£018. we find that logefy=S.1T2E0.15 πο107 vears). which is comparable to the core Lle burning lifetime of an sdB star (1.5.«107 vears. Dorman. Rood. O'Connell. 1993).," Using $P_{\rm sd}/P=0.57\pm 0.13$, we find that $\log t_{\rm sd}=8.17\pm
0.15$ $t_{\rm sd}\approx 1.5\times 10^8$ years), which is comparable to the core He burning lifetime of an sdB star $\approx1.5\times
10^8$ years, Dorman, Rood, O'Connell 1993)."
 Thus. possibly. before he sdD star has evolved to a white ecwart. there will be a urther episode of mass exchange.," Thus, possibly before the sdB star has evolved to a white dwarf, there will be a further episode of mass exchange."
 What IKPD 0422|5421 will then look like. and whether it can be identified. with any presently known class of objects. takes the subject well »vond the expertise or inclination of the present authors to speculate.," What KPD 0422+5421 will then look like, and whether it can be identified with any presently known class of objects, takes the subject well beyond the expertise or inclination of the present authors to speculate."
 The κα star WD 0422|5421 was found to be a short period detached: binary with a period of £2=0.0901795+(3.10') days., The sdB star KPD 0422+5421 was found to be a short period detached binary with a period of $P=0.0901795\pm (3\times 10^{-7})$ days.
 We argue that the companion star is a white dwarl since the companion star is not seen in the spectra. and the synthetic light curve model of the Ü and D light curves requires its radius to be on the order of =0.0172..," We argue that the companion star is a white dwarf since the companion star is not seen in the spectra, and the synthetic light curve model of the $U$ and $B$ light curves requires its radius to be on the order of $\approx 0.01\,R_{\odot}$."
 We derive component masses of Aa=0.7220.26M. and AMoaup=0.622018M.," We derive component masses of $M_{\rm
sdB}=0.72\pm 0.26\,M_{\odot}$ and $M_{\rm comp}= 0.62\pm
0.18\,M_{\odot}$."
 NPD 0422|5421 is one of a small number of known sdD. | white cwarf binary systems. and represents a poorly observed and short-Iived stage of binary star evolution., KPD 0422+5421 is one of a small number of known sdB + white dwarf binary systems and represents a poorly observed and short-lived stage of binary star evolution.
 The authors ave grateful to τον Salfer for supplying his gravity and temperature estimate of WD 0422|5421 before publication., The authors are grateful to Rex Saffer for supplying his gravity and temperature estimate of KPD 0422+5421 before publication.
 We also thank Peter Egeleton for pointing out to us the intriguing ancl possibly. related: system. HD 185510., We also thank Peter Eggleton for pointing out to us the intriguing and possibly related system HD 185510.
 CIx thanks Ed ther for his hospitality at the University of Texas. (Austin): the director of the McDonald Observatory for a generous allotment oftelescope time: Rob Robinson and Rac Stiening for the use of the Stiening photometer: and the South African Foundation for Research Development for partially funding some of this work., CK thanks Ed Nather for his hospitality at the University of Texas (Austin); the director of the McDonald Observatory for a generous allotment of telescope time; Rob Robinson and Rae Stiening for the use of the Stiening photometer; and the South African Foundation for Research Development for partially funding some of this work.
 JO and RW acknowledge partial support from NASA erant NAG 5-3459 to the Pennsylvania State University., JO and RW acknowledge partial support from NASA grant NAG 5-3459 to the Pennsylvania State University.
 We are grateful to Ivan Llubeny for instruction in the use of his stellar atmosphere codes TLUSTY. SYNSPEC. and. ROTLENS.," We are grateful to Ivan Hubeny for instruction in the use of his stellar atmosphere codes , , and ."
 This research. has 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."
scales. analogous to the finding of Einastoctal.(2011) that the component superclusters of the Sloan Great Wall have cillerent evolutionary histories.,"scales, analogous to the finding of \citet{Einasto2011} that the component superclusters of the Sloan Great Wall have different evolutionary histories."
 Even if this were true. homogeneity asserts that any. global property of sulficientlv large volumes should. be the same within the expected statistical variations. so the density of quasars in these LOCs would remain clistinetive.," Even if this were true, homogeneity asserts that any global property of sufficiently large volumes should be the same within the expected statistical variations, so the density of quasars in these LQGs would remain distinctive."
 Of course. unknown observational biases or selection elfecets could conceivably also allect. the overall dimensions of the LOCs.," Of course, unknown observational biases or selection effects could conceivably also affect the overall dimensions of the LQGs."
 In finding compatibility. of LOQGs with concordance cosmology Milleretal.(2004) noted. that they hack not considered questions of shape and topology., In finding compatibility of LQGs with concordance cosmology \citet{Miller2004} noted that they had not considered questions of shape and topology.
 Given both the large sizes and elongated morphology that we find for Ul.28 and Ul.11 we have begun a programme to re-investigate compatibility. using the full set of LOCs that we find from the DRTQSO catalogue.," Given both the large sizes and elongated morphology that we find for U1.28 and U1.11 we have begun a programme to re-investigate compatibility, using the full set of LQGs that we find from the DR7QSO catalogue."
 LEC received. partial support [rom Center of Excellence in Astrophysics and Associated Technologies (PED 06)., LEC received partial support from Center of Excellence in Astrophysics and Associated Technologies (PFB 06).
 The many authors of GNU/Linux. IH. €GGobi. and TOPCAT are gratefully acknowledged.," The many authors of GNU/Linux, R, GGobi and TOPCAT are gratefully acknowledged."
 This research. has used the SDSS DIUQSO catalogue (Schneiderοἱal.2010)., This research has used the SDSS DR7QSO catalogue \citep{Schneider2010}.
. Funding for the SDSS and SDSS-LE has been provided bv the Alfred P. Sloan. Foundation. the Participating Institutions. the National Science. Foundation. the U.S.," Funding for the SDSS and SDSS-II has been provided by the Alfred P. Sloan Foundation, the Participating Institutions, the National Science Foundation, the U.S."
After finishing the final draft of this manuscript we found out that Chadidetal.(2011). had just. published. the results of their analysis of the same public data set. of ColtoT 105288363.,After finishing the final draft of this manuscript we found out that \citet{cha11} had just published the results of their analysis of the same public data set of CoRoT 105288363.
 While in some points. our results are in agreement with theirs. there are also some aspects which differ.," While in some points, our results are in agreement with theirs, there are also some aspects which differ."
 Those aspects refer to the following items. cliscussecl in a detailed wav in the text.," Those aspects refer to the following items, discussed in a detailed way in the text."
 Chacidetal.(2011). state that the behaviour of CoRGL 105288363 is a strong support for the scenario described by Stothers(2006)., \citet{cha11} state that the behaviour of CoRoT 105288363 is a strong support for the scenario described by \citet{sto}.
. We note. however. that a varving magnetic field modulating the convection is still an unproven hypothesis. anc that there is currently no model that can quantitatively describe the interaction between such a field and the turbulent 3achler&Iollath(2011). use the amplitude equation formalism {ο study the interaction between the fundamental mode and the 9th. overtone in a 9:2 resonance. and find stochastic. chaotic behaviour.," We note, however, that a varying magnetic field modulating the convection is still an unproven hypothesis, and that there is currently no model that can quantitatively describe the interaction between such a field and the turbulent \citet{buchler11} use the amplitude equation formalism to study the interaction between the fundamental mode and the 9th overtone in a 9:2 resonance, and find stochastic, chaotic behaviour."
 The low aniplitucles of possible excited: overtones and. the fact that period. doubling is not observed in all Blazhko stus. however. seem to make the 9:2 resonance as sole cause for the Blazhko elect unlikely.," The low amplitudes of possible excited overtones and the fact that period doubling is not observed in all Blazhko stars, however, seem to make the 9:2 resonance as sole cause for the Blazhko effect unlikely."
 On the other hand. the resonance that was found in the simulations of Szabóetal...(2010). and. Buchler&IWwolléth(2011) was startlingly strong. ancl considering the fact that. period. doubling is à temporary phenomenon which is present only during certain phases. one might expect that the occurence rate increases when more continuous and accurate satellite data become Lt is. of course. possible that the two elfects. the variable turbulent convection as well as the 9:2 resonance. coexist to explain the various observed. phenomena.," On the other hand, the resonance that was found in the simulations of \citet{szabo10} and \citet{buchler11} was startlingly strong, and considering the fact that period doubling is a temporary phenomenon which is present only during certain phases, one might expect that the occurence rate increases when more continuous and accurate satellite data become It is, of course, possible that the two effects, the variable turbulent convection as well as the 9:2 resonance, coexist to explain the various observed phenomena."
 The observed behaviour of ColtoT. 105288363 certainly. poses a significant problem to all models which require a clock-work like behaviour. and favours explanations that are capable of producing irregular Chadidetal.(2011). report a long-term: modulation periotof 151 4 7 d. We note that the detection of periods exceeding the time base of the availabe data (145 d)l has to be taken with great. caution.," The observed behaviour of CoRoT 105288363 certainly poses a significant problem to all models which require a clock-work like behaviour, and favours explanations that are capable of producing irregular \citet{cha11} report a long-term modulation period of 151 $\pm$ 7 d. We note that the detection of periods exceeding the time base of the availabe data (145 d) has to be taken with great caution."
" Frequencies which are separated by a value of 17,yocl/NT (Loumos&DecningLOTS)..."," Frequencies which are separated by a value of $\nu_{1}-\nu_{2} < 1/\Delta T $ \citep{lou78}. \ref{lightcurve},"
 Ll. 5. 7. ασιαetal.(2011)..., \ref{fp} \ref{loop} \citet{cha11}.
While the validity of this scaling law applies to a coronal loop in hvdrostatic equilibrium. 1 mieht also approximately apply (o a flaring loop near the peak time. because both (1) the enerev and (2) momentum equations are nearly balanced near the flare peak. (,"While the validity of this scaling law applies to a coronal loop in hydrostatic equilibrium, it might also approximately apply to a flaring loop near the peak time, because both (1) the energy and (2) momentum equations are nearly balanced near the flare peak. ("
1) Energv equation: In the initial heating phase. the heating term is larger than the combined conductive and radiative loss term. ie. 75;>(I44E44). because the average flare temperature is rising (41(/)/dl> 0) due to excessive heating.,"1) Energy equation: In the initial heating phase, the heating term is larger than the combined conductive and radiative loss term, i.e., $E_H > (E_{cond}+E_{rad})$, because the average flare temperature is rising $dT(t)/dt > 0$ ) due to excessive heating."
 After the peak time. the conductive and radiative losses exceed (he decreasing or stopped heating rate. ie. Brag). because the flare plasma is cooling (dP(αικ0).," After the peak time, the conductive and radiative losses exceed the decreasing or stopped heating rate, i.e., $E_H < (E_{cond}+E_{rad})$ , because the flare plasma is cooling $dT(t)/dt < 0$ )."
" Therelore. (here is a balance between heating and loss terms. i.e. Ly£2(Leong4E,,4). near the flare peak time I, [defined by the peak in emission measure £M,ο= f,)]."," Therefore, there is a balance between heating and loss terms, i.e., $E_H \approx (E_{cond}+E_{rad})$ , near the flare peak time $t_p$ [defined by the peak in emission measure $EM_p=EM(t=t_p$ )]."
 Hydrodynanmic simulations bv Jakimiec et al. (, Hydrodynamic simulations by Jakimiec et al. (
"1992) show that the RTV scaling law predicts a maximunm electron density ny that can be reached in a flare loop if a constant heating rate is applied sufficiently long alter the maximum temperature 7, is reached.",1992) show that the RTV scaling law predicts a maximum electron density $n_p$ that can be reached in a flare loop if a constant heating rate is applied sufficiently long after the maximum temperature $T_p$ is reached.
" In Appendix A we show how the agreement between the maximum electron density n, and the RTV-precdicted value scales with the heating duration. but is nearly independent of the maximum temperature 7, and heating rale Ey⋅ ⋖⋡∃"," In Appendix A we show how the agreement between the maximum electron density $n_p$ and the RTV-predicted value scales with the heating duration, but is nearly independent of the maximum temperature $T_p$ and heating rate $E_H$. ("
↕⋝⇀∖↕∪∐∐↲∐⊓∐∐≼↲≺⇂∏≀↧↴∐∪∐∶⊳↔⊲≼↲≺∢∪∐≼∐⋡∖↽⋅≀↧↴↥⋟∖⊽∪⊔∐↲∐↓∪∐∐↲∐⊓∐∐≼↲≺⇂∏≀↧↴∐∪∐↕⋟∖⊽∐≼↲≀↧↴∏⋡∖↽∣↽≻≀↕↴↥≀↧↴∐≺∢≼↲≼⇂ alter a flare loop is filled.,"2) Momentum equation: Secondly, also the momentum equation is nearly balanced after a flare loop is filled."
 The losses are dominated by thermal conduction at high flare temperatures (sav at J2LO MIN: Asclivanden Alexander 2001). and the loop filling time is much shorter Gn the order of της1 min: see hydrodynamic simulations by MacNeice οἱ al.," The losses are dominated by thermal conduction at high flare temperatures (say at $T\gapprox 10$ MK; Aschwanden Alexander 2001), and the loop filling time is much shorter (in the order of $\tau_{fill}
\lapprox 1$ min; see hydrodynamic simulations by MacNeice et al."
 1984: Nagai Eimslie 1984: Fisher et al., 1984; Nagai Emslie 1984; Fisher et al.
 1985a.b.c; Mariska Polancl 1935: Yokovama Shibata 1998. 2001: Hori et al.," 1985a,b,c; Mariska Poland 1985; Yokoyama Shibata 1998, 2001; Hori et al."
 1997. 1998) than the conductive or radiative cooling time (in the order of 7.47Z10 min: Antiochos Sturrock 1978: Culhane οἱ al.," 1997, 1998) than the conductive or radiative cooling time (in the order of $\tau_{cool}
\gapprox 10$ min; Antiochos Sturrock 1978; Culhane et al."
 1994: Aschwancen Alexander 2001). and thus pressure gradients resulting from the chromospheric evaporation process are largely balanced out so that the momentum equation is approximately [ulfillecd.," 1994; Aschwanden Alexander 2001), and thus pressure gradients resulting from the chromospheric evaporation process are largely balanced out so that the momentum equation is approximately fulfilled."
 Also the assumption of a constant pressure (made in the REV law) is better fulfilled in the hot soft X-ray emitting flare loops (because of the larger pressure scale heights. Ap> L) than in the cooler EUV-emitting coronal loops (where often Ay« £L).," Also the assumption of a constant pressure (made in the RTV law) is better fulfilled in the hot soft X-ray emitting flare loops (because of the larger pressure scale heights, $\lambda_T > L$ ) than in the cooler EUV-emitting coronal loops (where often $\lambda_T < L$ )."
" Applying the standard RTV scaling law to flare loops now. with 7,7ως. bv inserting the expression for the thermal pressure. we obtain a scaling law forthe peak density 7, as a function of the peak temperature T,p aud loop length L."," Applying the standard RTV scaling law to flare loops now, with $T_p \approx T_{max}$, by inserting the expression for the thermal pressure, we obtain a scaling law forthe peak density $n_p$ as a function of the peak temperature $T_p$ and loop length $L$ ,"
Herschel. tracing the dust heated by newly formed stars and the diffuse interstellar radiation field. and (111) CO and eemission tracing the neutral molecular and atomic gas.,"Herschel, tracing the dust heated by newly formed stars and the diffuse interstellar radiation field, and (iii) CO and emission tracing the neutral molecular and atomic gas."
 The rest of the paper is organized as follows: Sec., The rest of the paper is organized as follows: Sec.
 2 presents details of our observations and ancillary data. Sec.," 2 presents details of our observations and ancillary data, Sec."
 3 states the basic results of the spectroscopic observations and a qualitative and a quantitative comparison of the aand eemission with all other available tracers and their correlation., 3 states the basic results of the spectroscopic observations and a qualitative and a quantitative comparison of the and emission with all other available tracers and their correlation.
 Sec., Sec.
 4 studies the role of aas an indicator of the star formation rate (SFR) and Sec., 4 studies the role of as an indicator of the star formation rate (SFR) and Sec.
 5 analyzes the energy balance in the mapped region., 5 analyzes the energy balance in the mapped region.
 Sec., Sec.
 6 presents a detailed analysis of the emission from the ppeak position in 3302 in terms of models of PDRs., 6 presents a detailed analysis of the emission from the peak position in 302 in terms of models of PDRs.
 In Sec., In Sec.
 7 we summarize and discuss the major findings of the paper., 7 we summarize and discuss the major findings of the paper.
 A region extending over 2x2' aaround the rregion BCLMP 302 in the northern. arm of 333. was observed with the 5x5 pixel integral field unit (IFU) of the PACS Spectrometer using the wavelength switching (WS) mode in combination with observations of an emission free off source position outside of the galaxy at RA/Dec (J2000) = 22.58717/30.6404., A region extending over $\times$ around the region BCLMP 302 in the northern arm of 33 was observed with the $\times$ 5 pixel integral field unit (IFU) of the PACS Spectrometer using the wavelength switching (WS) mode in combination with observations of an emission free off source position outside of the galaxy at RA/Dec (J2000) = $^\circ$ $^\circ$.
 The field of view of the IFU is 47747” wwith 9744 pixels (Poglitschetal.2010)., The field of view of the IFU is $\times$ with 4 pixels \citep{poglitsch2010}.
". We used the Ist and 3rd order gratings to observe the um. um. um. um. um. and um Hines. with the shortest possible observing time (1 line repetition. | cycle). and with a reference position at RA = 0130""20:99. Dee= 30738/25744 (2000)."," We used the 1st and 3rd order gratings to observe the $\,\mu$ m, $\,\mu$ m, $\,\mu$ m, $\,\mu$ m, $\,\mu$ m, and $\,\mu$ lines, with the shortest possible observing time (1 line repetition, 1 cycle), and with a reference position at RA = $^{\rm h}$ $^{\rm m}$ 9, Dec = 4 (J2000)."
 The reference position was selected based on aand maps., The reference position was selected based on and maps.
" For all the lines a 3x3 raster was observed on a eegrid with the IFU centered at R.A. = 01 34"" 05:99 18766 (2000) with position angle. P. A.=22."," For all the lines a $\times$ 3 raster was observed on a grid with the IFU centered at R.A. = $^{h}$ $^{m}$ 9 6 (J2000) with position angle, P.A. =."
5°.. The resulting footprintis shown in reffig jiresult.., The resulting footprint is shown in \\ref{fig_ciiresult}.
 TheF WH Mbeamsizeo fthePACS spect ometeris nnearó3 ym aand nnear [58pm ((E.Sturm. priv.," The FWHM beam size of the PACS spectrometer is near $\,\mu$ and near $\,\mu$ (E.Sturm, priv."
 comm.), comm.).
 The lines are unresolved. as the spectral resolution of PACS is. larger than ffor all lines.," The lines are unresolved, as the spectral resolution of PACS is larger than for all lines."
 The observations were performed on January 7. 2010 and the total observing time was hhours for all the 6 lines.," The observations were performed on January 7, 2010 and the total observing time was hours for all the 6 lines."
 The PACS spectra were reduced using HIPE version 3.0 CIB 1452 (Ottetal.2010)., The PACS spectra were reduced using HIPE version 3.0 CIB 1452 \citep{ott2010}.
. The WS data reduction pipeline was custom-made by the NASA Herschel Science. Center (NHSC) helpdesk., The WS data reduction pipeline was custom-made by the NASA Herschel Science Center (NHSC) helpdesk.
 The data were exported to FITS cubes which were later analyzed using internally developed IDL routines to extract the line intensity maps., The data were exported to FITS cubes which were later analyzed using internally developed IDL routines to extract the line intensity maps.
 Using PACS. the uim. um. um. and the um hhave been detected.," Using PACS, the $\,\mu$ m, $\,\mu$ m, $\,\mu$ m, and the $\,\mu$ have been detected."
 The lines of um aand the umm wwere not detected.," The lines of $\,\mu$ and the $\,\mu$ were not detected."
 The peak and Io noise limits of the intensities of theu]..i].. um aand um sspectra. at the ppeak position. are presented in Table l..," The peak and $1\sigma$ noise limits of the intensities of the, $\,\mu$ and $\,\mu$ spectra, at the peak position, are presented in Table \ref{tab_linerms}."
 Fig., Fig.
 3 show the observed PACS spectra at the position of the ppeak position., \ref{fig_pacsspec} show the observed PACS spectra at the position of the peak position.
 The peak integrated intensities were derived by first fitting and subtracting a polynomial baseline of 2nd order and then fitting a Gaussian., The peak integrated intensities were derived by first fitting and subtracting a polynomial baseline of 2nd order and then fitting a Gaussian.
 Due to unequal coverage of different positions. as seen from the grid of observed positions shown in reftigiiresult. thermsachievedisnotuniform.," Due to unequal coverage of different positions, as seen from the grid of observed positions shown in \\ref{fig_ciiresult}, the rms achieved is not uniform."
Ivariesbyaboutafactorof., It varies by about a factor of 3 over the entire map.
" Using HIFI we have observed a single spectrum at the peak position (R.A. = 01 34006779 Dee = 3074723711 (12000) of the eemission from 330? I reftig,spec).", Using HIFI we have observed a single spectrum at the peak position (R.A. = $^{\rm h}$ $^{\rm m}$ 79 Dec = 1 (J2000)) of the emission from 302 \\ref{fig_allspec}) ).
 TheHIFIspectrumweastakenon0lAugust2010duringonel , The HIFI spectrum was taken on 01 August 2010 during one hour of observing time using the load chop mode with the same reference position as was used for the PACS observation.
llineis1900536.9 MHz. knowntowithinanuncertaintyo f 1.3MH 0.2 s," The frequency of the line is 1900536.9 MHz, known to within an uncertainty of 1.3 MHz (0.2 ) \citep{Cooksy1986}."
s," The blue shifted line required to tune the local oscillator to 1899.268 GHz, about the highest frequency accessible to HIFI."
pectrawererecordedusingthewidebandacoustoopticalspectrometer.," The spectra were recorded using the wide band acousto optical spectrometer, covering a bandwidth of GHz for each polarization with a spectral resolution of MHz."
 co combiningbothpolarizations.," We calculated the noise-weighted averaged spectrum, combining bothpolarizations."
A fringe fitting o(6Eig.," A fringe fitting tool available within HIPE was used to subtract standing waves, subsequently the data were exported to CLASS for further analysis."
" "" MOM 3)."," Next, a linear baseline was subtracted and the spectrum was rebinned to a velocity resolution of 0.63 (Fig. \ref{fig_allspec}) )."
Wescaledtheresultingdatatothemainbeamscaleusingabeame f à surface accuracy of 7=3.8 ym (Olberg2010).," We scaled the resulting data to the main beam scale using a beam efficiency of, using the Ruze formula with the beam efficiency for a perfect primary mirror $\eta_{{\rm mb},0}=0.76$ and a surface accuracy of $\sigma=3.8\,\mu$ m \citep{OlbergBeam2010}."
. The measured peak temperature from a Gaussian fit is KK and the rms (lor limit) is mmK at kkms7! resolution. which is consistent within of the rms. predicted by HSPOT.," The measured peak temperature from a Gaussian fit is K and the rms $1\sigma$ limit) is mK at $^{-1}$ resolution, which is consistent within of the rms predicted by HSPOT."
" The half power beam width (HPBW) is 12.2"".", The half power beam width (HPBW) is $12.2''$ .
 For comparison with the sspectrum we have observed spectra of the (2-1) and (1-0), For comparison with the spectrum we have observed spectra of the (2–1) and (1–0)
limit on the CMB temperature at ;.= [.08 set by Luetal.(1996).. constrains the CMB temperature empirically.,"limit on the CMB temperature at $z = 4.08$ set by \cite*{Lu96}, constrains the CMB temperature empirically."
 If we assume that the CMB temperature varies monotonically with then the measure of Luetal.(1996) gives a<1.05 between ;:=QO and :=Ls. that is Τον;11.55 K at;=3.0513.," If we assume that the CMB temperature varies monotonically with $z$: then the measure of \cite*{Lu96} gives $\alpha < 1.05$ between $z =0$ and $z = 4.08$, that is $T_{\rm CMB} < 11.85$ K at $z = 3.0543$."
 This implies that byC(144) must be lower than 9.1«10| +. ie. only a minor contribution to the derived population ratio in our system. where ;V4δω for is between 0.02 and 0.03 (Table 2)).," This implies that $b_{01}U(\nu_{01})$ must be lower than $ 9.1 \times
10^{-4}$ $^{-1}$, i.e. only a minor contribution to the derived population ratio in our system, where $N_1/N_0$ for is between 0.02 and 0.03 (Table \ref{tab:CII}) )."
 As a first approximation therefore this term can be neglected and the same is true for byyU0)=shyiUor)., As a first approximation therefore this term can be neglected and the same is true for $b_{10}U(\nu_{01}) = \frac{1}{2} b_{01}U(\nu_{01})$.
 An infrared photon flux with an intensity comparable to the one measured in the Galactic plane (Bennettetal.1992.. Kogutetal. 1996)) would correspond to an excitation rate at least 2 orders of magnitude smaller than the one for à CMB photon flux with Τον=10 K and can be ignored.," An infrared photon flux with an intensity comparable to the one measured in the Galactic plane \cite{Bennett92}, \cite{Kogut96}) ) would correspond to an excitation rate at least 2 orders of magnitude smaller than the one for a CMB photon flux with $T_{\rm CMB}
= 10$ K and can be ignored."
 The other important type of photon excitation is UV photon pumping., The other important type of photon excitation is UV photon pumping.
 After the absorption of a photon an atom will usually cascade back through a variety of states. sometimes reaching levels that could not be populated by direct radiative upward transition from the ground state.," After the absorption of a photon an atom will usually cascade back through a variety of states, sometimes reaching levels that could not be populated by direct radiative upward transition from the ground state."
" If i represents all the quantum numbers for one of the upper levels. reached by photon absorption. the transition rate from level 0 to level I. is given by (Spitzer 1978)): where e,,4 Is the fraction of downward transitions from level i» that populate level 1. when the atom first reaches the group of lower levels."," If $m$ represents all the quantum numbers for one of the upper levels, reached by photon absorption, the transition rate from level 0 to level 1, is given by \cite{spitzer}) ): where $\epsilon_{m1}$ is the fraction of downward transitions from level $m$ that populate level 1, when the atom first reaches the group of lower levels."
" For transitions within a multiplet the values of e,,; are tabulated (e.g. Allen 1963)).", For transitions within a multiplet the values of $\epsilon_{mj}$ are tabulated (e.g. \cite{allen}) ).
" To evaluate Ly we considered all the direct upward transitions from ground state 252P. longwards of 900As: 27P!> 35 (AA903.00 D. 2PPο (AALOB7. 1036), 2D""> 2D (AA1335. 1331)."," To evaluate $\Gamma_{01}$ we considered all the direct upward transitions from ground state $2p^2P^0$, longwards of 900: $^2P^0 \rightarrow$ $^2P$ $\lambda\lambda903,904$ ), $^2P^0 \rightarrow$ $^2S$ $\lambda\lambda1037,1036$ ), $^2P^0 \rightarrow$ $^2D$ $\lambda\lambda1335,1334$ )."
" For the UV field we adopted the Milky Way spectral energy distribution (SED) given by Black(1987) with a Milky Way intensity at 912 oof LT«10Perecὃς'Uz! (Mathisetal. 1983). obtaining: that is xy,=2.3«104|."," For the UV field we adopted the Milky Way spectral energy distribution (SED) given by \cite*{Black87} with a Milky Way intensity at 912 of $4.7 \times 10^{-19}{\rm erg~cm^{-2}~s^{-1}~Hz^{-1}}$ \cite{Mathis83}) ), obtaining: that is $\gamma_{01} = 2.3\times 10^{-4}$."
 Since the UV pumping excitation rate 210 1s of the same order of magnitude as 70. the UV de-excitation term can also be omitted in ((4). for any likely UV flux intensity and SED.," Since the UV pumping de-excitation rate $\gamma_{10}$ is of the same order of magnitude as $\gamma_{01}$, the UV de-excitation term can also be omitted in (4), for any likely UV flux intensity and SED."
 The final shape of the balance equation ts: with the collisional term being dominated by the electrons or hydrogen atoms contributionaccording to the absorber ionization degree., The final shape of the balance equation is: with the collisional term being dominated by the electrons or hydrogen atoms contributionaccording to the absorber ionization degree.
 The UV field can also be expressed in terms of the hydrogen density through the tonization parameter. U—o()f/nyc. where o(IT) ts the surface flux of hydrogen-ionizing photons.," The UV field can also be expressed in terms of the hydrogen density through the ionization parameter, $U = \phi({\rm H})/n_{\rm H}c$, where $\phi({\rm H})$ is the surface flux of hydrogen-ionizing photons."
 If we assume a UV flux having the Milky Way SED given in Black(1987). 294=0.znU and: where f; is the fraction of particle j with respect to the hydrogen density.," If we assume a UV flux having the Milky Way SED given in \cite*{Black87}, $ \gamma_{01} = 0.7 n_{\rm H} U$ and: where $f_j$ is the fraction of particle $j$ with respect to the hydrogen density."
 From the ratio observed in our data at 2=3.0513 SOM derived log=3.2 and assuming solar abundance ratios a consistent fit to the data was obtained with log/=2.8 and Z~O.001Z.. (SOM)., From the ratio observed in our data at $z=3.0543$ SOM derived $\log U \geq -3.2$ and assuming solar abundance ratios a consistent fit to the data was obtained with $\log U = -2.8$ and $Z \sim 0.001 Z_{\odot}$ (SOM).
 As shown in Table 2.. the derived excitation temperature for the levels ./=3/2 and J=1/2 of for the damped Lya absorber toward 2619 at 2=3.0513 is between 19.6 and 21.6 K. We can use these values and expression | to constram the density and UV flux at the absorber.," As shown in Table \ref{tab:CII}, the derived excitation temperature for the levels $J= 3/2$ and $J=1/2$ of for the damped $\alpha$ absorber toward $-$ 2619 at $z=3.0543$ is between 19.6 and 21.6 K. We can use these values and expression \ref{eq:final} to constrain the density and UV flux at the absorber."
 However. in order to evaluate ((11)). we need to make an assumption about the ionization. degree of the gas in the absorber.," However, in order to evaluate \ref{eq:final}) ), we need to make an assumption about the ionization degree of the gas in the absorber."
 Since the efficiency of the electron and hydrogen collisional is very different. the collisional excitation term depends critically on the gas ionization degree.," Since the efficiency of the electron and hydrogen collisional is very different, the collisional excitation term depends critically on the gas ionization degree."
 To illustrate this we will consider two limiting cases for the absorbing cloud: a quasi-neutral gas. for which ».<0.002 my. and a tonized plasma.," To illustrate this we will consider two limiting cases for the absorbing cloud: a quasi-neutral gas, for which $n_e < 0.002~n_{\rm H}$ , and a ionized plasma."
 In the result of ((11)) is plotted. in the case of a quasi-neutral gas. where the collisional excitation ts due to collisions with atomic hydrogen.," In \\ref{fig:TexNeutral} the result of \ref{eq:final}) ) is plotted, in the case of a quasi-neutral gas, where the collisional excitation is due to collisions with atomic hydrogen."
 The horizontal continuous, The horizontal continuous
vyvalues and a decline in Που mictallicities above |Fe/H]~—1 (Origlia et al. 2003)).,values and a decline in for metallicities above $\feh \sim -1$ (Origlia et al. \cite{origlia03}) ).
 Johuson et al. (2009)).," Johnson et al. \cite{johnson09}) ),"
 on the other hand. fud that in uw CCen red giauts iucreases from about 0.2 ddex at [Fe/H]~1.7 to ddex at [Fe/U]~1.0.," on the other hand, find that in $\omega$ Cen red giants increases from about $-0.2$ dex at $\feh \sim -1.7$ to dex at $\feh \sim -1.0$."
 A similar lucrease Is not secu for the low-a halo stars., A similar increase is not seen for the $\alpha$ halo stars.
 Enhancements of Na aud a Na-O uiticorrelation are preseut in all wellstudicd elobular ¢‘listers (Carretta ct al. 2009)), Enhancements of Na and a Na-O anticorrelation are present in all well-studied globular clusters (Carretta et al. \cite{carretta09}) )
 aud may be caused by the chemical enrichment frou iutermecdiate-niass ACB stars undergoing lot-bottoni hydrogen burning., and may be caused by the chemical enrichment from intermediate-mass AGB stars undergoing hot-bottom hydrogen burning.
 According to the lvdrodvnamical simulations of D'Exrcole et al. (2008)).," According to the hydrodynamical simulations of D'Ercole et al. \cite{dercole08}) ),"
 the gas ejected from these ACB stags collects iu the cluster core via cooling flows. which may explain the differeuce in bbetween stars remaiiue i ως σαι itself aud those originating in the progeutor galaxy.," the gas ejected from these AGB stars collects in the cluster core via cooling flows, which may explain the difference in between stars remaining in $\omega$ Cen itself and those originating in the progenitor galaxy."
 We conclude that he derived abundance ratios provide clear evidence of two «istinct opulatious of stars that are among the most metal-rich in the Galactic halo., We conclude that the derived abundance ratios provide clear evidence of two distinct populations of stars that are among the most metal-rich in the Galactic halo.
 The reason that previous studies have failed o detect this dichotomy may be ascribed to he lower precision of the abundances for less houogeneous saniples of stars. aud ereater focus on metal-poor stars.," The reason that previous studies have failed to detect this dichotomy may be ascribed to the lower precision of the abundances for less homogeneous samples of stars, and greater focus on metal-poor stars."
 The La stars niv be ancient disk or buee stars ‘heated’ to halo kincimatics by merge satellite ealaxies or they could be the first stars formed iu a dissipative collapse of a xoto-Cialactic eas cloud., The $\alpha$ stars may be ancient disk or bulge stars `heated' to halo kinematics by merging satellite galaxies or they could be the first stars formed in a dissipative collapse of a proto-Galactic gas cloud.
 The ow-a stars are probable accreted from dwarf galaxies. and some are Likely to be associated with the w CCen progenitor galaxy.," The $\alpha$ stars are probably accreted from dwarf galaxies, and some are likely to be associated with the $\omega$ Cen progenitor galaxy."
 Further studies of possible correlations between the abundance ratios aud orbital paraueters of the stars may help us to clarify the origin of the two populations., Further studies of possible correlations between the abundance ratios and orbital parameters of the stars may help us to clarify the origin of the two populations.
0.3 to6.,0.3 to.
3!.. XRT data have been processed with the package v. 6.6.1 and corresponding calibration files: standard filtering and screening criteria have been applied., XRT data have been processed with the package v. 6.6.1 and corresponding calibration files: standard filtering and screening criteria have been applied.
" Piled-up Window Timing (WT) data have been corrected following the prescriptions by Romanoetal.(2006),, while piled-up Photon Counting (PC) data have been extracted from an annular region whose inner radius has been derived comparing the observed to the nominal point spread function (PSF, Morettietal.2005;; Vaughanetal. 2006))."," Piled-up Window Timing (WT) data have been corrected following the prescriptions by \cite{Romano06}, while piled-up Photon Counting (PC) data have been extracted from an annular region whose inner radius has been derived comparing the observed to the nominal point spread function (PSF, \citealt{Moretti05}; \citealt{Vaughan06}) )."
 The background is estimated from a source-free portion of the sky and then subtracted., The background is estimated from a source-free portion of the sky and then subtracted.
" The 0.3-10 keV background subtracted, PSF and vignetting corrected light-curve of each GRB has been re-binned so as to assure a minimum signal-to-noise (SN) equal to 4."," The 0.3-10 keV background subtracted, PSF and vignetting corrected light-curve of each GRB has been re-binned so as to assure a minimum signal-to-noise (SN) equal to 4."
 The count-rate light-curves are calibrated into luminosity light-curves using a time dependent count-to-flux conversion factoras described in Marguttietal.(2010a)., The count-rate light-curves are calibrated into luminosity light-curves using a time dependent count-to-flux conversion factoras described in \cite{Margutti10a}.
. This procedure produces luminosity curves where the possible spectral evolution of the source is properly taken into account., This procedure produces luminosity curves where the possible spectral evolution of the source is properly taken into account.
 Each 0.3-10 keV XRT light-curve is calibrated in the common rest frame energy band defined by the redshift distribution of the sample which turns out to be 2.2-14.4 keV. This allows us to make a direct comparison between the X-ray afterglows of different bursts while avoiding extrapolation of the signal to an unobserved energyband?., Each 0.3-10 keV XRT light-curve is calibrated in the common rest frame energy band defined by the redshift distribution of the sample which turns out to be 2.2-14.4 keV. This allows us to make a direct comparison between the X-ray afterglows of different bursts while avoiding extrapolation of the signal to an unobserved energy.
. We produced a software to automatically identify the smooth continuum underlying the X-ray afterglow of GRBs with superimposed flaring activity., We produced a software to automatically identify the smooth continuum underlying the X-ray afterglow of GRBs with superimposed flaring activity.
" The procedure is based on the x? statistics and can be described as a two-step process: A first blind fit of the entire X-ray afterglow (continuum plus flares) is done in log-log units using a power-law, smoothly joint broken power-law or double broken power-law models."," The procedure is based on the $\chi^2$ statistics and can be described as a two-step process: A first blind fit of the entire X-ray afterglow (continuum plus flares) is done in log-log units using a power-law, smoothly joint broken power-law or double broken power-law models."
" If the P-value (probability of obtaining a result at least as extreme as the one that is actually observed) associated to this fit is lower than5%,, the data point with the largest residual is removed from the light-curve and a new fit is performed."," If the P-value (probability of obtaining a result at least as extreme as the one that is actually observed) associated to this fit is lower than, the data point with the largest residual is removed from the light-curve and a new fit is performed."
 This process is repeated until a P- >5% is obtained., This process is repeated until a P-value $>5\%$ is obtained.
 The F-test is used to choose between the different nested models when necessary., The F-test is used to choose between the different nested models when necessary.
 The best fitting model satisfying the P-value condition is identified with the underlying continuum associated to a particular GRB ray afterglow (red dot-dashed line in Fig., The best fitting model satisfying the P-value condition is identified with the underlying continuum associated to a particular GRB X-ray afterglow (red dot-dashed line in Fig.
 1 for the case of 0060607A) and is subtracted from the original light-, \ref{Fig:example} for the case of 060607A) and is subtracted from the original light-curve.
 The resulting residuals and respective errors (both the statistical uncertainty associated to the original light-curve bins and the one coming from the continuum are properly taken into account and propagated) constitute the candidate X-ray flaring component associated to a GRB: this is shown in the inset of Fig., The resulting residuals and respective errors (both the statistical uncertainty associated to the original light-curve bins and the one coming from the continuum are properly taken into account and propagated) constitute the candidate X-ray flaring component associated to a GRB: this is shown in the inset of Fig.
 1 for 00606074. Note that no particular flare functional shape is assumed in this analysis., \ref{Fig:example} for 060607A. Note that no particular flare functional shape is assumed in this analysis.
 This method allows us to account for small variations superimposed to the continuum: in C10 flares were instead visually identified and then fitted with a specific profile., This method allows us to account for small variations superimposed to the continuum: in C10 flares were instead visually identified and then fitted with a specific profile.
 The average GRB 2.2-14.4 keV rest frame flaring component for any time interval t;— is computed as: where N is the number of GRB displaying a positive flaring component at a minimum 2c significance during t;—ts; Li is the 2.2-14.4 keV luminosity of the flaring component of the i* burst evaluated at at t=(t;+t5)/2., The average GRB 2.2-14.4 keV rest frame flaring component for any time interval $t_i-t_f$ is computed as: where N is the number of GRB displaying a positive flaring component at a minimum $\sigma$ significance during $t_i-t_f$; $L_{i}$ is the 2.2-14.4 keV luminosity of the flaring component of the $i^{th}$ burst evaluated at at $t=(t_i+t_f)/2$.
 Linear interpolation is used when necessary., Linear interpolation is used when necessary.
" The uncertainty affecting (L) is found through standard propagation of the uncertainties affecting each flaring component, as determined in the previous paragraph."," The uncertainty affecting $\langle L \rangle$ is found through standard propagation of the uncertainties affecting each flaring component, as determined in the previous paragraph."
 The typical uncertainty on each GRB flaring component data point is of the order of 30% the value of the GRB afterglow light-curve the flare subtraction., The typical uncertainty on each GRB flaring component data point is of the order of $30\%$ the value of the GRB afterglow light-curve the flare subtraction.
" An example is shown in Fig. 1,,"," An example is shown in Fig. \ref{Fig:example},"
 inset., inset.
 The use of the mean in Eq., The use of the mean in Eq.
" 1 prevents (L) from being dominated by a single bright event (the use of the median has been proven to lead to very similar results); instead, the linear mean of the excesses leads to (L) values biassed towards the bright end of the luminosity distribution of the excesses at any time t, and is therefore discarded."," \ref{Eq:avelum} prevents $\langle L \rangle$ from being dominated by a single bright event (the use of the median has been proven to lead to very similar results); instead, the linear mean of the excesses leads to $\langle L \rangle$ values biassed towards the bright end of the luminosity distribution of the excesses at any time $t$ , and is therefore discarded."
 0050904 shows an extended flaring activity, 050904 shows an extended flaring activity
constrain its rotation and to get clues on its internal structure and orientation.,constrain its rotation and to get clues on its internal structure and orientation.
the chromosphere of four giant stars.,the chromosphere of four giant stars.
 In. Sect., In Sect.
 5. à semi-empirical model of the chromosphere of & Cet Is used to interpret spectroscopic and interferometric data.," 5, a semi-empirical model of the chromosphere of $\beta$ Cet is used to interpret spectroscopic and interferometric data."
 Conclusions are presented in the last section., Conclusions are presented in the last section.
 Seven K giant stars. listed in Table |.. were observed at medium and high spectral resolution with the Visible spEctroGraph And polarimeter (VEGA. ?)) integrated within the CHARA array at Mount Wilson Observatory (California. USA. ?)).," Seven K giant stars, listed in Table \ref{table:1}, were observed at medium and high spectral resolution with the Visible spEctroGraph And polarimeter (VEGA, \citealt{mourard09}) ) integrated within the CHARA array at Mount Wilson Observatory (California, USA, \citealt{brummelaar05}) )."
 We used two criteria to select the giant stars for the program: These two criteria ensured the optimal operation of VEGA when no fringe tracker is used., We used two criteria to select the giant stars for the program: These two criteria ensured the optimal operation of VEGA when no fringe tracker is used.
 Apart from B Ophicus. : Cepher. and 109 Herculis. which were Observed in the continuum only. all stars were observed in both the continuum and chromospheric spectral lines.," Apart from $\beta$ Ophicus, $\iota$ Cephei, and 109 Herculis, which were observed in the continuum only, all stars were observed in both the continuum and chromospheric spectral lines."
" We selected four lines with chromospheric cores: the H, Balmer line and the infrared tripet lines (849.8 nm. 854.2 nm and 866.2 nm)."," We selected four lines with chromospheric cores: the $_{\alpha}$ Balmer line and the infrared tripet lines $849.8$ nm, $854.2$ nm and $866.2$ nm)."
 Details of the observations can be found in Table 2.., Details of the observations can be found in Table \ref{table:2}.
 The shortest baseline of the CHARA array (SIS2. 33 m) was used during the observations and the data were recorded in two spectral bands of the continuum simultaneously using the two detectors of VEGA (blue and red detectors).," The shortest baseline of the CHARA array (S1S2, $33$ m) was used during the observations and the data were recorded in two spectral bands of the continuum simultaneously using the two detectors of VEGA (blue and red detectors)."
 Calibrator stars were also observed ῃ the continuum in order to calibrate the measurements., Calibrator stars were also observed in the continuum in order to calibrate the measurements.
 We used the following sequence of observations when only one calibrator was available and the sequence when two calibrators were available., We used the following sequence of observations when only one calibrator was available and the sequence when two calibrators were available.
 We used the medium spectral resolution mode of VEGA (R= 5.000) for the observations in the continuum.," We used the medium spectral resolution mode of VEGA $R=5,000$ ) for the observations in the continuum."
 The high spectral resolution mode of VEGA (R= 30.000) was used for the observations of spectral lines for which calibrators were not required because the visibility in. the spectral lines were calibrated by the measurements in. the continuum. close to the spectral lines (see Sect. 2.2)).," The high spectral resolution mode of VEGA $R=30,000$ ) was used for the observations of spectral lines for which calibrators were not required because the visibility in the spectral lines were calibrated by the measurements in the continuum, close to the spectral lines (see Sect. \ref{seqdataproc}) )."
" We selecte the calibrators using the SearchCal developed at JMMC (?).. providing an estimate of the limb-darkened (LD) angular diameter (06,55)."," We selecte the calibrators using the SearchCal developed at JMMC \citep{bonneau06}, providing an estimate of the limb-darkened (LD) angular diameter $\theta_{LD}$ )."
 The uniform-disk (UD) angular diameter (05/5) is required to estimate the transfer function of the instrument at each wavelength., The uniform-disk (UD) angular diameter $\theta_{UD}$ ) is required to estimate the transfer function of the instrument at each wavelength.
 The UD angular diameter of each calibrator is derived using 655 with the linear LD coefficients given by ? and ?.., The UD angular diameter of each calibrator is derived using $\theta_{LD}$ with the linear LD coefficients given by \citet{claret95} and \citet{diaz95}.
 For each calibrator. Table presents θε at 620 nm and 790 nm.," For each calibrator, Table \ref{table:3} presents $\theta_{UD}$ at 620 nm and 790 nm."
 The data processing of VEGA is composed of two parts., The data processing of VEGA is composed of two parts.
 First. the data in the stellar continuum (at medium spectral resolution) are processed with the power spectral method giving the squared visibilities.," First, the data in the stellar continuum (at medium spectral resolution) are processed with the power spectral method giving the squared visibilities."
 Second. the processing of the data in spectral lines is based on the cross-spectrum method. which provides differential visibilities and phases across the lines (see ? for details).," Second, the processing of the data in spectral lines is based on the cross-spectrum method, which provides differential visibilities and phases across the lines (see \citealt{mourard09} for details)."
 For this analysis anc in the case of observations at medium spectral resolution. we divide the whole spectral band recorded by the red detector into four spectral channels of 10 nm centered onthe wavelengths 775 nm. 785 nm. 795 nm. and 805 nm.," For this analysis and in the case of observations at medium spectral resolution, we divide the whole spectral band recorded by the red detector into four spectral channels of $10$ nm centered onthe wavelengths $775$ nm, $785$ nm, $795$ nm, and $805$ nm."
 For the blue detector. the continuum ts visible only at 625.5 nm. we then use only one spectral channel (with a spectral bandwidth of 15 nm).," For the blue detector, the continuum is visible only at $625.5$ nm, we then use only one spectral channel (with a spectral bandwidth of $15$ nm)."
 The processing of these five spectral bands give the squared visibilities used to constrain the LD angular diameter (see Sect. 3.1))., The processing of these five spectral bands give the squared visibilities used to constrain the LD angular diameter (see Sect. \ref{angdiamsusec}) ).
 The cross-spectrum method is applied to the data recorded in the high spectral resolution mode., The cross-spectrum method is applied to the data recorded in the high spectral resolution mode.
" We compute the complex differential visibility between a large spectral channel used às reference (centered αἱ ο) and a narrow spectral band (centered at +) sliding in the reference spectral channel where V, and 6, represent the visibility and the phase of the fringe patterns at the wavelength tj.", We compute the complex differential visibility between a large spectral channel used as reference (centered at $\lambda_1$ ) and a narrow spectral band (centered at $\lambda_2$ ) sliding in the reference spectral channel where $V_{\lambda_i}$ and $\phi_{\lambda_i}$ represent the visibility and the phase of the fringe patterns at the wavelength $\lambda_i$ .
 The width of the sliding narrow spectral band (Als hereafter) is chosen to ensure a sufficient signal-to-noise ratio (S/N) for the estimation of, The width of the sliding narrow spectral band $\Delta\lambda_2$ hereafter) is chosen to ensure a sufficient signal-to-noise ratio $(S/N)$ for the estimation of
Often. hardness of spectra is described by the ratio of count rates or fluxes. A. in two channels.,"Often, hardness of spectra is described by the ratio of count rates or fluxes, $R$, in two channels."
 However. even if energy or photon fluxes are used. this value depends on the width of the channels. and is just specific to a given instrument.," However, even if energy or photon fluxes are used, this value depends on the width of the channels, and is just specific to a given instrument."
 On the other hand. a hardness ratio can be uniquely converted to the equivalent spectral index. I. corresponding to a power law spectrum with the same ratio of Huxes as the actual data.," On the other hand, a hardness ratio can be uniquely converted to the equivalent spectral index, $\Gamma$, corresponding to a power law spectrum with the same ratio of fluxes as the actual data."
 Thus. this gives a local spectral slope. which= ean usually be interpreted physically. unlike the hardness ratio. which is an abstract and instrument-dependent quantity.," Thus, this gives a local spectral slope, which can usually be interpreted physically, unlike the hardness ratio, which is an abstract and instrument-dependent quantity."
 Note hat this index is a unique and monotonous function of the flux jardness ratio., Note that this index is a unique and monotonous function of the flux hardness ratio.
 This index is analogous to colours used in stellar astrophysics. which also describe a spectrum in an instrument-independent way.," This index is analogous to colours used in stellar astrophysics, which also describe a spectrum in an instrument-independent way."
 This technique was introduced by 2032. who. 1owever. gave no details of the calculation.," This technique was introduced by Z02, who, however, gave no details of the calculation."
 We present them here., We present them here.
 The condition of the ratio of the energy fluxes in a power law spectrum being equal to that in the observed spectrum. A. can be written as. where the channel boundary energies are CEj..Bo). (Εν.Ej.," The condition of the ratio of the energy fluxes in a power law spectrum being equal to that in the observed spectrum, $R$, can be written as, where the channel boundary energies are $(E_1,\,E_2)$, $(E_3,\,E_4)$ ."
 For adjacent channels with increasing energy. Εν= (assumed in this work). which. however. is not required in this formalism.," For adjacent channels with increasing energy, $E_2=E_3$ (assumed in this work), which, however, is not required in this formalism."
 We define. In order to tind ΓΙΑ). we need to solve. Equation CÀ3)) can be readily solved by Newton's method. in which case the derivative is also needed. The uncertaintyof E is given by AU=AR/|uR/dI]. where AR is the uncertainty of &.," We define, In order to find $\Gamma(R)$, we need to solve, Equation \ref{f_g}) ) can be readily solved by Newton's method, in which case the derivative is also needed, The uncertaintyof $\Gamma$ is given by $\Delta \Gamma = \Delta R / \vert {\rm d} R/{\rm d} \Gamma\vert$, where $\Delta R$ is the uncertainty of $R$."
 We find. If R is given by the ratio of photon. rather than energy. fluxes. then 2—T is replaced by |—T in equationCAL) and e=1—TI.," We find, If $R$ is given by the ratio of photon, rather than energy, fluxes, then $2-\Gamma$ is replaced by $1-\Gamma$ in equation\ref{eq:hr}) ) and $g=1-\Gamma$ ."
 The special cases of I2 are now forT =|., The special cases of $\Gamma=2$ are now for $\Gamma=1$ .
HeCdTe arrays). being sampled continuously during the exposure. the information lost is only tliat fraction accumulatedeffer the CR hit.,"HgCdTe arrays), being sampled continuously during the exposure, the information lost is only that fraction accumulated the CR hit."
 There should be little difficulty ideutifviug cosmic rays in space-borne unages. as the vast majority of hits cover mauy pixels aud deposit thousauds of electrous.," There should be little difficulty identifying cosmic rays in space-borne images, as the vast majority of hits cover many pixels and deposit thousands of electrons."
 On the grouud. cosmic rays cause negligible loss of information.," On the ground, cosmic rays cause negligible loss of information."
 Iu theSNAP mission it will be important to derive accurate colors for resolved galaxies so as o obtain photometric redshilt estimates for lost galaxies., In the mission it will be important to derive accurate colors for resolved galaxies so as to obtain photometric redshift estimates for host galaxies.
 This places performance requirements on the S/N of galaxy photometry., This places performance requirements on the $S/N$ of galaxy photometry.
 For a galaxy with known intrinsic Πιν distribution g(x). the yest possible S/N on the total Hux is derivable through the same Fisher information formalism as Or point sources.," For a galaxy with known intrinsic flux distribution $g({\bf x})$, the best possible $S/N$ on the total flux is derivable through the same Fisher information formalism as for point sources."
 This is equivalent to measuring the flux ina Wiener-filtered image., This is equivalent to measuring the flux in a Wiener-filtered image.
 Iu. practice. i0wever. galaxies come iu au iufinite variety of shapes. so oue caunotpriori choose the ideal ilter for each image.," In practice, however, galaxies come in an infinite variety of shapes, so one cannot choose the ideal filter for each image."
 More practical is to measure the flux through some precletermined aperture ol shape w(x): This weighted flux is not. useful for studies requiring absolute total luminosities for galaxies. but will provide very accurate galaxy colors if matched apertures are used iu different wavelenetl bands.," More practical is to measure the flux through some predetermined aperture of shape $w({\bf x})$: This weighted flux is not useful for studies requiring absolute total luminosities for galaxies, but will provide very accurate galaxy colors if matched apertures are used in different wavelength bands."
 E will assuime. for simplicity. that both the galaxy aud the weight are circularly syοίτς.," I will assume, for simplicity, that both the galaxy and the weight are circularly symmetric."
 Then a simple propagation of errors gives a S/N ratio for the fie. when Nyquist-sampled. of The first teri in the utunerator is the source shot noise. the second terim is from the white-noise backgrouud arising from sky (or dark. read) counts of (4 per uuit area.," Then a simple propagation of errors gives a $S/N$ ratio for the $f_w$, when Nyquist-sampled, of The first term in the numerator is the source shot noise, the second term is from the white-noise background arising from sky (or dark, read) counts of $n$ per unit area."
 A useful. choice. of. weight. is. the Gaussian.. w=ePopp7/77.," A useful choice of weight is the Gaussian, $w=e^{-r^2/2\sigma^2}$."
 The Fourier4. transform. is. of. course also a Gaussian. aud can be evaluated for any candidate ePSF P! aud galaxy. profile g(r).," The Fourier transform is of course also a Gaussian, and can be evaluated for any candidate ePSF $P^\prime$ and galaxy profile $g(r)$."
 The size of the weight function σ can be adjusted to optimize the S/N for each galaxy on an image., The size of the weight function $\sigma$ can be adjusted to optimize the $S/N$ for each galaxy on an image.
 Clearly the effect of finite resolution iu the PSF or PRE is to remove high-& information which uuieht be present in the galaxy image., Clearly the effect of finite resolution in the PSF or PRF is to remove $k$ information which might be present in the galaxy image.
 If the galaxy. scale is larger than the PSF. then the PSF is irrelevant.," If the galaxy scale is larger than the PSF, then the PSF is irrelevant."
 Stall galaxies reduce to the point-source Limit., Small galaxies reduce to the point-source limit.
" Below I evaluate the resultant S/N for exponential«lisk galaxies (gyx€ "").", Below I evaluate the resultant $S/N$ for exponential-disk galaxies $g\propto e^{-\alpha r}$ ).
 Iu this case a Craussian-welghted. flux measurement is ouly a few percent noisier than the optimally-weiglitecd, In this case a Gaussian-weighted flux measurement is only a few percent noisier than the optimally-weighted
"the depth of the potential well, more than the total mass, plays a key rólle in determining the strength of the UV excess.","the depth of the potential well, more than the total mass, plays a key rôlle in determining the strength of the UV excess."
" A multiple regression upon —My and logoo confirms that all of the trend is accounted for by logco, with no significant residual luminosity dependence."," A multiple regression upon $-M_V$ and $\log{\sigma_0}$ confirms that all of the trend is accounted for by $\log{\sigma_0}$, with no significant residual luminosity dependence."
 The correlations with central (FUV-NUV) are slightly stronger (SRCC = -0.86; p=10? for the logao case) which is unsurprising as loggo is a central measurement., The correlations with central (FUV-NUV) are slightly stronger (SRCC = -0.86; $p = 10^{-9}$ for the $\log{\sigma_0}$ case) which is unsurprising as $\log{\sigma_0}$ is a central measurement.
 Figure 6 shows the logarithmic colour gradients plotted against My and logco., Figure \ref{fig:GradientPlots} shows the logarithmic colour gradients plotted against $M_V$ and $\log{\sigma_0}$.
" All elliptical galaxies have strong positive Vpw , but the gradient does not correlate strongly with either My or Ίοσσο."," All elliptical galaxies have strong positive $\nabla_{FN}$ , but the gradient does not correlate strongly with either $M_V$ or $\log{\sigma_0}$ ."
" Vw; is predominantly negative, and correlates strongly with My and in particular with logco (SRCC 20.80; p« 0.0001), in the sense that the largest negative gradients are in galaxies with logao«2.35."," $\nabla_{NJ}$ is predominantly negative, and correlates strongly with $M_V$ and in particular with $\log{\sigma_0}$ (SRCC = 0.80; $p < 0.0001$ ), in the sense that the largest negative gradients are in galaxies with $\log{\sigma_0} < 2.35$."
" Colour gradients in early-type galaxies, particularly those involving the bluer observable bands, have long been thought to be a measure of a metallicity gradient, and thus the dependence of gradients upon kinematic parameters has been used as a probe of galaxy formation models."," Colour gradients in early-type galaxies, particularly those involving the bluer observable bands, have long been thought to be a measure of a metallicity gradient, and thus the dependence of gradients upon kinematic parameters has been used as a probe of galaxy formation models."
 Peletier et al. (, Peletier et al. (
"1990) find no significant correlation between Vvg and My, although they cover a wider range in absolute magnitude than we do and there is some indication in their Figure 12 that the gradient is strongest for galaxies with —20>Mp-—24 and weaker at both brighter magnitudes, and at fainter ones which we do not cover.","1990) find no significant correlation between $\nabla_{UR}$ and $M_V$, although they cover a wider range in absolute magnitude than we do and there is some indication in their Figure 12 that the gradient is strongest for galaxies with $-20 > M_B > -24$ and weaker at both brighter magnitudes, and at fainter ones which we do not cover."
 La Barbera et al. (, La Barbera et al. (
"2005) find no significant dependence of Vr upon Mr, however they do find an environmental dependence, with steeper gradients away from rich clusters.","2005) find no significant dependence of $\nabla_{gr}$ upon $M_R$, however they do find an environmental dependence, with steeper gradients away from rich clusters."
 La Barbera et al. (, La Barbera et al. (
2010) define a composite colour gradient V. and find very weak dependence upon loggo in a sample of early-type galaxies from the Sloan Digital Sky Survey (SDSS).,2010) define a composite colour gradient $\nabla_*$ and find very weak dependence upon $\log{\sigma_0}$ in a sample of early-type galaxies from the Sloan Digital Sky Survey (SDSS).
 Tortora et al. (, Tortora et al. (
2010) analyse the colour gradients in a much larger sample of galaxies from SDSS.,2010) analyse the colour gradients in a much larger sample of galaxies from SDSS.
" For the early-type galaxies in their sample they find a complex dependence of V,; upon logoo, with the strongest gradients occurring at logo~2.0, with weaker gradients in higher velocity dispersion galaxies, and weaker or even positive gradients at logoo«2.0."," For the early-type galaxies in their sample they find a complex dependence of $\nabla_{gi}$ upon $\log{\sigma_0}$, with the strongest gradients occurring at $log{\sigma} \sim 2.0$, with weaker gradients in higher velocity dispersion galaxies, and weaker or even positive gradients at $\log{\sigma_0} \leq 2.0$."
 Tortora et al. (, Tortora et al. (
2010) and La Barbera et al. (,2010) and La Barbera et al. (
"2010) concur that metallicity gradients are the main driver of the colour gradients, with age a secondary factor.","2010) concur that metallicity gradients are the main driver of the colour gradients, with age a secondary factor."
 This is borne out by a number of studies of spectroscopically determined gradients in stellar population parameters (Kobayashi Arimoto 1999; Ogando et al., This is borne out by a number of studies of spectroscopically determined gradients in stellar population parameters (Kobayashi Arimoto 1999; Ogando et al.
 2005; Spolaor et al., 2005; Spolaor et al.
 2009)., 2009).
 A further complication with the interpretation of our trend of VΝο with logoo as evidence of a dependence on the metallicity gradient is the possibility that the UV excess stars will contribute some flux in the NUV band as well as in FUV., A further complication with the interpretation of our trend of $\nabla_{NJ}$ with $\log{\sigma_0}$ as evidence of a dependence on the metallicity gradient is the possibility that the UV excess stars will contribute some flux in the NUV band as well as in FUV.
" Although Vrw does not correlate strongly with logc, the gobal value of (FUV-NUV) does, and this, combined with the overall positive gradients in this colour could lead to some leakage of the FUV excess flux into the NUV band, preferentially at small radii and for galaxies with large logco."," Although $\nabla_{FN}$ does not correlate strongly with $\log{\sigma}$, the gobal value of (FUV-NUV) does, and this, combined with the overall positive gradients in this colour could lead to some leakage of the FUV excess flux into the NUV band, preferentially at small radii and for galaxies with large $\log{\sigma_0}$."
 Rawle et al. (, Rawle et al. (
"2008) discuss the effect of the FUV excess upon the scatter in the (NUV-J) colour-magnitude relation, and show that it does not contribute significantly to this scatter.","2008) discuss the effect of the FUV excess upon the scatter in the (NUV-J) colour-magnitude relation, and show that it does not contribute significantly to this scatter."
" For all except four of our galaxies there are estimates of stellar population parameters log(Age), [Z/H] and [o/Fe] in the literature."," For all except four of our galaxies there are estimates of stellar population parameters $\log{(Age)}$, [Z/H] and $\alpha$ /Fe] in the literature."
" There are differences in methodology between these sources, in terms of the spatial extent of the region observed, and the model spectra used to convert line index measurements into population parameters."," There are differences in methodology between these sources, in terms of the spatial extent of the region observed, and the model spectra used to convert line index measurements into population parameters."
 Ogando et al. (, Ogando et al. (
2010) use the SSP models of Thomas et al. (,2010) use the SSP models of Thomas et al. (
2003) to convert the line strength measurements presented in Ogando et al. (,2003) to convert the line strength measurements presented in Ogando et al. (
2008) to population parameters.,2008) to population parameters.
 These line strengths were measured from spectra taken with a 4.1 x 2.5 arcsec slit., These line strengths were measured from spectra taken with a 4.1 x 2.5 arcsec slit.
 Annibali et al. (, Annibali et al. (
"2007) use a new set of a-enhanced SSP models, to fit to line strength measurements presented in Rampazzo et al. (","2007) use a new set of $\alpha$ -enhanced SSP models, to fit to line strength measurements presented in Rampazzo et al. ("
2005) and Annibali et al. (,2005) and Annibali et al. (
"2006), these line index measurements are within an aperture of R./8.","2006), these line index measurements are within an aperture of $R_e/8$."
 Annibali et al. (, Annibali et al. (
"2007) present values of Z, which we convert to [Z/H] using Zo = 0.018.","2007) present values of Z, which we convert to [Z/H] using $_{\odot}$ = 0.018."
 Kuntschner et al. (, Kuntschner et al. (
"2010) use the SAURON integral-field spectrograph to derive line index maps, and fit the stellar population models of Schiavon (2007) to these.","2010) use the SAURON integral-field spectrograph to derive line index maps, and fit the stellar population models of Schiavon (2007) to these."
" They present derived stellar population parameters within R./8 and Γο, we use the former."," They present derived stellar population parameters within $R_e/8$ and $R_e$, we use the former."
 Spolaor at el. (, Spolaor at el. (
"2010) report stellar population parameters as derived in a number of papers by the Swinburne group, for the galaxies in our sample the original sources are Sánnchez-Blázzquez et al. (","2010) report stellar population parameters as derived in a number of papers by the Swinburne group, for the galaxies in our sample the original sources are Sánnchez-Blázzquez et al. ("
2007; Reda at al. (,2007; Reda at al. (
2007); Spolaor et al. (,2007); Spolaor et al. (
2008); Brough et al. (,2008); Brough et al. (
2007) and Proctor Sansom (2002).,2007) and Proctor Sansom (2002).
" Their values are measured within R./8, and measured line indices are converted to stellar population parameters using the models of Thomas et al. ("," Their values are measured within $R_e/8$, and measured line indices are converted to stellar population parameters using the models of Thomas et al. ("
2003).,2003).
 Serra et al. (, Serra et al. (
"2008) present stellar population parameters within R./16 derived using the stellar population models of Bruzual Charlot (2003), and the models of the index response to [E/Fe] presented by Lee et al. (","2008) present stellar population parameters within $R_e/16$ derived using the stellar population models of Bruzual Charlot (2003), and the models of the index response to [E/Fe] presented by Lee et al. ("
2009).,2009).
 Thomas et al. (, Thomas et al. (
"2005) measure line indices within Rz,/10, and derive stellar population using their own SSP models.","2005) measure line indices within $R_e/10$, and derive stellar population using their own SSP models."
 Loubser et al. (, Loubser et al. (
"2009) measure line indices within a slit of 1 arcsec by R./8, and convert these to population parameters usingthe SSP models of Thomas et al. (","2009) measure line indices within a slit of 1 arcsec by $R_e/8$, and convert these to population parameters usingthe SSP models of Thomas et al. ("
2003).,2003).
" Given the differences in observational technique, region and SSP models, we attempt to adjust the population parameters to a standard scale, using galaxies in common between the samples (all of the samples are of course larger than just those galaxies for which there are GALEX data)."," Given the differences in observational technique, region and SSP models, we attempt to adjust the population parameters to a standard scale, using galaxies in common between the samples (all of the samples are of course larger than just those galaxies for which there are GALEX data)."
" However even with the complete samples, the overlaps between studies are small andinconsistent."," However even with the complete samples, the overlaps between studies are small andinconsistent."
" Moreover the studywith the largest sample, and the most galaxies in common with other studies (Thomas et al."," Moreover the studywith the largest sample, and the most galaxies in common with other studies (Thomas et al."
 2005) is also the one for which the rms differences with other studies are the greatest., 2005) is also the one for which the rms differences with other studies are the greatest.
images in the Schwarzschild lensing.,images in the Schwarzschild lensing.
 However. such treatment is applicable only. when the deflection angle is an odd function of r.," However, such treatment is applicable only when the deflection angle is an odd function of $r$."
 If (he source and lens ave completely aligned. along the line of sight. the image is expected to be circular (an Einstein ring).," If the source and lens are completely aligned along the line of sight, the image is expected to be circular (an Einstein ring)."
" The Einstein radius Hj. which is defined as the radius of the circular image on the lens plane. is obtained from Equation (5)) with ο)=0 as The image positions can then be calculated from and where 0—b/D,=r/D, is the angle between the image and lens. and 05;=Rp/Dy, is the angular Einstein radius."," The Einstein radius $R_E$, which is defined as the radius of the circular image on the lens plane, is obtained from Equation \ref{eqn:proj}) ) with $\beta = 0$ as The image positions can then be calculated from and where $\theta = b / D_L \thickapprox r / D_L$ is the angle between the image and lens, and $\theta_E = R_E / D_L$ is the angular Einstein radius."
" Using reduced parameters 3=3/065 aud 0=0/05. Equations (8)) aud (9)) become simple cubic Formulas: and As the discriminant of Equation (10)) is —4SOT<0, Equation (10)) has two conjugate complex solutions aud a real solution: with."," Using reduced parameters $\hat{\beta} = \beta / \theta_E$ and $\hat{\theta} = \theta / \theta_E$, Equations \ref{eqn:imgeq}) ) and \ref{eqn:imgeq2}) ) become simple cubic formulas: and As the discriminant of Equation \ref{eqn:poly}) ) is $-4\hat{\beta}^3 - 27 < 0$, Equation \ref{eqn:poly}) ) has two conjugate complex solutions and a real solution: with,"
"racial extent of the quasars region observed. along the line of sight at time is (2???) where rg, is the neutral fraction and five=6Rae)fe is the age of the quasar corresponding to the time when photons that reach f at time / were emitted.","radial extent of the quasars region observed along the line of sight at time is \citep{white2003,haiman2002,fan2006}
 where $x_{\HIeqn}$ is the neutral fraction and $t_{\rm age} = t-R(t_{\rm age})/c$ is the age of the quasar corresponding to the time when photons that reach $R$ at time $t$ were emitted."
 Phe value of F2 ds subject to large uncertainties. including the quasar lifetime ancl luminosity. as well as the clumpiness and ionisation state of the surrounding IGM.," The value of $R$ is subject to large uncertainties, including the quasar lifetime and luminosity, as well as the clumpiness and ionisation state of the surrounding IGM."
 Beeause of the method used. to include the ionising eLllect of quasars at locations bevond the sphere that would be generated by the quasar alone in a fully neutral ancl uniform ICM. we put Hy=d when calculating 2.," Because of the method used to include the ionising effect of quasars at locations beyond the sphere that would be generated by the quasar alone in a fully neutral and uniform IGM, we put $x_{\HIeqn} = 1$ when calculating $R$."
 We assume that the mass of the central black hole (BIL) scales as a power law with the circular. velocity ο of the host halo. Adinxec (?7)..," We assume that the mass of the central black hole (BH) scales as a power law with the circular velocity $v_c$ of the host halo, $M_{\rm bh} \propto v_c^{\gamma}$ \citep{wl2003}. ."
" The circular velocity of a halo of mass AZ, at redshift z varies as e(z)xAHFey|2)?2 (2), which leads to à power law relation between BIT mass and halo mass. AdinxAL)τα2377"," The circular velocity of a halo of mass $M_{\rm h}$ at redshift $z$ varies as $v_c(z) \propto M_{\rm h}^{1/3}(1+z)^{1/2}$ \citep{bl2001}, which leads to a power law relation between BH mass and halo mass, $M_{\rm bh} \propto M_{\rm h}^{5/3}(1+z)^{5/2}$."
", Uo the DE powering this quasar shines at a fraction η of its Eeldington luminosity. then we have LyxgMiunMg?(3|z)77, wheres=5 for simple mocels of feedback-limited. accretion (2).."," If the BH powering this quasar shines at a fraction $\eta$ of its Eddington luminosity then we have $L_{\rm q} \propto \eta M_{\rm bh} \propto \eta M_{\rm h}^{\gamma/3}(1+z)^{\gamma/2}$, where $\gamma = 5$ for simple models of feedback-limited accretion \citep{wl2003}."
" Since LyxINS. using Equation (9)) leads to (in proper units) 1n models where BL growth is limited by feedback over the gas dynamical time (e.g...2).. the quasar lifetime £4 1s proportional to //. and so f/f,x(1|z)17 [p lou].."," Since $L_{\rm q} \propto \dot{N_{\gamma}}$, using Equation \ref{RqWhite}) ) leads to (in proper units) In models where BH growth is limited by feedback over the gas dynamical time \citep[e.g.,][]{wl2003}, the quasar lifetime $t_{\rm q}$ is proportional to $H$ , and so $t_{\rm q} \propto (1+z)^{-3/2}$ for $z \gg 1$."
" Equating fase with /, and substituting into Equation (10)) elves Llere we adopt a quasar luminosity fraction of η=0.1 (27) and an index value 7=5 (?)..", Equating $t_{\rm age}$ with $t_{\rm q}$ and substituting into Equation \ref{2ndlastRq}) ) gives Here we adopt a quasar luminosity fraction of $\eta = 0.1$ \citep{martini2001} and an index value $\gamma = 5$ \citep{wp2006}. .
 Since the relative quasar contribution to reionisation is expected to be < 14 per cent (?7).. we have chosen to investigate fiducial relativeES quasar ionisation contributions of 10 per cent.," Since the relative quasar contribution to reionisation is expected to be $\lsim$ 14 per cent \citep{srbin2007}, we have chosen to investigate fiducial relative quasar ionisation contributions of 10 per cent."
 We therefore choose Ry ancl Adywin such that at overlap (24;=6 in our model) the ionising contribution from a population of quasars alone relative to the ionising contribution of galaxies takes a value Qo/Q.=0.1.," We therefore choose $R_0$ and $M_{\rm h,min}$ such that at overlap $z_{\rm ov} = 6$ in our model) the ionising contribution from a population of quasars alone relative to the ionising contribution of galaxies takes a value $Q_{\rm q}/Q_{\star} = 0.1$."
 We reiterate that the prescription. outlined. in this section is intended το provide a parametorisecd mocel. loosely based. on underlying physical models. rather than a complete physical model of all the processes involved.," We reiterate that the prescription outlined in this section is intended to provide a parameterised model, loosely based on underlying physical models, rather than a complete physical model of all the processes involved."
 The parameterisation Chosen provides a general set of models within which various different possible physical models are contained., The parameterisation chosen provides a general set of models within which various different possible physical models are contained.
 “Phe plausibility. of this modelling can be investigated by determining the values of various physical parameters for cach of our simulation. which we clenote by SI. S2 and 583 (in order of decreasing minimum halo mass).," The plausibility of this modelling can be investigated by determining the values of various physical parameters for each of our simulation, which we denote by S1, S2 and S3 (in order of decreasing minimum halo mass)."
 We do this for quasar dutw evele fou by noting that the fraction of active quasars with absolute bolometric magnitude My:26 Is where AL25 is the host halo mass corresponding to à quasar with absolute bolometric magnitude My=26., We do this for quasar duty cycle $f_{\rm on}$ by noting that the fraction of active quasars with absolute bolometric magnitude $\mathcal{M}_{\rm B} \leq -26$ is where $M_{-26}$ is the host halo mass corresponding to a quasar with absolute bolometric magnitude $\mathcal{M}_{\rm B} = -26$ .
" In order to solve for fi, we note that for 2ὃν1. Equating A? in Equation (9)) with à, in Equation (11)) enables Aly, to be written as a function. of f, and the model dependent: Ay. whieh can then be substituted in Equation (12))."," In order to solve for $f_{\rm on}$ we note that for $z \gg 1$, Equating $R$ in Equation \ref{RqWhite}) ) with $R_{\rm q}$ in Equation \ref{Rq}) ) enables $M_{\rm h}$ to be written as a function of $f_{\rm on}$ and the model dependent $R_0$, which can then be substituted in Equation \ref{eq:fon}) )."
" We use ceCMp10"" w;= 6(7?7) and N.CM=26)2. 10ss+ (2).", We use $\psi(\mathcal{M}_{\rm B} \leq -26) = 0.65 \times 10^{-9}$ $^{-3}$ at $z = 6$ \citep{fan2001} and $\dot{N}_{\gamma}(\mathcal{M}_{\rm B} = -26) = 2 \times 10^{57}$ $^{-1}$ \citep{white2003}.
 Figure 2 shows the solution for fi [rom Equation (12)) eraphically for fyz3. 4 and proper Alpe.," Figure \ref{fon} shows the solution for $f_{\rm on}$ from Equation \ref{eq:fon}) ) graphically for $R_0 \approx 3$, 4 and proper Mpc."
 Table 1 gives the values of these parameters used in simulations SL$3., Table \ref{tab:param} gives the values of these parameters used in simulations S1–S3.
" We find duty eveles of ~10.7 4.LO7. which at z&6 corresponds to &,~10"" 410* vvr."," We find duty cycles of $\sim 10^{-3}$ $4\times 10^{-2}$, which at $z \approx 6$ corresponds to $t_{\rm q} \sim 10^6$ $4\times 10^7$ yr."
 These are comparable to the Salpeter time for doubling of BLE mass and to estimates of quasar lifetime (?).., These are comparable to the Salpeter time for doubling of BH mass and to estimates of quasar lifetime \citep[][]{martini2004}.
 The prescription outlined above does not consider the effect of recombinations within cquasar-generated regions., The prescription outlined above does not consider the effect of recombinations within quasar-generated regions.
 ὃν investigating recombinations in inhomogeneous [fossil ionisecl regions around quasars. ο established. that most of the regions formed. curing hyelrogen reionisation remain hiehlyionised throughout the entire. reionisation process.," By investigating recombinations in inhomogeneous fossil ionised regions around quasars, \cite{furl2008} established that most of the regions formed during hydrogen reionisation remain highlyionised throughout the entire reionisation process."
 “They found that the ionising background. from ealaxies inside the fossilregions. together with any residual low-level emission. [rom the BIL past its bright quasar phase. eLlicienthy suppresses recombinations during hvdrogen reionisation.," They found that the ionising background from galaxies inside the fossilregions, together with any residual low-level emission from the BH past its bright quasar phase, efficiently suppresses recombinations during hydrogen reionisation."
 Thus. even though quasars are thought to be transient. the assumption of fully ionised relic bubbles vieldsa good approximationat the level of our model.," Thus, even though quasars are thought to be transient, the assumption of fully ionised relic bubbles yieldsa good approximationat the level of our model."
lt is ue to the fact that we are using a volume-limited sample. and have introduced à cutoll of —20.45 in r-band magnitude.,"It is due to the fact that we are using a volume-limited sample, and have introduced a cutoff of $-20.45$ in r-band magnitude."
 The radio-optical correlation at low (Liuica*107 )) luminosities means that in cliscarding objects with lower r-band magnitudes we are also excluding some radio sources., The radio-optical correlation at low $\LGHz < 10^{23}$ ) luminosities means that in discarding objects with lower r-band magnitudes we are also excluding some radio sources.
 1n Figure6 3. we illustrate this effect by binningm our sample1 in radio luminosity. and plotting distributions of r-band magnitudes in each bin.," In Figure \ref{fig:opticalDropouts} we illustrate this effect by binning our sample in radio luminosity, and plotting distributions of r-band magnitudes in each bin."
" Phere are relatively few objects with Mr21.8 at higher radio luminosities: however. when Lite,€4075 there is a significant contribution from galaxies with Alp> PLS. sugeesting that there is also an appreciable number of low luminosity radio sources with AM,>20.45 that we have discarded. clue to our optical completeness limit."," There are relatively few objects with $\Mr\/ > -21.8$ at higher radio luminosities; however, when $\LGHz < 10^{23}$ there is a significant contribution from galaxies with $\Mr\/>-21.8$ , suggesting that there is also an appreciable number of low luminosity radio sources with $\Mr\/>-20.45$ that we have discarded due to our optical completeness limit."
 Not surprisingly. the dilference between our Luminosity function and those for [Dux-limited. samples becomes apparent at around 10777," Not surprisingly, the difference between our luminosity function and those for flux-limited samples becomes apparent at around $10^{22.5}$."
 It is worth noting that the --break is not the only way to separate star-forming galaxies from radio AGNs in our sample., It is worth noting that the -break is not the only way to separate star-forming galaxies from radio AGNs in our sample.
 Another possible methocl is to use the concentration index. defined as the ratio of Petrosian light radii. Co=PsuIos.," Another possible method is to use the concentration index, defined as the ratio of Petrosian light radii, $C = R_{50}/R_{90}$."
 Low values of (κου 0.33) correspond to earlv-tvpe galaxies. which are more likely to harbour powerful racio sources. while late-tvpe (star-forming) galaxies are those with €'-0.315.," Low values of $C$ $< \sim 0.33$ ) correspond to early-type galaxies, which are more likely to harbour powerful radio sources, while late-type (star-forming) galaxies are those with $C>0.375$."
 We have carried out the subsequent analysis using both the Oy(4000) and concentration index diagnostics to split. up our sample. ancl found: very. similar results.," We have carried out the subsequent analysis using both the $\Dn\/(4000)$ and concentration index diagnostics to split up our sample, and found very similar results."
 Therefore. only the --break demarcation findings are presented. here.," Therefore, only the -break demarcation findings are presented here."
 We construct a bivariate luminosity function by binning all objects in stellar mass ancl plotting the radio luminosity. unction. for each bin., We construct a bivariate luminosity function by binning all objects in stellar mass and plotting the radio luminosity function for each bin.
 Fractions plotted: represent the number of AGNs in a certain stellar mass range brighter han a given racio luminosity. divided by the total number of objects in that stellar mass bin.," Fractions plotted represent the number of AGNs in a certain stellar mass range brighter than a given radio luminosity, divided by the total number of objects in that stellar mass bin."
 Again. corrections using he V.Vi method (Condon1989). were applied.," Again, corrections using the $V/\Vmax\/$ method \cite{Condon89} were applied."
 As r and Ζ spectroscopic bands largely track the same population of old bulge stars. we expect completeness in r-band to imply hat our sample is also complete in z-band. and hence in," As r and z spectroscopic bands largely track the same population of old bulge stars, we expect completeness in r-band to imply that our sample is also complete in z-band, and hence in"
is dominated by observational errors rather than systematic ones stemming from the choice of WD cooling curves (Salarisetal.2009).,is dominated by observational errors rather than systematic ones stemming from the choice of WD cooling curves \citep{Salaris09}.
". On the other hand, WD cooling ages vary more when interpolated using different models."," On the other hand, WD cooling ages vary more when interpolated using different models."
" To account for this systematic uncertainty, we assume a fractional error of 50% for the WD cooling ages obtained using the Fontaineetal.(2001) models."," To account for this systematic uncertainty, we assume a fractional error of $50\%$ for the WD cooling ages obtained using the \citet{Fontaine01} models."
" For nominal choices of WD composition, envelope thickness, energy loss rates, and opacities, the fractional uncertainty in WD cooling ages for the remainder of our sample is still «20% in most cases."," For nominal choices of WD composition, envelope thickness, energy loss rates, and opacities, the fractional uncertainty in WD cooling ages for the remainder of our sample is still $<20\%$ in most cases."
 Improvements in the observational measurements of aand aand the removal of cluster non-members from the sample more than offset any systematic affect stemming from the inclusion of a second group of WD cooling curve models in our analysis., Improvements in the observational measurements of and and the removal of cluster non-members from the sample more than offset any systematic affect stemming from the inclusion of a second group of WD cooling curve models in our analysis.
 Progenitor lifetimes are a function of WD cooling ages and cluster ages., Progenitor lifetimes are a function of WD cooling ages and cluster ages.
 The accuracy of our initial-final mass data set depends critically on constraining the cluster distance and age., The accuracy of our initial-final mass data set depends critically on constraining the cluster distance and age.
" Although a few clusters have measured parallaxes, the distances to most are inferred by main sequence fitting."," Although a few clusters have measured parallaxes, the distances to most are inferred by main sequence fitting."
" Photometric methods can be used to infer composition, reddening, and distance even for clusters with limited membership and spectroscopic data (see for example Pinsonneault et al."," Photometric methods can be used to infer composition, reddening, and distance even for clusters with limited membership and spectroscopic data (see for example Pinsonneault et al."
" 1997, 2004: An et al."," 1997, 2004; An et al."
 2007)., 2007).
" Well-studied open clusters can have spectroscopic metallicity measurements, extinction inferred from polarization studies, and both radial velocity and proper motion membership data."," Well-studied open clusters can have spectroscopic metallicity measurements, extinction inferred from polarization studies, and both radial velocity and proper motion membership data."
" The uncertainties in these basic cluster parameters are largely determined by the available information in cach specific cluster, which we discuss below, and set the uncertainty in the inferred cluster distance."," The uncertainties in these basic cluster parameters are largely determined by the available information in each specific cluster, which we discuss below, and set the uncertainty in the inferred cluster distance."
" The measured distance to the cluster yields the turnoff luminosity, which can be combined with isochrones to determine the age and the mass-main sequence lifetime relationship (see S09 for a discussion)."," The measured distance to the cluster yields the turnoff luminosity, which can be combined with isochrones to determine the age and the mass-main sequence lifetime relationship (see S09 for a discussion)."
 The errors in these two aspects are distinct in nature., The errors in these two aspects are distinct in nature.
 Core size deserves special comment., Core size deserves special comment.
" Convective core overshooting (Chiosi&Maeder1986) or rotational mixing (Maeder&Meynet2000) are difficult to model theoretically, and both mechanisms have the practical effect of extending the main sequence lifetime by providing extra fuel."," Convective core overshooting \citep{Chiosi86} or rotational mixing \citep{Maeder00} are difficult to model theoretically, and both mechanisms have the practical effect of extending the main sequence lifetime by providing extra fuel."
 The main sequence is broader than that predicted by models without overshoot or mixing. which is evidence that this phenomenon is real to some degree (e.g. 2010).," The main sequence is broader than that predicted by models without overshoot or mixing, which is evidence that this phenomenon is real to some degree \citep[\eg][]{Andersen91, Torres10}."
". Asa result, the main sequence lifetime-mass relationship has substantial theoretical uncertainties."," As a result, the main sequence lifetime-mass relationship has substantial theoretical uncertainties."
" However, the total fuel burned is more reliable than the cluster ages because the helium core mass at the end of the main sequence is less than the minimum core mass required for helium ienition."," However, the total fuel burned is more reliable than the cluster ages because the helium core mass at the end of the main sequence is less than the minimum core mass required for helium ignition."
" Uncertainties related to overshooting will thus be important for intermediate mass stars, where the star leaves the main sequence with a helium core greater than the minimum required"," Uncertainties related to overshooting will thus be important for intermediate mass stars, where the star leaves the main sequence with a helium core greater than the minimum required"
 , 
"zone, the instability does not work, because the front does not propagate, and we have only marginally stable solutions (Szuszkiewicz&Miller (1997))).","zone, the instability does not work, because the front does not propagate, and we have only marginally stable solutions \cite{szusz97}) )."
" For instance, Hameuryetal.(2009) find that the heating and cooling fronts of the ionization instability do not propagate strongly enough in their model to account for the large luminosity oscillations in AGN."," For instance, \cite{hameury09} find that the heating and cooling fronts of the ionization instability do not propagate strongly enough in their model to account for the large luminosity oscillations in AGN."
 The outburst cycles caused by the ionization instability are very sensitive to the viscosity parameter., The outburst cycles caused by the ionization instability are very sensitive to the viscosity parameter.
" As was shown already for dwarf novae, the amplitudes consistent with observations can be obtained only in the models with non-constant viscosity, ie. o in the hot state must be larger than in the cold state."," As was shown already for dwarf novae, the amplitudes consistent with observations can be obtained only in the models with non-constant viscosity, i.e. $\alpha$ in the hot state must be larger than in the cold state."
" This may also be the case in the X- binaries, however in AGN the situation may be different (see Janiuketal.(2004)))."," This may also be the case in the X-ray binaries, however in AGN the situation may be different (see \cite{janiuk04}) )."
" 'Thus the fast radiation pressure instability is expected to operate for an average accretion rate higher than a certain lower limit, mostly dependent on the adopted viscous scaling."," Thus the fast radiation pressure instability is expected to operate for an average accretion rate higher than a certain lower limit, mostly dependent on the adopted viscous scaling."
 There is an upper limit as well if the outflow is strongly increasing function of the Eddington ratio., There is an upper limit as well if the outflow is strongly increasing function of the Eddington ratio.
 The ionization instability should operate if the disc is large enough to show a partially ionized hydrogen zone for a given average accretion rate., The ionization instability should operate if the disc is large enough to show a partially ionized hydrogen zone for a given average accretion rate.
 Confronting the observational constraints with model predictions will in turn allow us to find the constraints for the viscosity parameterization and the role of the outflow., Confronting the observational constraints with model predictions will in turn allow us to find the constraints for the viscosity parameterization and the role of the outflow.
 In the next sections we make a preliminary step in this direction., In the next sections we make a preliminary step in this direction.
 The Galactic X-ray binary systems are variable in wide range of timescales., The Galactic X-ray binary systems are variable in wide range of timescales.
" First, the transient X-ray sources undergo their X-ray active states on timescales of years."," First, the transient X-ray sources undergo their X-ray active states on timescales of years."
" Second, some sources exhibit X-ray periodic variability on timescales of months."," Second, some sources exhibit X-ray periodic variability on timescales of months."
" Third, some of the most luminous sources are variable in timescales of tens-thousands of seconds."," Third, some of the most luminous sources are variable in timescales of tens-thousands of seconds."
 Finally many of X-ray binaries undergo quasi-periodic oscillations.," Finally, many of X-ray binaries undergo quasi-periodic oscillations."
 Direct comparison of these data and the models is not simple: even the identification of the type of the variability with the mechanism is not unique., Direct comparison of these data and the models is not simple: even the identification of the type of the variability with the mechanism is not unique.
" In Table 1 we summarize the properties of the exemplary, best studied X-ray binary sources found in the literature, which in our opinion may display radiation pressure or ionization instability."," In Table \ref{tab:binaries} we summarize the properties of the exemplary, best studied X-ray binary sources found in the literature, which in our opinion may display radiation pressure or ionization instability."
" We list their characteristic variability timescales and amplitudes, estimated Eddington ratios and disc sizes, as well as we indicate a possible instability mechanism responsible for the variability, whenever it is in agreement with our computations presented in Sec. 2.."," We list their characteristic variability timescales and amplitudes, estimated Eddington ratios and disc sizes, as well as we indicate a possible instability mechanism responsible for the variability, whenever it is in agreement with our computations presented in Sec. \ref{sec:results}."
" The maximum disc radius is estimated based on the of the Roche lobe size, from the simplified formula given by Paczyriski(1971), whenever we had the data for the system orbital parameters."," The maximum disc radius is estimated based on the of the Roche lobe size, from the simplified formula given by \cite{pac71}, whenever we had the data for the system orbital parameters."
 The estimates given in Table 1 should be treated only as indications., The estimates given in Table \ref{tab:binaries} should be treated only as indications.
 The information comes mostly from the literature and frequently relies on quantitative description., The information comes mostly from the literature and frequently relies on quantitative description.
 All objects in Table 1 display ionization instability since we selected objects classified as X-ray novae and there is very little uncertainty in the establishment of their nature.," All objects in Table \ref{tab:binaries}
 display ionization instability since we selected objects classified as X-ray novae and there is very little uncertainty in the establishment of their nature."
 Most of them show large amplitude outbursts lasting days., Most of them show large amplitude outbursts lasting days.
 The ratio of the maximum to minimum flux was estimated from the peak and the emission level at the end of the outburst so it represents the lower limit - only some of the X-ray novae have clear detection in the quiescence., The ratio of the maximum to minimum flux was estimated from the peak and the emission level at the end of the outburst so it represents the lower limit - only some of the X-ray novae have clear detection in the quiescence.
" Clearly, more careful observational analysis is needed to better study the amplitude pattern during the ionization instabilities in these sources."," Clearly, more careful observational analysis is needed to better study the amplitude pattern during the ionization instabilities in these sources."
" However, from the comparison of the accretion rates and disc sizes with our map presented in Figure 1 one can already infer some information about the individual sources."," However, from the comparison of the accretion rates and disc sizes with our map presented in Figure \ref{fig:topo}
 one can already infer some information about the individual sources."
" All but one among the sources are well within the instability strip, so the instability operates as expected."," All but one among the sources are well within the instability strip, so the instability operates as expected."
 Only, Only
Diffuse dust in our Galaxy is heated by the interstellar radiation field at temperatures ranging from 10K to 100K. dding on both the dimensions of the grains and the radiative environment.,"Diffuse dust in our Galaxy is heated by the interstellar radiation field at temperatures ranging from $\,$ K to $\,$ K, ding on both the dimensions of the grains and the radiative environment."
 Their emission ts thus in the sub-mm/FIR range. and is partially polarized. due to the alignment of grains in the Galactic magnetic field (see e.g. Draine 2003)).," Their emission is thus in the sub-mm/FIR range, and is partially polarized, due to the alignment of grains in the Galactic magnetic field (see e.g. Draine \cite{Drai03}) )."
 While FIR emission of interstellar dust has been mapped quite accurately. even at high galactic latitudes. by the IRAS. ISO and Spitzer surveys. and is investigated in deep detail by the currently operating Herschel observatory. its polarization properties are almost unknown (see e.g. Hildebrand Kirby 2004).," While FIR emission of interstellar dust has been mapped quite accurately, even at high galactic latitudes, by the IRAS, ISO and Spitzer surveys, and is investigated in deep detail by the currently operating Herschel observatory, its polarization properties are almost unknown (see e.g. Hildebrand Kirby \cite{Hild04}) )."
 The interest in. studying dust polarization is. twofold., The interest in studying dust polarization is twofold.
 An accurate measurement of dust polarization at different wavelengths is important to better understand the nature and structure of dust grains and to probe the magnetic field of our Galaxy (see e.g. Vaillancourt 2008.. Hildebrand et al. 2000)).," An accurate measurement of dust polarization at different wavelengths is important to better understand the nature and structure of dust grains and to probe the magnetic field of our Galaxy (see e.g. Vaillancourt \cite{Vail09}, Hildebrand et al. \cite{Hild00}) )."
 Moreover. polarized dust emission is an important ttamminnatting foreground in precision measurements of the polarization of the cosmic microwave background (CMB). the current ambitious target of CMB measurements.," Moreover, polarized dust emission is an important ting foreground in precision measurements of the polarization of the cosmic microwave background (CMB), the current ambitious target of CMB measurements."
 Measuring it at wavelengths where dust polarization is dominant. 15 mandatory in order to correct the cossmolloggiceal signal at the level required to measure B-modes (see e.g. Lazarian et al. 2009))., Measuring it at wavelengths where dust polarization is dominant is mandatory in order to correct the cal signal at the level required to measure B-modes (see e.g. Lazarian et al. \cite{Laza09}) ).
 Diffuse Galactic sources act as a foreground. mimicking the cosmological polarized signal (Hanany and Rosenkranz 2003:: Ponthieu and Martin 2006)): at frequencies above GHz interstellar dust is the dominant foreground.," Diffuse Galactic sources act as a foreground, mimicking the cosmological polarized signal (Hanany and Rosenkranz \cite{Hanany03a}; ; Ponthieu and Martin \cite{Ponthieu06}) ): at frequencies above $\,$ GHz interstellar dust is the dominant foreground."
 Therefore. an accurate knowlledgge of the polarization of these sources Is necessary to perform a precise measurement of the polarized cosmological siggnal.," Therefore, an accurate ge of the polarization of these sources is necessary to perform a precise measurement of the polarized cosmological gnal."
 This is made difficult by the small amplitude of the polarized signal of these sources. and by its strong angular and wavelength dependence (Tucci et al. 2003).," This is made difficult by the small amplitude of the polarized signal of these sources, and by its strong angular and wavelength dependence (Tucci et al. \cite{Tucci05}) )."
 Several experiments are planned to measure dust polarization at high galactie latitudes., Several experiments are planned to measure dust polarization at high galactic latitudes.
 The Planck satellite (Tauber et al. 2010)), The Planck satellite (Tauber et al. \cite{Taub10}) )
 and in particular the High Frequency Instrument (Lamarre et al. 2010)), and in particular the High Frequency Instrument (Lamarre et al. \cite{Lama10}) )
 is performing a shallow whole-sky survey of dust polarization at frequencies up to GHz.," is performing a shallow whole-sky survey of dust polarization at frequencies up to $\,$ GHz."
 Here. the signal due to interstellar dust at high galactic latitudes is roughly similar to the level of CMB anisotropy (Μας! et al. 2001:2006..," Here, the signal due to interstellar dust at high galactic latitudes is roughly similar to the level of CMB anisotropy (Masi et al. \cite{Masi01, Masi06},"
 Ponthieu et al. 2005))., Ponthieu et al. \cite{Ponthieu05}) ).
 Higher frequency surveys are more sensitive to the polarization signal of interstellar dust (see Fig. 1))., Higher frequency surveys are more sensitive to the polarization signal of interstellar dust (see Fig. \ref{fig:0}) ).
" PILOT is a stratospheric balloon borne experiment for the measurement of the polarization of the continuum emission in the diffuse interstellar medium at frequencies around 545 and GHz. with a resolution of 3.29% and 1.44"". respectively (Bernard et al. 2007))."," PILOT is a stratospheric balloon borne experiment for the measurement of the polarization of the continuum emission in the diffuse interstellar medium at frequencies around 545 and $\,$ GHz, with a resolution of $3.29'$ and $1.44'$, respectively (Bernard et al. \cite{Bernard07}) )."
 The experiment is optimized to perform surveys of the polarized dust emission along the Galactic plane but also at high Galactic latitudes. where previous experiments have provided only upper limits (Benoit et al. 2004::," The experiment is optimized to perform surveys of the polarized dust emission along the Galactic plane but also at high Galactic latitudes, where previous experiments have provided only upper limits (Benoit et al. \cite{Benoit04};"
 Ponthieu et al. 2005))., Ponthieu et al. \cite{Ponthieu05}) ).
 PILOT will probe the large scale distribution of the galactic magnetic field and the alignment properties of the dust grams providing strong constraints for dust models. in particular via the dependence on frequency of the degree of polarization.," PILOT will probe the large scale distribution of the galactic magnetic field and the alignment properties of the dust grains providing strong constraints for dust models, in particular via the dependence on frequency of the degree of polarization."
 BLAST-pol. the polarization sensitive version of the very successful BLAST balloon telescope. will also search. for polarized dust emission in the galactic plane in the southern hemisphere (Marsden et al. 2008..," BLAST-pol, the polarization sensitive version of the very successful BLAST balloon telescope, will also search for polarized dust emission in the galactic plane in the southern hemisphere (Marsden et al. \cite{Marsden08},"
 Fissel et al. 2010))., Fissel et al. \cite{Fissel10}) ).
 SPIDER (Crill et al. 2008)), SPIDER (Crill et al. \cite{Crill08}) )
 is an ambitious balloon-borne instrument aimed at the measurement of the polarization of the CMB by means of large arrays of polarization sensitive detectors., is an ambitious balloon-borne instrument aimed at the measurement of the polarization of the CMB by means of large arrays of polarization sensitive detectors.
 Working in the stratosphere. SPIDER can cover high frequencies (1n particular two bands at 225 GHz and 275 GHz). to monitor polarized emission from interstellar dust where 1t is not negligible with respect to the polarized component of the CMB.," Working in the stratosphere, SPIDER can cover high frequencies (in particular two bands at 225 GHz and 275 GHz), to monitor polarized emission from interstellar dust where it is not negligible with respect to the polarized component of the CMB."
 EBEX is also a balloon-borne CMB polarimeter. aiming at smaller angular scales and with a dust monitor at 410 GHz (Oxley et al. 2004)).," EBEX is also a balloon-borne CMB polarimeter, aiming at smaller angular scales and with a dust monitor at 410 GHz (Oxley et al. \cite{Oxle04}) )."
 There are different instrumental techniques to detect faint polarized signals., There are different instrumental techniques to detect faint polarized signals.
 Polarizzattion-sensitive bolometers (PSB. Jones et al. 2003))," tion-sensitive bolometers (PSB, Jones et al. \cite{Jones03}) )"
 have been used in B03 (Masi et al. 2006)), have been used in B03 (Masi et al. \cite{Masi06}) )
 and Planck (Delabrouille and Kaplan 2001))., and Planck (Delabrouille and Kaplan \cite{Delabrouille01}) ).
 Orthomode transducers have been used in WMAP (Jarosiket al. 2003)), Orthomode transducers have been used in WMAP (Jarosiket al. \cite{Jarosik03}) )
 and other coherent systems., and other coherent systems.
 Astack of a rotating birefringent crystal followed by a static polarizer implements a Stokes polarimeter. and has been used," A stack of a rotating birefringent crystal followed by a static polarizer implements a Stokes polarimeter, and has been used"
The emission coefficients that appear in the previous equation include only line processes: in order to reproduce the observed Q/7 profile. we need to add the contribution of the continuum.,"The emission coefficients that appear in the previous equation include only line processes: in order to reproduce the observed $Q/I$ profile, we need to add the contribution of the continuum."
" Assuming such contributions to be constant across the line we have where the superscripts ""c and ο recall that. the corresponding quantities refer. to continuum. and. line processes. respectively."," Assuming such contributions to be constant across the line we have where the superscripts $c$ ” and $\ell$ ” recall that the corresponding quantities refer to continuum and line processes, respectively."
 The quantities and ej will be considered as free parameters to be adjustedες in order to reproduce the observed polarization profile.," The quantities $\varepsilon_{X}^{\, \rm c}$ and $\varepsilon_I^{\,\rm c}$ will be considered as free parameters to be adjusted in order to reproduce the observed polarization profile."
 The effect of collisions will be neglected throughout this investigation., The effect of collisions will be neglected throughout this investigation.
 It should be observed that given the extreme weakness of the absorption features of this lithium doublet in the solar intensity spectrum. the optically thin slab model is expected to be a rather good approximation for the modelling of these lines.," It should be observed that given the extreme weakness of the absorption features of this lithium doublet in the solar intensity spectrum, the optically thin slab model is expected to be a rather good approximation for the modelling of these lines."
 Applying Eq. (7)).," Applying Eq. \ref{eq:QsuI_cont}) ),"
 and assuming the following values of the free parameters e;=200x 2. where ερ is the maximum value of ει in the wavelength range considered. and a Doppler width of 60 mÁ.. we obtain the two-peak &j.linear polarization profile shown in Figure 3..," and assuming the following values of the free parameters $\varepsilon_I^{\, \rm c}=200 \times \varepsilon_I^{\, \rm max}$ , where $\varepsilon_I^{\, \rm max}$ is the maximum value of $\varepsilon_I^{\, \ell}$ in the wavelength range considered, $\varepsilon_Q^{\, \rm c}=10^{-4} \times 
\varepsilon_I^{\, \rm c}$ , and a Doppler width of 60 we obtain the two-peak linear polarization profile shown in Figure \ref{fig:best_fit}."
 As can be observed in the left panel of Figure 4.. modifying the continuum contribution to the intensity we modify the amplitude of the two peaks. without changing the shape of the profile.," As can be observed in the left panel of Figure \ref{fig:cont-doppler}, modifying the continuum contribution to the intensity we modify the amplitude of the two peaks, without changing the shape of the profile."
" The value thàt we assumed for the continuum contribution (ej=200x appears to be a good choice. given the weakness of the &j"")line. and that it allows the peak falling at shorter wavelengths (hereafter referred to as the ""blue peak”) to reach the same amplitude (0.02%)) as the central peak of the Q/T profile observed by Stenfloetal.(2000) (see Figure 1))."," The value that we assumed for the continuum contribution $\varepsilon_I^{\, \rm c}=200 \times \varepsilon_I^{\, \rm max}$ ) appears to be a good choice, given the weakness of the line, and that it allows the peak falling at shorter wavelengths (hereafter referred to as the “blue peak”) to reach the same amplitude ) as the central peak of the $Q/I$ profile observed by \citet{Ste00} (see Figure \ref{fig:observation}) )."
" The value of e, has been adjusted in order to obtain in the far wings the same continuum polarization level as in the observation (0.01%)."," The value of $\varepsilon_Q^{\, \rm c}$ has been adjusted in order to obtain in the far wings the same continuum polarization level as in the observation )."
 Particularly interesting is the sensitivity of the theoretical O/T profile to the value of the Doppler width., Particularly interesting is the sensitivity of the theoretical $Q/I$ profile to the value of the Doppler width.
 The profile shown in Figure 3. has been obtained assuming a Doppler width of 60mA. which corresponds. neglecting microturbulent velocities. to a temperature of 3000 K. For the sake of simplicity. we consider the same Doppler width for the two isotopes. despite their mass difference (=14%)-.," The profile shown in Figure \ref{fig:best_fit} has been obtained assuming a Doppler width of 60, which corresponds, neglecting microturbulent velocities, to a temperature of 3000 K. For the sake of simplicity, we consider the same Doppler width for the two isotopes, despite their mass difference $\approx 14\%$."
. Profiles obtained assuming different values of the Doppler width are plotted in the right panel of Figure 4.., Profiles obtained assuming different values of the Doppler width are plotted in the right panel of Figure \ref{fig:cont-doppler}.
 Interestingly. we observe that the two-peak structure gradually disappears as the Doppler width is increased: the observation of a two-peak structure could thus provide precise information concerning the thermal properties of the thin atmospheric region where this weak lithium doublet is formed.," Interestingly, we observe that the two-peak structure gradually disappears as the Doppler width is increased: the observation of a two-peak structure could thus provide precise information concerning the thermal properties of the thin atmospheric region where this weak lithium doublet is formed."
 In the left panel of Figure 5.. the same Q// profile as in Figure 3 is plotted together with the profiles expected in the hypothetical cases where only °Li (short-dashed line) or only *Li Cong-dashed line) were present.," In the left panel of Figure \ref{fig:2is-hfs}, the same $Q/I$ profile as in Figure \ref{fig:best_fit} is plotted together with the profiles expected in the hypothetical cases where only $^6$ Li (short-dashed line) or only $^7$ Li (long-dashed line) were present."
 The polarization peak that is obtained when a single isotope is considered (either *Li or /Li) is. in both cases. due to the corresponding D» line.," The polarization peak that is obtained when a single isotope is considered (either $^6$ Li or $^7$ Li) is, in both cases, due to the corresponding $_2$ line."
" Indeed. the polarization signals produced by the D, lines. both of *Li and of ""Li t are several orders of magnitude smaller. and cannot be appreciated on this plot."," Indeed, the polarization signals produced by the $_1$ lines, both of $^6$ Li and of $^7$ Li, are several orders of magnitude smaller, and cannot be appreciated on this plot."
 It is then clear that the physical origin of the two-peak structure of the Q/7 profile that we have obtained within our modelling assumption lies in the isotopic shift between the two lithium isotopes: the two peaks are nothing else but the signals produced by the D» lines of “Lit (blue peak) and of Li (red peak). which fall at different wavelengths because of the isotopic shift. and which are weighted by the isotopic abundances.," It is then clear that the physical origin of the two-peak structure of the $Q/I$ profile that we have obtained within our modelling assumption lies in the isotopic shift between the two lithium isotopes: the two peaks are nothing else but the signals produced by the $_2$ lines of $^7$ Li (blue peak) and of $^6$ Li (red peak), which fall at different wavelengths because of the isotopic shift, and which are weighted by the isotopic abundances."
 It is also interesting to observe that the signals due to the D> lines of the two Isotopes. as found 1n the hypothetical cases where only °Li or only ‘Li were present. do not have the same amplitude (see the left panel of Figure 5)).," It is also interesting to observe that the signals due to the $_2$ lines of the two isotopes, as found in the hypothetical cases where only $^6$ Li or only $^7$ Li were present, do not have the same amplitude (see the left panel of Figure \ref{fig:2is-hfs}) )."
 This is due to the fact the two isotopes have different HFS (in particular different nuclear spin quantum numbers. and different HFS constants. producing different splittings among the various F- levels. so that °Li is less depolarized by HFS than Li).," This is due to the fact the two isotopes have different HFS (in particular different nuclear spin quantum numbers, and different HFS constants, producing different splittings among the various $F$ -levels, so that $^6$ Li is less depolarized by HFS than $^7$ Li)."
 As a consequence. the relative amplitude of the two peaks that we have found in the resulting Q// profile does not reflect only the different abundances of the two lithium isotopes. but also their different HFS.," As a consequence, the relative amplitude of the two peaks that we have found in the resulting $Q/I$ profile does not reflect only the different abundances of the two lithium isotopes, but also their different HFS."
 Therefore. the presence of HFS in the two isotopes has to be taken into account in order to obtain the correct relative amplitude between the two peaks.," Therefore, the presence of HFS in the two isotopes has to be taken into account in order to obtain the correct relative amplitude between the two peaks."
 This is clearly shown in the right panel of Figure 5.. where the profiles obtained including (solid line) and neglecting (dashed line) HFS are plotted.," This is clearly shown in the right panel of Figure \ref{fig:2is-hfs}, where the profiles obtained including (solid line) and neglecting (dashed line) HFS are plotted."
" Although the D, and D» lines (both of ?Li and ""Li 0 are very close to each other. so that quantum interferences between the upper levelsof the corresponding transitions are not negligible. it is important to point out that. becauseof the weakness of these lines. the continuum level is rapidly reached both in the intensity spectrum.and in the Q/7 spectrum. and the spectral signatures of such coherences are"," Although the $_1$ and $_2$ lines (both of $^6$ Li and $^7$ Li ) are very close to each other, so that quantum interferences between the upper levelsof the corresponding transitions are not negligible, it is important to point out that, becauseof the weakness of these lines, the continuum level is rapidly reached both in the intensity spectrum,and in the $Q/I$ spectrum, and the spectral signatures of such coherences are"
progenitor stars (Woosley Weaver 1985]).,progenitor stars (Woosley Weaver \cite{woo85}) ).
 Iu both nodels the column deusity. is. about 2 aud the asma temperature is 0.6 keV. For the BS model we get a total N-ray huuinositv of ere +. for the RS nodel it is eve 1.," In both models the column density is about $^{-2}$ and the plasma temperature is 0.6 keV. For the BS model we get a total X-ray luminosity of erg $^{-1}$, for the RS model it is erg $^{-1}$."
 RS models with differcut ügher metallicities vield unacceptable fits., RS models with different higher metallicities yield unacceptable fits.
 The fit of the X-ray spectimm with a two-component LOC (RS|PO) is oulv slightly better than the fits., The fit of the X-ray spectrum with a two-component model (RS+PO) is only slightly better than the fits.
 Nevertheless. from the points mentioned ivove and the plivsical victure discussed m Sect.," Nevertheless, from the points mentioned above and the physical picture discussed in Sect."
 1. this model serves as the best explanation for the observed soft X-ray cussion., \ref{discussion} this model serves as the best explanation for the observed soft X-ray emission.
" Hydrosen colin density (N1,23.:3 ?) and power-law spectral dex (V=2.6) lie within the expected range (as discussed for the single power Luv above}.", Hydrogen column density $N_\mathrm{H}$ $^{-2}$ ) and power-law spectral index $\Gamma$ =2.6) lie within the expected range (as discussed for the single power law above).
 The plasina temperature of 0.3 keV fits with the observed values of other galaxies (0.8. NGC 253: Forbes et al. 1999: , The plasma temperature of 0.3 keV fits with the observed values of other galaxies (e.g. NGC 253: Forbes et al. \cite{for99}; ;
NCC 1808: Junkes et al. 1995))., NGC 1808: Junkes et al. \cite{jun95}) ).
 The total keV. Iuninositv for this model amounts to creyes with contribution from the RS compoucut., The total 0.1--2.4 keV luminosity for this model amounts to erg $^{-1}$ with contribution from the RS component.
 The spectral fit together with the residuals is plotted iu Fie. 6.., The spectral fit together with the residuals is plotted in Fig. \ref{rspomodel}.
 From the quality of the spectral fits alone there is no significance for favoring a sinele-component model or a combination of two compoucuts., From the quality of the spectral fits alone there is no significance for favoring a single-component model or a combination of two components.
 Due to the lack of any spatial information in the PSPC image there is no possibility to distinguish: between different spatial and spectral components simnultaucously., Due to the lack of any spatial information in the PSPC image there is no possibility to distinguish between different spatial and spectral components simultaneously.
 So oulv the combined information from the PSPC ancl IIRI data allows a more detailed interpretation of ιο ταν results., So only the combined information from the PSPC and HRI data allows a more detailed interpretation of the X-ray results.
 Several points speak against the single PO commpoucut model. as discussed in Sect. 3.2..," Several points speak against the single PO component model, as discussed in Sect. \ref{specfit}."
 A more likely scenario is a composition of several ciffercut emission sources. like au active nucleus; ΗΝΤΝDs. aud supernova roiinauts (SNRs).," A more likely scenario is a composition of several different emission sources, like an active nucleus, HMXBs, and supernova remnants (SNRs)."
 Iu the following we will therefore discuss acomposite, In the following we will therefore discuss acomposite
using the 6 element Institut de Radioastronomie Milliméttrique (IRAM) Plateau de Bure Interferometer.,using the 6 element Institut de Radioastronomie Milliméttrique (IRAM) Plateau de Bure Interferometer.
 The receivers were tuned to the para-H}8O 311-220 transition at 203.407498 GHz (1.47 mm).," The receivers were tuned to the $_2^{18}$ O $3_{1,3}-2_{2,0}$ transition at 203.407498 GHz (1.47 mm)."
 The correlator was set up with one unit with a bandwidth of 36 MHz (53 s~')) centered on this frequency providing a spectral resolution on 460 channels of 0.078 MHz (0.11 s!))., The correlator was set up with one unit with a bandwidth of 36 MHz (53 ) centered on this frequency providing a spectral resolution on 460 channels of 0.078 MHz (0.11 ).
 The source was observed in two configurations: in the C configuration on 02 December 2008 and in the B configuration on 11 and 13 January 2009., The source was observed in two configurations: in the C configuration on 02 December 2008 and in the B configuration on 11 and 13 January 2009.
 About 11 hours were spent in each configuration (including time used on gain calibrators etc.)., About 11 hours were spent in each configuration (including time used on gain calibrators etc.).
 When combined these two configurations cover baselines with lengths from 17.8 to 452 m (12 to 308 ka)., When combined these two configurations cover baselines with lengths from 17.8 to 452 m (12 to 308 $\lambda$ ).
 The data were calibrated and imaged using the CLIC and MAPPING packages from the IRAM GILDAS software., The data were calibrated and imaged using the CLIC and MAPPING packages from the IRAM GILDAS software.
" The calibration followed the standard approach: the absolute flux calibration was established through observations of MWC 349, the bandpass by observations of the strong quasar 3c454.3 and the complex gains by regular observations of the nearby quasar J0336--323 (approximately 0.8 Jy at 1.45 mm)."," The calibration followed the standard approach: the absolute flux calibration was established through observations of MWC 349, the bandpass by observations of the strong quasar 3c454.3 and the complex gains by regular observations of the nearby quasar J0336+323 (approximately 0.8 Jy at 1.45 mm)."
 Integrations with clearly deviating amplitudes and/or phases were flagged and the continuum was subtracted prior to Fourier transformation of the line data., Integrations with clearly deviating amplitudes and/or phases were flagged and the continuum was subtracted prior to Fourier transformation of the line data.
" With natural weighting the resulting beam size is 0.67""x at a position angle of 36.7°; the field of view is ((HPBW) at 1.45 mm.", With natural weighting the resulting beam size is $\times$ at a position angle of $^\circ$; the field of view is (HPBW) at 1.45 mm.
" The resulting RMS noise level is 11 mJy beam""! channel""! for the line data using natural weighting.", The resulting RMS noise level is 11 mJy $^{-1}$ $^{-1}$ for the line data using natural weighting.
 The continuum sensitivity is limited by the dynamical range of the interferometer and the resulting RMS noise level is a few mJy beam|., The continuum sensitivity is limited by the dynamical range of the interferometer and the resulting RMS noise level is a few mJy $^{-1}$.
 Fig., Fig.
 1. shows the continuum image of the observed region around IRAS4B. As seen both IRASAB and its nearby companion IRAS4B’ are clearly detected in the image., \ref{iras4b_cont} shows the continuum image of the observed region around IRAS4B. As seen both IRAS4B and its nearby companion $'$ are clearly detected in the image.
 Table 1 lists the results of elliptical Gaussian fits to the two sources: both are resolved with fluxes in agreement with the results from assuming that it has its origin in thermal dust continuum emission with Ενοv with 3.," Table \ref{cont_table}
 lists the results of elliptical Gaussian fits to the two sources: both are resolved with fluxes in agreement with the results from \cite{prosacpaper} assuming that it has its origin in thermal dust continuum emission with $F_\nu\propto \nu^{\alpha}$ with $\alpha\approx 2.5-3$ ."
" The continuum peak is clearly offset by 5—7"" from the emission at 3.6—24 um seen in the Spitzer Space Telescope images of IRAS4B; an indication that the Spitzer emission has its origin in material heated by the protostellar outflow even at 24 um (see also Jorgensenetal.2007b and Fig.", The continuum peak is clearly offset by $''$ from the emission at 3.6–24 $\mu$ m seen in the Spitzer Space Telescope images of IRAS4B; an indication that the Spitzer emission has its origin in material heated by the protostellar outflow even at 24 $\mu$ m (see also \citealt{scubaspitz} and Fig.
 2 of Allenetal. 2007))., 2 of \citealt{allenppv}) ).
 Fig., Fig.
 2 shows the spectrum toward the continuum peak of IRASAB. A number of lines are clearly detected as listed in Table 3 — including the targeted Hj*0 345-255 line.," \ref{spectrum} shows the spectrum toward the continuum peak of IRAS4B. A number of lines are clearly detected as listed in Table \ref{line_id} – including the targeted $_2^{18}$ O $3_{1,3}-2_{2,0}$ line."
 For the line identification we used, For the line identification we used
despite later observations of the svstem which support larger values (e.g. Hvnes et 11998).,despite later observations of the system which support larger values (e.g. Hynes et 1998).
 In addition. we must also consider the effects. of irradiation-driven circulation over the surface of. the secondary.," In addition, we must also consider the effects of irradiation-driven circulation over the surface of the secondary."
 “Phe transfer of heated material [from the irradiated regions towards the inner Lagrangian point. and therefore within the discs shadow. would produce similar consequences for the radial velocity curve as a small-aneled disc.," The transfer of heated material from the irradiated regions towards the inner Lagrangian point, and therefore within the disc's shadow, would produce similar consequences for the radial velocity curve as a small-angled disc."
 Furthermore. the obvious asvmimetry of the data around: orbital phase 0.5 (when the illuminated hemisphere is directed towards the line-of-sight) possibly mav be explained by non-axially svinmetric circulation induced by the Coriolis force.," Furthermore, the obvious asymmetry of the data around orbital phase 0.5 (when the illuminated hemisphere is directed towards the line-of-sight) possibly may be explained by non-axially symmetric circulation induced by the Coriolis force."
 Although a detailed discussion is bevond the scope of this paper. it has been shown that such circulation elfects are significant.," Although a detailed discussion is beyond the scope of this paper, it has been shown that such circulation effects are significant."
 For example. Schandl. AMever-LHofmeister AMever (1997). used horizontal heat transfer in their modelling of the visual lighteurve of CAL ST: also. the analysis of the optical lighteurve ofLIZ Herculis. bv Wippenhahn Thomas (1979). required. circulation to explain the shape of the lighteurve at minimum.," For example, Schandl, Meyer-Hofmeister Meyer (1997) used horizontal heat transfer in their modelling of the visual lightcurve of CAL 87; also, the analysis of the optical lightcurve of HZ Herculis, by Kippenhahn Thomas (1979), required circulation to explain the shape of the lightcurve at minimum."
 Another unique feature in the radial velocity curve of X-Rav Nova Seo 1904 is the high heliocentric raclial velocity of approximately -1505., Another unique feature in the radial velocity curve of X-Ray Nova Sco 1994 is the high heliocentric radial velocity of approximately -150.
 After correction for the peculiar motion of the Sun ancl dillerential. Galactic rotation. the magnitude of the space velocity of X-Ray Nova Sco 1994 stands out as being being much higher than anv other dynamically identified: Galactic black hole candidate.," After correction for the peculiar motion of the Sun and differential Galactic rotation, the magnitude of the space velocity of X-Ray Nova Sco 1994 stands out as being being much higher than any other dynamically identified Galactic black hole candidate."
 Brandt. Podsiadlowski Sigurdsson (1995) give an explanation of the high space velocity. of X-Ray Nova Sco 1994 in terms of a delayed. black bole creation. which appears to favour the production of a relatively low black hole mass.," Brandt, Podsiadlowski Sigurdsson (1995) give an explanation of the high space velocity of X-Ray Nova Sco 1994 in terms of a delayed black hole creation, which appears to favour the production of a relatively low black hole mass."
 In this scenario. the initial collapse leads to the formation of a neutron star. allowing for a kick normally. associated with a neutron star formation.," In this scenario, the initial collapse leads to the formation of a neutron star, allowing for a kick normally associated with a neutron star formation."
 The neutron star is then converted into a black hole due either to subsequent accretion of matter or a phase transition in the compact object., The neutron star is then converted into a black hole due either to subsequent accretion of matter or a phase transition in the compact object.
 According to the. stripped-giant models for the companion star (lxing 1993: Brandt. Pocdsiadlowski Sigurdsson 1995) the maximum mass of the secondary. is 2.3M.," According to the stripped-giant models for the companion star (King 1993; Brandt, Podsiadlowski Sigurdsson 1995) the maximum mass of the secondary is 2.3."
.. Our lower limit for the secondary star mass of Alo>»L4 implies that a maximum of ~0.9 hhas therefore been available for accretion onto the black hole., Our lower limit for the secondary star mass of $M_{2}>1.4$ implies that a maximum of $\sim0.9$ has therefore been available for accretion onto the black hole.
 Since in the phase transition scenario. the black hole would initially be formed with a relatively low mass (<Al... Brown Bethe 1994). there is insulliciont matter available to form the observed lower limit for the compact object of 4.1M.," Since in the phase transition scenario, the black hole would initially be formed with a relatively low mass $<2$, Brown Bethe 1994), there is insufficient matter available to form the observed lower limit for the compact object of 4.1."
.. The alternative hypothesis in which the black hole in X-Ray Nova Sco is formed via an intermediate neutron star stage. and then converted. into a black hole by subsequent. accretion of supernova material. therefore appears more Consistent with our mass limits.," The alternative hypothesis in which the black hole in X-Ray Nova Sco is formed via an intermediate neutron star stage, and then converted into a black hole by subsequent accretion of supernova material, therefore appears more consistent with our mass limits."
 Llowever. other possible scenarios. such as a prompt black hole formation with an associated Blaauw-Bocrsma kick (see Drandt Poclsiacllowski 1995). cannot be ruled out at this stage.," However, other possible scenarios, such as a prompt black hole formation with an associated Blaauw-Boersma kick (see Brandt Podsiadlowski 1995), cannot be ruled out at this stage."
 We have reanalysed the published outburst absorption line racial velocity data of X-Ray Nova Sco 1994., We have reanalysed the published outburst absorption line radial velocity data of X-Ray Nova Sco 1994.
 We find that as the N-ray source was active during the observations. one has to moclel the ellects of X-ray irradiation of the secondary star when interpreting the radial velocity curve. since the irradiation will alfect. the strength of the absorption lines.," We find that as the X-ray source was active during the observations, one has to model the effects of X-ray irradiation of the secondary star when interpreting the radial velocity curve, since the irradiation will affect the strength of the absorption lines."
 The observed outburst radial velocity data is fitted using the X-ray heating mocel and 90 per cent confidence solutions are obtained in the (As.9) plane.," The observed outburst radial velocity data is fitted using the X-ray heating model and 90 per cent confidence solutions are obtained in the $K_{2},q$ ) plane."
 Assuming a binary mass ratio of 3.0 and the inclination range of 63.77 to 70.77. we derive imits on the masses of the binary components: 4.16.7 aand 1.4«Mo<2.2 for the black hole and secondary star. respectively (90 per cent confidence).," Assuming a binary mass ratio of 3.0 and the inclination range of $63.7^\circ$ to $70.7^\circ$, we derive limits on the masses of the binary components: $4.1<M_{1}<6.7$ and $1.4<M_{2}<2.2$ for the black hole and secondary star, respectively (90 per cent confidence)."
 This lower limit for the black hole mass is consistent with the idea that it was formed as the result of he post-supernova collapse of a neutron star., This lower limit for the black hole mass is consistent with the idea that it was formed as the result of the post-supernova collapse of a neutron star.
 We urge. future spectroscopic observations of X-Ray ova Sco 1994 in.quieseenec.. which will enable the true racial velocity of the secondary star ancl also the binary mass ratio to be determined directly.," We urge future spectroscopic observations of X-Ray Nova Sco 1994 in, which will enable the true radial velocity of the secondary star and also the binary mass ratio to be determined directly."
 Phese parameters are crucial in establishing the true masses of the binary components., These parameters are crucial in establishing the true masses of the binary components.
 We would like to thank Jerry Orosz for providing the racial velocity cata., We would like to thank Jerry Orosz for providing the radial velocity data.
at hieh redshift.,at high redshift.
 Repeating this uecasuremenut for similar sinulatious reveals that high redshift (2>10) massive inflow on this scale is a common occurrence for galaxies iu this mass range (a few «101 M at +=4)., Repeating this measurement for similar simulations reveals that high redshift $z>10$ ) massive inflow on this scale is a common occurrence for galaxies in this mass range (a few $\times 10^{11}$ $_\odot$ at $z = 5$ ).
 One might wonder whether the rapid contraction of a cluster by a factor of a few would lead to a high rate of interactions of single objects with hard binaries aud rence to re-expausion of the cluster due to energy iuput., One might wonder whether the rapid contraction of a cluster by a factor of a few would lead to a high rate of interactions of single objects with hard binaries and hence to re-expansion of the cluster due to energy input.
 The degree to which this cau happen clearly depends ou the fraction of stars or black holes im lard binaries hat are not so hard that they meree or collide rapidly., The degree to which this can happen clearly depends on the fraction of stars or black holes in hard binaries that are not so hard that they merge or collide rapidly.
 This fraction. iu turn. depends on the velocity dispersion: uegher velocity dispersion neans fewer hard binarics and less binary binding cucreyv that cau poteutially be apped to hold off core collapse.," This fraction, in turn, depends on the velocity dispersion; higher velocity dispersion means fewer hard binaries and less binary binding energy that can potentially be tapped to hold off core collapse."
 We also note that close hree-body interactions between objects of comparable nass vield a thermal distribution of ecceutricities., We also note that close three-body interactions between objects of comparable mass yield a thermal distribution of eccentricities.
 The eccentricities that eive pericenter distances low cnough or collisious (0.01 AU for solartype stars) or fast uereer by eravitational radiation (also ~0.01 AU for 10AL. black holes to merece in a few million years or ess) will destroy the binary., The eccentricities that give pericenter distances low enough for collisions $\sim 0.01$ AU for solar-type stars) or fast merger by gravitational radiation (also $\sim 0.01$ AU for $\sim 10~M_\odot$ black holes to merge in a few million years or less) will destroy the binary.
 If the hard-soft boundary isa<l1 AU (corresponding to a velocity dispersion a~30 lau | for sobu-tvpe stars or 6100 kins | oy LOAM. black holes). the available binding euergv in binaries is less than the binding energy m the singles. rence binary-sinele juteractious are incfiicicut at holding off core collapse.," If the hard-soft boundary is $a<1$ AU (corresponding to a velocity dispersion $\sigma\sim 30$ km $^{-1}$ for solar-type stars or $\sigma\sim 100$ km $^{-1}$ for $\sim
10~M_\odot$ black holes), the available binding energy in binaries is less than the binding energy in the singles, hence binary-single interactions are inefficient at holding off core collapse."
" We now consider the subsequent evolution of the stellar cluster. assunius it has received a significant influx of eas, and perhapsstars."," We now consider the subsequent evolution of the stellar cluster, assuming it has received a significant influx of gas, and perhapsstars."
" Begining with a uuclear stellar cluster of iiass AL.gaeo1. and halfauass radius n~1. the infall of10* M. of eas is likely to shrink the cluster by a factor of about ten leaving initial core deusities à~108 stars ο and velocity dispersious ονyy~3050,"," Beginning with a nuclear stellar cluster of mass $M_{\rm
c,6}æ\sim 1$, and half-mass radius $r_{\rm h} \sim 1$, the infall of $10^7$ $_{\odot}$ of gas is likely to shrink the cluster by a factor of about ten leaving initial core densities $n \sim 10^8$ stars $^{-3}$ and velocity dispersions $v_{\infty,10} \sim 30 -
50$."
 Several thines will happen to this cluster., Several things will happen to this cluster.
 Wider binaries Q@vlich previously heated the cluster aud prevented core collapse) will now be soft. as the velocity dispersion has increased. and will be broken up.," Wider binaries (which previously heated the cluster and prevented core collapse) will now be soft, as the velocity dispersion has increased, and will be broken up."
 Any remaining binarics will quickly iuerge via gravitational radiatiou spiral or collisions., Any remaining binaries will quickly merge via gravitational radiation inspiral or collisions.
 Merger products are likely to be retained with the cluster owing to the increased escape speed which resulted from the eas infall., Merger products are likely to be retained with the cluster owing to the increased escape speed which resulted from the gas infall.
 It is possible that some merecr products will iu turi exchange into other binaries which then merece., It is possible that some merger products will in turn exchange into other binaries which then merge.
 This process may even be repeated., This process may even be repeated.
 After a short time. the cluster core will contain black holes having a broader range of masses thau previously.," After a short time, the cluster core will contain black holes having a broader range of masses than previously."
 Iu addition. all binaries will either have been broken up or have merged.," In addition, all binaries will either have been broken up or have merged."
 Thus the cluster core will lose its source of energy. aud will begin to contract approaching core collapse as it transfers energy to the cluster halo via two-body scattering.," Thus the cluster core will lose its source of energy, and will begin to contract approaching core collapse as it transfers energy to the cluster halo via two-body scattering."
 As the core density increases. runasvay inerecrs between any non-conipact stars within the cluster core becomes possible (Quinlan&Shapiro 1990).," As the core density increases, runaway mergers between any non-compact stars within the cluster core becomes possible \citep{Quinlan90}."
. Such collisious may produce more massive stars which in turn will add to the population of black holes and neutron stars., Such collisions may produce more massive stars which in turn will add to the population of black holes and neutron stars.
" At lieh densities, extremely close encounters occur between two compact objects where binary formation via eravitational wave emission is possible."," At high densities, extremely close encounters occur between two compact objects where binary formation via gravitational wave emission is possible."
 The timescale for capture is given by (Quinlan&Shapiro1989) where pao is the reduced mass of the two compac objects capturing each other Gu units of 10 AL. and Adjyy is their total mass (in units of 100. ADL).," The timescale for capture is given by \citep{Quinlan89}
 where $\mu_{10}$ is the reduced mass of the two compact objects capturing each other (in units of 10 $_\odot$ and $M_{100}$ is their total mass (in units of 100 $_\odot$ )."
 We see that when the core reaches densities »~1013 stars ?. which at a velocity dispersion e~300lis contains few<104AL... the formation of BU-BID binaries via gravitational wave enuüssion becomes possible on a timescale <AIvr.," We see that when the core reaches densities $n \sim 10^{12}$ stars $^{-3}$, which at a velocity dispersion $v\sim 300~{\rm km~s}^{-1}$ contains ${\rm
few}\times 10^4~M_\odot$, the formation of BH-BH binaries via gravitational wave emission becomes possible on a timescale $\ltorder$ Myr."
 These binaries will lave verv shor inspiral timescales., These binaries will have very short inspiral timescales.
 We will thus have a second phase of black hole binary mcrecrs., We will thus have a second phase of black hole binary mergers.
 The black holes produced via such mergers lave greater mass. and given tha tho timescale for merecrs scales as AL122) hügberanass black holes have a ereater probability of ruergcr aud the process will run away.," The black holes produced via such mergers have greater mass, and given that the timescale for mergers scales as $M^{-12/7}$, higher-mass black holes have a greater probability of merger and the process will run away."
 Such a phase may well collec most of the mass of the stellar black holes iuto a single black hole. which would have several percent of the original stellar mass of the svsteu for typical initial mass functions.," Such a phase may well collect most of the mass of the stellar black holes into a single black hole, which would have several percent of the original stellar mass of the system for typical initial mass functions."
 Thus the mass of the single black hole could be ~10? ML. for situations similar to that shown in Figure 2. (assmniusg most iufalliug eas reaches the centro), Thus the mass of the single black hole could be $\sim 10^5$ $_\odot$ for situations similar to that shown in Figure \ref{fig:dmb2} (assuming most infalling gas reaches the centre).
 Accretion of neutron stars aud white dwarfs. oth of which will be swallowed whole for M=10°AZ... could boost the mass by a factor of a few. depending ou he mass at which stellar interactions in the radius of influence of the black hole become an effective heating source for the cluster (6.9... Culletal.2008)).," Accretion of neutron stars and white dwarfs, both of which will be swallowed whole for $M\gtorder 10^5~M_\odot$ , could boost the mass by a factor of a few, depending on the mass at which stellar interactions in the radius of influence of the black hole become an effective heating source for the cluster (e.g., \citealt{Gill08}) )."
" Overall. he time elapsed. between mass infall to the production of the seed SMDBIT will be roughly 100 Myr (0.8. Quinlan&Shapiro 19903). hence producing a ~10AL. black tole within ~208LOO Myr after the begiuniusg of the universe,"," Overall, the time elapsed between mass infall to the production of the seed SMBH will be roughly 100 Myr (e.g., \citealt{Quinlan90}) ), hence producing a $\sim 10^5~M_\odot$ black hole within $\sim 300-400$ Myr after the beginning of the universe."
 Growth to ~10?AZ. via gas accretion could ren occur by a redshift 2~6 via standard. Eddiuston-nuited accretion onto moderately spinning black holes., Growth to $\sim 10^9~M_\odot$ via gas accretion could then occur by a redshift $z\sim 6$ via standard Eddington-limited accretion onto moderately spinning black holes.
 Iu sunny. we have proposed that clusters with initial structural paraiueters similar to current-dav nuclear clusters παν (1) form as part of carly hierarchical moreime. (2) accrete gas with a total mass comparable to or greater than that of the cluster at redshifts + 10. (3) contract as a result so that binary heating is iuncffectivo. (1) uudergo core collapse to a deusity high enough that stellarmass black holes merge. aud thus (5) have most of the iiass originally in stellarauass black holes collect in to a sinele black hole that could be ~10°AL. or larger.," In summary, we have proposed that clusters with initial structural parameters similar to current-day nuclear clusters may (1) form as part of early hierarchical merging, (2) accrete gas with a total mass comparable to or greater than that of the cluster at redshifts $z>10$ , (3) contract as a result so that binary heating is ineffective, (4) undergo core collapse to a density high enough that stellar-mass black holes merge, and thus (5) have most of the mass originally in stellar-mass black holes collect in to a single black hole that could be $\sim 10^5~M_\odot$ or larger."
 This black hole would therefore be a hiehauass seed that could comfortably grow to superlassive size bv the observed redshifts 2~ 6., This black hole would therefore be a high-mass seed that could comfortably grow to supermassive size by the observed redshifts $z\sim 6$ .
 We eratefully acknowledge the hospitality. of the Aspen Center for Physics., We gratefully acknowledge the hospitality of the Aspen Center for Physics.
 IBD was supported by the Swedish Research Council (eraut. 2008-L089)., MBD was supported by the Swedish Research Council (grant 2008-4089).
 ICAL was supported in part byNASA erant NNNOSATI20G. JAIB acknowledges NASA award NNNIOACSIC., MCM was supported in part byNASA grant NNX08AH29G. JMB acknowledges NASA award NNX10AC84G.
frequency for Alog is the same as for Alog«eg.,frequency for $\Delta \log A$ is the same as for $\Delta \log a_f$.
 Then we approximately set as where AlogPy is taken from the SN Ia reeion in Figure Al2 and the factor of 2/3 comes from the conversion between the period aud the separation., Then we approximately set as where $\Delta \log P_0$ is taken from the SN Ia region in Figure \ref{ztotreg100} and the factor of $2/3$ comes from the conversion between the period and the separation.
 Substituting AlogA= 0.6:2/3. Ag=2.6/L1891.2/5.60 0.37. Afy=LIS. Mg=5.60. we obtain myp.o.sο=0.0006.," Substituting $\Delta \log A = 0.6 \cdot 2/3$ $\Delta q = 2.6/4.48-1.2/5.60=0.37$ , $M_A = 4.48$, $M_B = 5.60$, we obtain $\nu_{{\rm WD},0.8-0.9} = 0.0006$."
 The SNe Ia rates for other WD iuass intervals are summarized in Table A5.., The SNe Ia rates for other WD mass intervals are summarized in Table \ref{tbl_realizaton_frequency}. .
 Then. the suuunatiou of SN Ta rates for three intervals (8O9AL.. 0.9LOA... and 1.0 LALAL..) gives pq=0.0017 vet. which is laree enough to explain the dominant part of the SN Ta rate in our Galaxy.," Then, the summation of SN Ia rates for three intervals $0.8-0.9 M_\odot$, $0.9-1.0 M_\odot$, and $1.0-1.1 M_\odot$ ) gives $\nu_{\rm RG}= 0.0017$ $^{-1}$, which is large enough to explain the dominant part of the SN Ia rate in our Galaxy."
 If we further include the WD mass range of Ahwpo-llLeal... which is not shown in Table Ah.. the realization frequency increases to vac=0.0022 |.," If we further include the WD mass range of $M_{\rm WD,0}=1.1-1.2 M_\odot$, which is not shown in Table \ref{tbl_realizaton_frequency}, the realization frequency increases to $\nu_{\rm RG}= 0.0022$ $^{-1}$."
 Tere. we have not shown the region for Wyyp.y—1.2.M. in Fieure Al2 because Webbiuk et al," Here, we have not shown the region for $M_{\rm WD,0}=1.2 M_\odot$ in Figure \ref{ztotreg100} because Webbink et al."
/s (1983) eiipirical forimla is valid for May<2.5.3.04. with a degenerate helimm core.,"'s (1983) empirical formula is valid for $M_{2,0} < 2.5-3.0 M_\odot$ with a degenerate helium core."
 The range of Mp)=1.2.M.. exceeds this liit.," The range of $M_{\rm WD,0}=1.2 M_\odot$ exceeds this limit."
 To examine the effect of the stripping parameter. gag. we estimate the realization frequency for jae=0.3 as summarized in Table AG..," To examine the effect of the stripping parameter, $\eta_{\rm eff}$, we estimate the realization frequency for $\eta_{\rm eff}=0.3$ as summarized in Table \ref{tbl_realization_0.33}."
 The parameter region shrinks to oue third in area compared with the case of jig=1. so that the realization frequency is reduced to about one third of jj4p=1 case. Lo. vag=0.0008 5|.," The parameter region shrinks to one third in area compared with the case of $\eta_{\rm eff}=1$, so that the realization frequency is reduced to about one third of $\eta_{\rm eff}=1$ case, i.e., $\nu_{\rm RG}= 0.0008$ $^{-1}$."
 For the WDMS progenitors. IIKNUS98 have fouud a new evolutionary path. which las not been taken into account in the previous works (e.g. Rappaportctal. 1995: DiStefano&Rappaport 1991: Yuneclsouetal. 1996: Yuugelsoun&Livio 1998)). aud estimated the realization frequeucy to be as large as my=0.001 bL," For the WD+MS progenitors, HKNU98 have found a new evolutionary path, which has not been taken into account in the previous works (e.g., \cite{rap94}; \cite{dis94}; \cite{yun96}; \cite{yun98}) ), and estimated the realization frequency to be as large as $\nu_{\rm MS}= 0.001$ $^{-1}$."
 We briefly follow their new evolutionary path in the following aud discuss the total rate of SN Ta explosions: If acLA. ο WD is descending from au ACB star. its zZoro-age lad sequence nass is TAL. by equation L1)) and the binary separation is larger than e;~1350HR. if the secondary mass is ~BAL...," We briefly follow their new evolutionary path in the following and discuss the total rate of SN Ia explosions: If a $\sim 1 M_\odot$ C+O WD is descending from an AGB star, its zero-age main sequence mass is $\sim 7 M_\odot$ by equation \ref{mass_relation_CO}) ) and the binary separation is larger than $a_i \sim 1350 ~R_\odot$ if the secondary mass is $\sim 2 M_\odot$."
" Its separation shriuks oa,~TOR. after the conuuion euvelope evolution with ocr=Ll (the secoudary mass of —2M .).", Its separation shrinks to $a_f \sim 70 ~R_\odot$ after the common envelope evolution with $\alpha_{\rm CE} = 1$ (the secondary mass of $\sim 2 M_\odot$ ).
 Theu the orbital xeriod becomes Ly~I0 d aud too long to become an SN Ia as seen iu Fieure ALL., Then the orbital period becomes $P_0 \sim 40$ d and too long to become an SN Ia as seen in Figure \ref{zams10}.
 Therefore. the C|O WD comes not yon an ACB star having the radius of equation CI16)) but roni a helium star whose hydrogeu-ich envelope has been stripped away in the first common cuvelope evolution (at he red-eilant phase with a helium core).," Therefore, the C+O WD comes not from an AGB star having the radius of equation \ref{radius_AGB_CO}) ) but from a helium star whose hydrogen-rich envelope has been stripped away in the first common envelope evolution (at the red-giant phase with a helium core)."
 Then. we follow he evolution of a binary consisting of a helium star aud a nain-sequence star.," Then, we follow the evolution of a binary consisting of a helium star and a main-sequence star."
 To estimate the decrease iu the separation after the conumnion euvelope phase. weuse the radius to helium core nass (RyAF a.) relation. which are taken from tables eiven by Bressan et al. (," To estimate the decrease in the separation after the common envelope phase, weuse the radius to helium core mass $R_1 - M_{\rm 1,He}$ ) relation, which are taken from tables given by Bressan et al. ("
1993).,1993).
 As an example. let us consider a pair of TAL..|2.542. with the initial separation of a;~506OORi.," As an example, let us consider a pair of $7 M_\odot + 2.5 M_\odot$ with the initial separation of $a_i \sim 50-600 R_\odot$."
 The binary evolves to SN Ia through the following stages: Therefore. the above pair of £M...|2.54L.. cau be a progenitor of SNe Ta if the initial separation is between 5O 150R.. which initiates a common cuvelope evolution at the helium core mass of Aji.=1.01211... corresponding to the initial orbital period of Pung=055 d for the WD]MS systems in Figure. All..," The binary evolves to SN Ia through the following stages: Therefore, the above pair of $7 M_\odot + 2.5 M_\odot$ can be a progenitor of SNe Ia if the initial separation is between $50-150 R_\odot$ , which initiates a common envelope evolution at the helium core mass of $M_{\rm 1,He}= 1.0-1.2 M_\odot$, corresponding to the initial orbital period of $P_{\rm orb,0}= 0.5-5$ d for the WD+MS systems in Figure \ref{zams10}."
 Iu this case. we have AloeA=logl50500.5 and q=2.5/70.36 in equation (13)).," In this case, we have $\Delta \log A= \log 150 - \log 50= 0.5$ and $q= 2.5/7= 0.36$ in equation \ref{realization_frequency}) )."
 Calculating 25 pairs of AL; aud. M»o;. TIVNU99 have obtained the parameter region of SN Ta explosion as Aq= rl AloeA20.5. M4=5.5. aud Mp=8.5.," Calculating 25 pairs of $M_{1,i}$ and $M_{2,i}$, HKNU99 have obtained the parameter region of SN Ia explosion as $\Delta q= 0.4$ , $\Delta \log A= 0.5$, $M_A= 5.5$, and $M_B= 8.5$."
 Substituting hese values iuto equation (13)). we obtain the SN Ia rate of mys=0.001 31;," Substituting these values into equation \ref{realization_frequency}) ), we obtain the SN Ia rate of $\nu_{\rm MS}= 0.001$ $^{-1}$."
 Thus TISNT99 have shown that the yequency of the WD|MS systems is about one third of he inferred rate in our Galaxy. which is ον larger than hat of Yuneelsou Livio’s (1998) estimation.," Thus HKNU99 have shown that the frequency of the WD+MS systems is about one third of the inferred rate in our Galaxy, which is much larger than that of Yungelson Livio's (1998) estimation."
 It should )o noted that Yuneclsou Livio (1998) have obtained the firth rate of d«102 xv| for their models 15 aud 16 x rolaxiug all the coustraiuts on the mass ratio of their ynary models. although it is not a realistic case.," It should be noted that Yungelson Livio (1998) have obtained the birth rate of $1 \times 10^{-3}$ $^{-1}$ for their models 15 and 16 by relaxing all the constraints on the mass ratio of their binary models, although it is not a realistic case."
" The orbital velocity of the WD|MS systems dis much aster than that for the WD|RC svstems. ie.Όλο,~lan 3 for Mwp=1LOAL.. aud AAs=2.0M.. at he zero-age main sequence."," The orbital velocity of the WD+MS systems is much faster than that for the WD+RG systems, i.e.,$a \Omega_{\rm orb} \sim 400$km $^{-1}$ for $M_{\rm WD}= 1.0 M_\odot$ and $M_{\rm MS}= 2.0 M_\odot$ at the zero-age main sequence."
" Then. the switching from equation (12)) to equation (37)) occurs at ο~1.5e0,5,4, kins 1."," Then, the switching from equation \ref{fast_wind_angular_momentum}) ) to equation \ref{slow_wind_angular_momentum}) ) occurs at $v \sim 1.5 a \Omega_{\rm orb} \sim 600$ km $^{-1}$ ."
 This means that the wind velocity las to be ‘aster than 00 kan1 in order to avoid the formation of a colmmon envelope., This means that the wind velocity has to be faster than $\sim 600$ km$^{-1}$ in order to avoid the formation of a common envelope.
 Otherwise. winds carry large specific aneular momentum and drastically shorten the separation oeuhauce the mass trausfer aud to eventually form a," Otherwise, winds carry large specific angular momentum and drastically shorten the separation toenhance the mass transfer and to eventually form a"
metallicity and thus a two to three times higher metallicity than the X-ray emitting plasma observed in the quasi-quiescent phase.,metallicity and thus a two to three times higher metallicity than the X-ray emitting plasma observed in the quasi-quiescent phase.
" While the absolute scale is only moderately constrained as outlined in refana,, the relative changes are quite robust and also found in two temperature models or the analysis the MOS spectra."," While the absolute scale is only moderately constrained as outlined in \\ref{ana}, the relative changes are quite robust and also found in two temperature models or the analysis the MOS spectra."
" In the final stages of the flare evolution the remaining material has cooled down to temperatures below 1 keV, exhibits pre-flare metallicity and becomes rather indistinguishable from the quasi-quiescent plasma."," In the final stages of the flare evolution the remaining material has cooled down to temperatures below 1 keV, exhibits pre-flare metallicity and becomes rather indistinguishable from the quasi-quiescent plasma."
" Overall, we find a clear correlation between plasma temperature and metallicity, indicating a strong modification of the chemical composition during the event."," Overall, we find a clear correlation between plasma temperature and metallicity, indicating a strong modification of the chemical composition during the event."
" Notably, the derived characteristics are quite typical for a stellar flare (seee.g.?).."," Notably, the derived characteristics are quite typical for a stellar flare \citep[see e.g.][]{gue04}."
" The high resolution RGS spectrum from IQ Aur has unfortunately only moderate signal to noise and only rather strong lines were clearly detected even in the total spectrum; see refrgs, where we show a flux-converted spectrum of the merged RGS detectors from a The individual line fluxes are measured from the total count spectrum and cross-checked by using a broad and a narrower The lower panel of refrgs shows thevu triplet with applied model, confirming that the forbidden line (22.1 À)) is significantly stronger than the intercombination line (21.8 A))."," The high resolution RGS spectrum from IQ Aur has unfortunately only moderate signal to noise and only rather strong lines were clearly detected even in the total spectrum; see \\ref{rgs}, where we show a flux-converted spectrum of the merged RGS detectors from a The individual line fluxes are measured from the total count spectrum and cross-checked by using a broad and a narrower The lower panel of \\ref{rgs} shows the triplet with applied model, confirming that the forbidden line (22.1 ) is significantly stronger than the intercombination line (21.8 )."
" The S/N ratio is quite poor, even in the total spectrum that includes the flare."," The S/N ratio is quite poor, even in the total spectrum that includes the flare."
" We also investigated the quasi-quiescent and flare phase separately, but except stronger emission from ""hotter lines’ (e.g. xvi) during the flare, a quantitative comparison suffers from the low S/N. The fluxes of emission lines used in the subsequent analysis as determined from the total spectrum are given in Table 3.."," We also investigated the quasi-quiescent and flare phase separately, but except stronger emission from 'hotter lines' (e.g. ) during the flare, a quantitative comparison suffers from the low S/N. The fluxes of emission lines used in the subsequent analysis as determined from the total spectrum are given in Table \ref{linetab}."
 They are corrected for an absorbing column of logNy=20 cm and we expect a flare contribution of roughly We utilize the He-like triplet to determine the distance of the X-ray emitting plasma to the surface of IQ Aur., They are corrected for an absorbing column of $\log N_{\rm H}=20$ $^{-2}$ and we expect a flare contribution of roughly We utilize the He-like triplet to determine the distance of the X-ray emitting plasma to the surface of IQ Aur.
Ovi traces the location of plasma at temperatures around 2 MK, traces the location of plasma at temperatures around 2 MK
required range of the dust extinction is far larecr than that estimated from the observed Balmer decrements.,required range of the dust extinction is far larger than that estimated from the observed Balmer decrements.
 This discrepancy nav be more siguificaut if the iron depletion factor is smaller than the value we adopt here (= 0.1)., This discrepancy may be more significant if the iron depletion factor is smaller than the value we adopt here (= 0.1).
 Iu cold ISAL the iron depletion factor reaches down to 0.01 (o.e.. Jenkins. Savage. Spitzer 1986: Cowie Soneaila 1986).," In cold ISM, the iron depletion factor reaches down to 0.01 (e.g., Jenkins, Savage, Spitzer 1986; Cowie Songaila 1986)."
 Although this laree discrepancy might be explained by introducing dust ervains which are located iu/around IIINERs selectively. this idea has the following two serious problems.," Although this large discrepancy might be explained by introducing dust grains which are located in/around HINERs selectively, this idea has the following two serious problems."
 First. to assume the existence of additional dust erains m the immer part of NLRs may conflict with the idea that it is relatively difficult for the dust grams to survive under hieh photoionizing ux. although we cannot exclude the possibility that the eraius could survive under the high flux for some situations.," First, to assume the existence of additional dust grains in the inner part of NLRs may conflict with the idea that it is relatively difficult for the dust grains to survive under high photoionizing flux, although we cannot exclude the possibility that the grains could survive under the high flux for some situations."
 We should recall the histograms presented in Figure 1l. which sugeest that the IIINERs are located at the iunerinost region in NLRs.," We should recall the histograms presented in Figure 1, which suggest that the HINERs are located at the innermost region in NLRs."
" second, the inferred amount of extiuctiou. 2 mag x-leX 10 mag. corresponds to Nye~(515)«10?! cmi7."," Second, the inferred amount of extinction, 3 mag $\lesssim A_V \lesssim$ 10 mag, corresponds to $N_{\rm H^0} \sim (5 - 15) \times 10^{21}$ $^{-2}$."
 This is far larger than the column density toward the uuclci of type AGNs. derived. froun N-rav spectral analysis (Nye<10?4 7: ees Reynolds 1997: Leighly 1999).," This is far larger than the column density toward the nuclei of type 1 AGNs, derived from X-ray spectral analysis $N_{\rm H^0} 
\lesssim 10^{21}$ $^{-2}$; e.g., Reynolds 1997; Leighly 1999)."
 These considerations suggest that the observed data are hard to be explained by the dusty models even if verv large amount of dust shiclding oulv the IIENERs selectively is introduced., These considerations suggest that the observed data are hard to be explained by the dusty models even if very large amount of dust shielding only the HINERs selectively is introduced.
 Therefore. we couclude that the iron in the IIINERs is not depleted outo dust erains siguificautlv.," Therefore, we conclude that the iron in the HINERs is not depleted onto dust grains significantly."
 This conclusion is consistent with the idea that IIINERs are located at ti0 Ποιος of NLRs aud thus are hidden bv dusty tori when secu from a edge-on view toward dusty tor. which is miplie by the histograms presented in Fieure 1 (sec also Pier Voit 1995: Muraviuua Tanienchi 1998a. 19958b: Darth et al.," This conclusion is consistent with the idea that HINERs are located at the innermost of NLRs and thus are hidden by dusty tori when seen from a edge-on view toward dusty tori, which is implied by the histograms presented in Figure 1 (see also Pier Voit 1995; Murayama Taniguchi 1998a, 1998b; Barth et al."
 1999: Tran et al., 1999; Tran et al.
 2000: Nagao et al., 2000; Nagao et al.
 2000b. 20012. 20015).," 2000b, 2001a, 2001b)."
 Considerimg that the TINERs ave located closer than the dust-sublinatiou radius (1.0... tuner radius of dustv tori). we cau understaud the absence of mterual dust erains in DINERS aud the ACN-type dependence of the visibility of high-ionization Cluission lines as preseuted in Figure 1. sinultaucouslv.," Considering that the HINERs are located closer than the dust-sublimation radius (i.e., inner radius of dusty tori), we can understand the absence of internal dust grains in HINERs and the AGN-type dependence of the visibility of high-ionization emission lines as presented in Figure 1, simultaneously."
 Finally. we meutiou the possible difference iu the flux ratio of [Fe vii] AG6087/[Ne. v|A3126 between type 1 aud type 2 ACNs.," Finally, we mention the possible difference in the flux ratio of [Fe $\lambda$ 6087/[Ne $\lambda$ 3426 between type 1 and type 2 AGNs."
 Although the dierence is not significant as presented iu Figure 2. the dierence in the extinction-COIYOCed fiux ratio may be larger than that iu the observed flux ratio since the dust extiucion is huger onu average in type 2 AGNs than type 1 AGNs (e... Dahari De Robertis 1988).," Although the difference is not significant as presented in Figure 2, the difference in the extinction-corrected flux ratio may be larger than that in the observed flux ratio since the dust extinction is larger on average in type 2 AGNs than type 1 AGNs (e.g., Dahari De Robertis 1988)."
 This difference nay sugeest that ion im spaially exteuded TINERs is depleted outo dust exaius., This difference may suggest that iron in spatially extended HINERs is depleted onto dust grains.
 Such spatially extended HENERs have been predicted theoreically (e.g... IKorista Ferland 1989: Fergusou et al.," Such spatially extended HINERs have been predicted theoretically (e.g., Korista Ferland 1989; Ferguson et al."
 1997b). cunissivitics of high-lonization cussion lines at the exeuded TINERs are expected to be low.," 1997b), emissivities of high-ionization emission lines at the extended HINERs are expected to be low."
 Indeed such extended ΕΠΗΣΤΗΣ are observationallv detected only in few Sevfer ealaxies (e.8.. Golev et al.," Indeed such extended HINERs are observationally detected only in few Seyfert galaxies (e.g., Golev et al."
 1995: Muravauua ct al., 1995; Murayama et al.
 1998: Nagao e al., 1998; Nagao et al.
 2000a: Nelson et al., 2000a; Nelson et al.
 2000: Ikracimier Crenshaw 20003., 2000; Kraemer Crenshaw 2000).
 Therefore. the effect of the presence of the spatially extended IIINERs on the observed £nx ratios of [Fe AGOST/[Ne. VIA3126 seems to be low.," Therefore, the effect of the presence of the spatially extended HINERs on the observed flux ratios of [Fe $\lambda$ 6087/[Ne $\lambda$ 3426 seems to be low."
" However. as sugeested by he histograms preseuted in Figure 1l. the innermost region in NLRs. where stroug [Fe ASF ancl [Ne VJA3126 arise. is hidden as for type 2 AGNs,"," However, as suggested by the histograms presented in Figure 1, the innermost region in NLRs, where strong [Fe $\lambda$ 6087 and [Ne $\lambda$ 3426 arise, is hidden as for type 2 AGNs."
ed Iu this case. the effect of dusty aud spatially exteu MINERS hav eniergse iu the observed flux ratios of [Fe AGOST/|Ne vJA3126.," In this case, the effect of dusty and spatially extended HINERs may emerge in the observed flux ratios of [Fe $\lambda$ 6087/[Ne $\lambda$ 3426."
 Tus idea is a highly speculative one since the observed clitHereuce in the flix ratio of [Fe AG0ST/[Ne vJA3126. beween the type 1 aud the type 2 ACNs is mareinal., This idea is a highly speculative one since the observed difference in the flux ratio of [Fe $\lambda$ 6087/[Ne $\lambda$ 3426 between the type 1 and the type 2 AGNs is marginal.
 Ii order to examine this idea observationally. spatial distributions of the flux ratio of Τον AGUNT/[Ne v]A3126 shuld be investigated.," In order to examine this idea observationally, spatial distributions of the flux ratio of [Fe $\lambda$ 6087/[Ne $\lambda$ 3426 should be investigated."
 Such studies cau male it clear how t16 iron depletion onto dust grains depends ou, Such studies can make it clear how the iron depletion onto dust grains depends on
5truciu BStrueiu Cnr Cnr Cnr Cnr σα1) clussilü ον CLUSSS ciusshbxlO 1102 ciutir clurd cutis clures 12pt,"5truein 8truein cmr8 cmr8 cmr8 cmr8 cmr10 cmssi10 cmss10 cmss8 cmssbx10 2 cmti7 cmr6 cmti8 cmr8 \def\ref{\par\noindent\hangindent 15pt}
 = 12pt"
5truciu BStrueiu Cnr Cnr Cnr Cnr σα1) clussilü ον CLUSSS ciusshbxlO 1102 ciutir clurd cutis clures 12pt-,"5truein 8truein cmr8 cmr8 cmr8 cmr8 cmr10 cmssi10 cmss10 cmss8 cmssbx10 2 cmti7 cmr6 cmti8 cmr8 \def\ref{\par\noindent\hangindent 15pt}
 = 12pt"
ο) is a function with the asviptotie behavior,$g(r/r_0)$ is a function with the asymptotic behavior
derived from (he accretion shock. peaks in the range of 0.1 to O.OLL. (solar luminosity). alihough the precise value is hiehlv dependent on the assumed limiting gas accretion rate and the shock physics.,"derived from the accretion shock, peaks in the range of 0.1 to $0.01\,\rm L_\odot$ (solar luminosity), although the precise value is highly dependent on the assumed limiting gas accretion rate and the shock physics."
 A planet caught during this time period would be brighter than at anv other time during its evolution., A planet caught during this time period would be brighter than at any other time during its evolution.
 As gas aceretion is turned off in the core accretion ease. the luminosity rapidly collapses to between 10? to 10°L.. depending upon the mass.," As gas accretion is turned off in the core accretion case, the luminosity rapidly collapses to between $10^{-5}$ to $10^{-6}\,\rm L_\odot$, depending upon the mass."
 At this point. the more massive planets have lower entropies (Figure 2). because a proporüionatelv greater amount of (heir mass has passed through the shock and arrives with low entropy.," At this point, the more massive planets have lower entropies (Figure 2), because a proportionately greater amount of their mass has passed through the shock and arrives with low entropy."
 As a result. post-accretion Iuninosity decreases wilh increasing mass (Figure 3). a result that is entirely a consequence of our treatment of the accretion shock.," As a result, post-accretion luminosity decreases with increasing mass (Figure 3), a result that is entirely a consequence of our treatment of the accretion shock."
 With the lowest gravity. the 1Mj planet has the largest radius (Figure 3) and the highest post-lormation Iuninositv.," With the lowest gravity, the $1\,\rm M_J$ planet has the largest radius (Figure 3) and the highest post-formation luminosity."
 The ‘hot-start’ models begin with arbitrarily large initial luminosities. greater than 10!L.. that expeditiously decay away.," The `hot-start' models begin with arbitrarily large initial luminosities, greater than $10^{-4}\,\rm L_\odot$, that expeditiously decay away."
 Since these planets start fully formed. the choice of lime /=0 for comparison to the core accretion models is somewhat arbitrary.," Since these planets start fully formed, the choice of time $t=0$ for comparison to the core accretion models is somewhat arbitrary."
 In Figure 3. we equate the time of the first hot-start model to time /=0 for the core accretion model.," In Figure 3, we equate the time of the first hot-start model to time $t=0$ for the core accretion model."
" This allows the hot-start models a 2 to 3 million vear ""head start in their cooling aud consequently minimizes (he difference from the core accretion predicted. huminositv.", This allows the hot-start models a 2 to 3 million year `head start' in their cooling and consequently minimizes the difference from the core accretion predicted luminosity.
 Nevertheless. with the sole exception of the 1M; planet. all of the model core-accretion model planets arefainter immediately. after (he end of accretion than the comparable hot start model al the same age.," Nevertheless, with the sole exception of the $1\,\rm\, M_J$ planet, all of the model core-accretion model planets are immediately after the end of accretion than the comparable hot start model at the same age."
 A I0Mj model is over (wo orders of magnitude fainter than if il experienced a hot start.," A $10\rm\, M_J$ model is over two orders of magnitude fainter than if it experienced a hot start."
 The difference in £ falls with mass. reaching a [actor of two [for a 2Mj model.," The difference in $L$ falls with mass, reaching a factor of two for a $2\,\rm M_J$ model."
" The 1M, planet formed by core accretion is a factor of two brighter than produced by (he equivalent hot start."," The $1\,\rm M_J$ planet formed by core accretion is a factor of two brighter than produced by the equivalent hot start."
 In Figure 4 we set time /=0 for the hot-start evolution to coincide with the first core accretion model., In Figure 4 we set time $t=0$ for the hot-start evolution to coincide with the first post-formation core accretion model.
 In (his case. which maxinizes (he difference between the two approaches. (he hot-start Iuminosity is larger for every. planet mass. although the difference is again least for the lowest mass case.," In this case, which maximizes the difference between the two approaches, the hot-start luminosity is larger for every planet mass, although the difference is again least for the lowest mass case."
 As illustrated in this figure. the lower initial entropy of the core accreted planets manifests as both a smaller initial radius and a much smaller effective temperature. both of which lead {ο a smaller huninositv.," As illustrated in this figure, the lower initial entropy of the core accreted planets manifests as both a smaller initial radius and a much smaller effective temperature, both of which lead to a smaller luminosity."
 The hot start evolution predicts that the most massive models at 1 Myr have a radius over twice that of Jupiter's and an effective temperature exceeding 2000 Ix. By contrast. the core accretion calculation predicts Z2«1.5 and Tog<9001x for all cases.," The hot start evolution predicts that the most massive models at 1 Myr have a radius over twice that of Jupiter's and an effective temperature exceeding 2000 K. By contrast, the core accretion calculation predicts $R<1.5\,\rm R_J$ and $T_{\rm eff} < 900\,\rm K$ for all cases."
 Note (hat as the post-core accretion Iuminosity falls very slowly. (he curves almost seem flat on the log-log plot.," Note that as the post-core accretion luminosity falls very slowly, the curves almost seem flat on the log-log plot."
 This is because the small. cool. core accretion planets cool [ar more slowly than (he large. bright. hot-start. planets (see Eq.," This is because the small, cool, core accretion planets cool far more slowly than the large, bright, hot-start planets (see Eq."
 2)., 2).
"and assumes a flat A cosmology with O,,=0.3 and O4=0.7 (Perlmutterefal1997).",and assumes a flat $\Lambda$ cosmology with ${\Omega_m}=0.3$ and ${\Omega_{\Lambda}}=0.7$ \citep{per97}.
. Though the luminosity distance can be calculated via numerical integration. it is preferred to find an analyvlic expression. especially for use in (he Monte Carlo simulations where such a calculation will have to be perlormed many millions of times.," Though the luminosity distance can be calculated via numerical integration, it is preferred to find an analytic expression, especially for use in the Monte Carlo simulations where such a calculation will have to be performed many millions of times."
 Pen(1999) offers such an opüon: however this formula can differ [rom (he numerically determined value by over55., \citet{pen99} offers such an option; however this formula can differ from the numerically determined value by over.
.. Therefore. a polvnomial fit to the numerical solutions is determined [or the redshift range z=0-5 using the Wolfram Research program (o attain a more accurate approximation.," Therefore, a polynomial fit to the numerical solutions is determined for the redshift range z=0-5 using the Wolfram Research program to attain a more accurate approximation."
 When Iuminosity distances determined via this polynomial fit were compared wilh the results of a direct numerical integration. the agreement was well within 1% [or the range 0.1«z4.0: however. at redshifts outside of this range. the disagreement is as large as7%.," When luminosity distances determined via this polynomial fit were compared with the results of a direct numerical integration, the agreement was well within 1 for the range $ 0.1 < z <4.0 $; however, at redshifts outside of this range, the disagreement is as large as."
".. Column(9) shows the broadband spectral index. 0,; gives the index a,,."," Column(9) shows the broadband spectral index, $\alpha_{ro}$ gives the index $\alpha_{og}$ ."
 Cohunn (11) has the gamma-ray variability parameter. ὅμως. as determined by Nolanefa£.(2003).," Column (11) has the gamma-ray variability parameter, $\delta_{var}$, as determined by \citet{nol03}."
. Column (12) gives the references for the data. where the first relerence is for redshift. the second for the radio data. the third for (he optical magnitudes. ancl the fourth for (he gamma rav. data.," Column (12) gives the references for the data, where the first reference is for redshift, the second for the radio data, the third for the optical magnitudes, and the fourth for the gamma ray data."
 We have evaluated the strength and significance of correlation between monochromatic luminosities using non-paranmetric methods. such as IXendall's 7 (Ixendall&Gibbous1990).," We have evaluated the strength and significance of correlation between monochromatic luminosities using non-parametric methods, such as Kendall's $\tau$ \citep{ken90}."
. These results are given in Table 2., These results are given in Table 2.
 Colunn(1) gives the independent variable. Column(2) the dependent. variable. Column(3) the munuber in the sample. Column(4) Kendall's 7 statistic. Column(5) the probability of the null result for this statistic. Colunn(6) the Spearman's p statistic. Column(7) the probability of the null result for Chis statistic. aud in Column(8) the regression technique used.," Column(1) gives the independent variable, Column(2) the dependent variable, Column(3) the number in the sample, Column(4) Kendall's $\tau$ statistic, Column(5) the probability of the null result for this statistic, Column(6) the Spearman's $\rho$ statistic, Column(7) the probability of the null result for this statistic, and in Column(8) the regression technique used."
 Here BJ refers to the Ducklev-Janmes method. EM refers to the EM method and 5B relers to Schinitt’s binned regression. as explained in (he references (hat follow.," Here BJ refers to the Buckley-James method, EM refers to the EM method and SB refers to Schmitt's binned regression, as explained in the references that follow."
 Cohunn(9) gives the linear regression slope. and Column(10) the linear regression Y-intercept.," Column(9) gives the linear regression slope, and Column(10) the linear regression Y-intercept."
 In some cases. we have upper and lower limits to measured values. aud (ο take (his into account. we use the proper correlation ancl regression techniques as discussed in Feigelson&Nelson(1935):Isobe.(1986) and utilized in the ASURV software package.," In some cases, we have upper and lower limits to measured values, and to take this into account, we use the proper correlation and regression techniques as discussed in \citet{fei85,iso86} and utilized in the ASURV software package."
 In reporting correlation coefficients. probabiliGes and regression fits. we adopt the format of Muckeefa£.(1997).," In reporting correlation coefficients, probabilities and regression fits we adopt the format of \citet{muc97}."
. That is. we report three significant digits down to 0.100.," That is, we report three significant digits down to 0.100."
 Between 0.010. and 0.100 we report (wo signilicant dieits. and for all numbers with values below 0.01. we only report with 1 significant digit.," Between 0.010 and 0.100 we report two significant digits, and for all numbers with values below 0.01, we only report with 1 significant digit."
 Most notably. we determine that there is a strong correlation between gamma ray and radio luminosity (Figure 1 and Table2).," Most notably, we determine that there is a strong correlation between gamma ray and radio luminosity (Figure 1 and Table2)."
 In order to determine whether the correlation is, In order to determine whether the correlation is
Figue 16. shows the eCDFs for both the detected IIVSs aud the simulated IIVSs (obtained from different evolutionary models for the DDITI mechanisin listed iu Table 3)).,Figure \ref{fig:f16} shows the $v$ CDFs for both the detected HVSs and the simulated HVSs (obtained from different evolutionary models for the BBH mechanism listed in Table \ref{tab:tab3}) ).
 Our caleulatious show that the eCDE. (with cyomTOOkins 1j ds almost independent of the choice of the thickness of the disk where the IVS progenitors are originated., Our calculations show that the $v$ CDF (with $v^{\infty}_{\rm ej}>700 \kms$ ) is almost independent of the choice of the thickness of the disk where the HVS progenitors are originated.
 The simulated οςΤΕ is insensitive to the choice of the orientation of the disk relative to the BBII orbital plane., The simulated $v$ CDF is insensitive to the choice of the orientation of the disk relative to the BBH orbital plane.
 The simulated eCDF is also inscusitive to the choice of the mass ratio vy Gf p~0.01. 0.001) aud the eccentricity of the BBIT (alsoseeSesanaetal. 2007).," The simulated $v$ CDF is also insensitive to the choice of the mass ratio $\nu$ (if $\nu\sim0.01$$-$$0.001$ ) and the eccentricity of the BBH \citep[also
see][]{Sesana07}."
.. As seen from Figure 16. differeut models produce quite similar eCDFs aud too many IIVSs with ejection volocities substantially higher than those of the detected IIVSs.," As seen from Figure \ref{fig:f16}, different models produce quite similar $v$ CDFs and too many HVSs with ejection velocities substantially higher than those of the detected HVSs."
 The primary reason for these similarly flat eCDFs is as follows., The primary reason for these similarly flat $v$ CDFs is as follows.
 The IIVS proecuitors injected iuto the region with peuetratiou parameter n/appi< lean gain sole enerev duriug the dvnamical interaction with the DDBIL and the mean cnerev eain is mainly determined by the semüniajor axis aud the mass ratio of the BBIT (Quinlan1996:Sesanaetal.2006) but insensitive to the detailed values of αν.," The HVS progenitors injected into the region with penetration parameter $r_{\rm p}/a_{\rm BBH}\la 1$ can gain some energy during the dynamical interaction with the BBH, and the mean energy gain is mainly determined by the semimajor axis and the mass ratio of the BBH \citep{Quinlan96,Sesana06} but insensitive to the detailed values of $r_{\rm p}/a_{\rm BBH}$."
 With increasing αμ to be slightly larger than 1. the οποιον gain dramatically drops to 0 (alsoseeSesanactal.2006).," With increasing $r_{\rm p}/a_{\rm BBH}$ to be slightly larger than 1, the energy gain dramatically drops to 0 \citep[also see][]{Sesana06}."
. Therefore. the slope of the eCDF is larecly determined bv the scatters of the energv gaius around their rius value at Πα<Le aud the scatters are mainly due to the different orbital phases of the BBIT aud inscusitive to other BBII parameters.," Therefore, the slope of the $v$ CDF is largely determined by the scatters of the energy gains around their rms value at $r_{\rm p}/a_{\rm BBH}\la 1$, and the scatters are mainly due to the different orbital phases of the BBH and insensitive to other BBH parameters."
" Because of these features in the DDBII mechauisui even the ""BW model for injecting stars, which suppresses the ejection of IIVSs at the hüehl-velocity cud for the Το] mechauisii. does not lead to such a suppression for the BBIT mechanism here."," Because of these features in the BBH mechanism, even the “RW” model for injecting stars, which suppresses the ejection of HVSs at the high-velocity end for the TBK mechanism, does not lead to such a suppression for the BBH mechanism here."
 With the relevant parameters for the models listed πι Table 3. the simulated eCDFs are all quite fat (except for the “FIXN-UB™ model). compared with the observational oues.," With the relevant parameters for the models listed in Table \ref{tab:tab3}, , the simulated $v$ CDFs are all quite flat (except for the “FIX-UB” model), compared with the observational ones."
" For the first IWS population. our Is- tests find that the lisclihoods for the simulated ¢CDFs to be consistent with the observed distribution are 0.01. 0.158. 0.003. 0,0001. 0.0001. and 0.01 for the Εν ""EIX- ""CAVIT. πονοον ποον and ""ExpD models. respectively,"," For the first HVS population, our K-S tests find that the likelihoods for the simulated $v$ CDFs to be consistent with the observed distribution are $0.01$, $0.18$, $0.003$, $0.0001$, $0.0001$, and $0.01$ for the “FIX”, ``FIX-UB'', “GW1”, “GW2”, “GW3”, and “ExpD” models, respectively."
" Except the high likelihood for tle ""FIX-UB"" model. the other small νο likelihoods above sugecst that the first population of the detected IIVSs are uulikelv drawn from the same distributions as those obtained from the other models"," Except the high likelihood for the ``FIX-UB'' model, the other small K-S likelihoods above suggest that the first population of the detected HVSs are unlikely drawn from the same distributions as those obtained from the other models."
" For the secoud TVS population. the Ίο likelihoods are 0.03. 0.12. 0.05. 0.02. 0.02. and 0.01 for the ""FIN,. ""FIN-UD. παΝΕ, παον 0037. aud ""προ auodels respectively."," For the second HVS population, the K-S likelihoods are $0.03$, $0.12$, $0.05$, $0.02$, $0.02$, and $0.04$ for the “FIX”, “FIX-UB”, “GW1”, “GW2”, “GW3”, and “ExpD” models, respectively."
" Those likelihoods are not simall cnough so thato the eCDE for the second. IIVS populatiou does not appear to be inconsistent with those obtained from the those models. especially for the ""FIN-UD"" τος."," Those likelihoods are not small enough so that the $v$ CDF for the second HVS population does not appear to be inconsistent with those obtained from the those models, especially for the “FIX-UB” model."
 For the combined salple of both the first aud the second ITVS populations. the KS likelihoods are 0.03. 0.11. 0.003. 3.1073410. 7. and 0.02. respectively.," For the combined sample of both the first and the second HVS populations, the KS likelihoods are $0.03$, $0.14$, $0.003$, $3\times 10^{-5}$, $3\times 10^{-5}$, and $0.02$, respectively."
" For the models in which epp evolves with time (1.6.. the ""CW aud ""ExpD imodels). our simulations iudicate a weak dependence of the ejection velocity ou time. as the DDII sominajor axis is smaller at a later time and the corresponding ejection velocity is larger (see Equation (9))."," For the models in which $a_{\rm BBH}$ evolves with time (i.e., the “GW” and “ExpD” models), our simulations indicate a weak dependence of the ejection velocity on time, as the BBH semimajor axis is smaller at a later time and the corresponding ejection velocity is larger (see Equation (9))."
 Compared to our results. the simulated eCDE obtained by Sesanactal.(2007) is flatter. as in heir model the BBIT decays fast at large separations aud nore stars with hieher velocities are ejected at a later iue as the BBIT becomes harder.," Compared to our results, the simulated $v$ CDF obtained by \citet{Sesana07}
 is flatter, as in their model the BBH decays fast at large separations and more stars with higher velocities are ejected at a later time as the BBH becomes harder."
 Due to the Linited statistics of the detected IIVSs aud the uncertainties im he BBIL dynamical evolution model. it appears to be xemature to use the eCDF to distinguish whether the IIVSs ave ejected due to the TDBIS mechanisin or the DDBIT nechanisin aud also coustrain detailed dynamical models or Injecting stars to the MDII(s) vicinity.," Due to the limited statistics of the detected HVSs and the uncertainties in the BBH dynamical evolution model, it appears to be premature to use the $v$ CDF to distinguish whether the HVSs are ejected due to the TBK mechanism or the BBH mechanism and also constrain detailed dynamical models for injecting stars to the MBH(s) vicinity."
 Note that the ejected stars are actually moving in the Galactic potential. some of those with low velocities (0.9. ey<οΊαν 1j nay return to the vicinity of the BBU and receive further kicks. which iav chanee the final eCDE.," Note that the ejected stars are actually moving in the Galactic potential, some of those with low velocities (e.g., $v^{\infty}_{\rm ej}\la400 \kms$ ) may return to the vicinity of the BBH and receive further kicks, which may change the final $v$ CDF."
 We check this effect by doing additional threc-body experineuts for the ejected stars with low ejection velocities., We check this effect by doing additional three-body experiments for the ejected stars with low ejection velocities.
 We find that the obtained final eCDFs are ouly slightly fatter than those shown in Figure 16 and our conclusions above will not be siguificautlv affected., We find that the obtained final $v$ CDFs are only slightly flatter than those shown in Figure \ref{fig:f16} and our conclusions above will not be significantly affected.
Tn this paper. we first use three-body experiments to study the interactions of the MDII with binary stars bound to the MDII.,"In this paper, we first use three-body experiments to study the interactions of the MBH with binary stars bound to the MBH."
 We fud that the probability of ejecting IIVSs is substantially culanced due to multiple encounters between the MDBIT and the stellar binariesInjecting inte its vicinity even at a distance substantially, We find that the probability of ejecting HVSs is substantially enhanced due to multiple encounters between the MBH and the stellar binariesinjecting into its vicinity even at a distance substantially
 Over the past decade. hundreds of free-floating brown dwarfs have been discovered iu the field aud in voune clusters (Dagi2000).. but their origin remains a mystery.," Over the past decade, hundreds of free-floating brown dwarfs have been discovered in the field and in young clusters \citep{bas00}, but their origin remains a mystery."
 In recent numerical lvdrodvuaimical sinulatious of the fragmentation of molecular cloud cores. selt-exavitatiug objects can form with initial masses of oulv ~1 Maj. but these fragments coutiuue to accrete matter from their surrounding core. uxuallv. to the poiut of eventually reaching stellar masses (Boss2001:Dateetal.2003).," In recent numerical hydrodynamical simulations of the fragmentation of molecular cloud cores, self-gravitating objects can form with initial masses of only $\sim1$ $M_{\rm Jup}$, but these fragments continue to accrete matter from their surrounding core, usually to the point of eventually reaching stellar masses \citep{bos01,bat03}."
". As a result. the formation of brown dafs appears difficult to achieve in these models,"," As a result, the formation of brown dwarfs appears difficult to achieve in these models."
 Oue possible solution has heen proposed by BReipurth&Clarke(2001). aud Boss (2001).. and investigated further by Dateetal.(2002) aud Bouvier (2003)..," One possible solution has been proposed by \citet{rc01} and \citet{bos01}, , and investigated further by \citet{bat02} and \citet{kb03b}."
 Ta this scenario. dvuamical interactions within a small cluster of protostellar sources leads to the ejection of one of the embryos. which prematurely halts accretion from the core’s reservoir of infalline laterial aud results iu the formation of à brown dwiut.," In this scenario, dynamical interactions within a small cluster of protostellar sources leads to the ejection of one of the embryos, which prematurely halts accretion from the core's reservoir of infalling material and results in the formation of a brown dwarf."
 Models for the origin of brown dwarfs can be tested through measurements of the uniltiplicity of substellar objects., Models for the origin of brown dwarfs can be tested through measurements of the multiplicity of substellar objects.
 Significaut progress in this area has becu made through extcusive direct maging of primarics in the field with masses near aud below the bydrogen burning mass Bit (xocrucrCazisetal. 2003)..," Significant progress in this area has been made through extensive direct imaging of primaries in the field with masses near and below the hydrogen burning mass limit \citep{koe99,mar99,rei01b,clo02a,clo02b,clo03,bou03,bur03,fre03,giz03}."
 In the data from these surveys. binary brown cawarfs occur at a rate of !~Ly for separations of &>1 AU and exhibit κατα separations of «~WwAU.," In the data from these surveys, binary brown dwarfs occur at a rate of $\sim15$ for separations of $a>1$ AU and exhibit maximum separations of $a\sim20$ AU."
 Ejection models predict a somewhat lower binary fraction (~5% 1) but a similar maxima separation (a~10 AU) 2002).. Bureas, Ejection models predict a somewhat lower binary fraction $\sim5$ ) but a similar maximum separation $a\sim10$ AU) \citep{bat02}.
seretal.(2003) concluded that the observed maxiuiuun separation is not a reflection of disruption of wider binaries by mteractions with stars or molecular clouds. which is supported by the abseuce of wide binaries among substellar primaries in less evolved populations iu open clusters (Morti and in star-forming regious (Neuhauseretal.2002," \citet{bur03} concluded that the observed maximum separation is not a reflection of disruption of wider binaries by interactions with stars or molecular clouds, which is supported by the absence of wide binaries among substellar primaries in less evolved populations in open clusters \citep{mar98,mar00,mar03} and in star-forming regions \citep{neu02,bou04}."
:Bouvct200L).. Tustead. suggested that wide low-mass binaries do not form or are disrupted at a very carly stage.," Instead, \citet{bur03} suggested that wide low-mass binaries do not form or are disrupted at a very early stage."
 Ireport the discovery of the first known widely-separated binary brown chwart. which was found sereudipitously during observations of candidate vouus brown dwarfs in the Chamacleou I star-forming region.," I report the discovery of the first known widely-separated binary brown dwarf, which was found serendipitously during observations of candidate young brown dwarfs in the Chamaeleon I star-forming region."
 I prescut optical aud near-IR images aud optical spectroscopy of the compoucuts of the pair 2)). assess the membership of the two objects in Chamacleon T 3.1)) and iu a binary system 2.2)). place them on the IIertzspruug-BRussell CI-R) diagram 3.3)). estimate their masses from theoretical evolutionary models 3.1)). aud discuss the duplications of this new binarysvstem for theories of brown cwarf formation £)).," I present optical and near-IR images and optical spectroscopy of the components of the pair \ref{sec:obs}) ), assess the membership of the two objects in Chamaeleon I \ref{sec:mem}) ) and in a binary system \ref{sec:bin}) ), place them on the Hertzsprung-Russell (H-R) diagram \ref{sec:ext}) ), estimate their masses from theoretical evolutionary models \ref{sec:mass}) ), and discuss the implications of this new binarysystem for theories of brown dwarf formation \ref{sec:dis}) )."
On the other hand. there is an alternative way in cases of mild pile-up: using only PATTERNZO events.,"On the other hand, there is an alternative way in cases of mild pile-up: using only $PATTERN$ =0 events."
 After such a filtering. we compared the PATTERN=O0 spectra extracted in à circular region with the spectra extracted in annular regions (hence free of pile-up) using the usual PATTERN=0-12 for MOS and PATTERN=0-4 for pn.," After such a filtering, we compared the $PATTERN$ =0 spectra extracted in a circular region with the spectra extracted in annular regions (hence free of pile-up) using the usual $PATTERN$ =0–12 for MOS and $PATTERN$ =0–4 for pn."
 This check was done by fitting simple 2-temperatures models on LWmedium data from Rev. 0156. LWthick data from Rev. 0731 and SWthick data from Rev. 1814 for EPIC-pn. and on LW«thick data from Rev. 0156 and SW«thick data from Rev. 1814 for EPIC-MOS.," This check was done by fitting simple 2-temperatures models on LW+medium data from Rev. 0156, LW+thick data from Rev. 0731 and SW+thick data from Rev. 1814 for EPIC-pn, and on LW+thick data from Rev. 0156 and SW+thick data from Rev. 1814 for EPIC-MOS."
 The comparison is excellent for EPIC-MOS data: fluxes and count rates differ by <1% and best-fit spectral parameters are within the errors., The comparison is excellent for EPIC-MOS data: fluxes and count rates differ by $<$ and best-fit spectral parameters are within the errors.
 The remaining difference can be attributed to the slightly larger noise in the spectra extracted in annuli and from calibration differences., The remaining difference can be attributed to the slightly larger noise in the spectra extracted in annuli and from calibration differences.
 The comparison ts less perfect for EPIC-pn. especially for the data taken with the medium filter: the flux differences reaches in this case. and best-fit spectral parameters are at 2-σ from each other.," The comparison is less perfect for EPIC-pn, especially for the data taken with the medium filter: the flux differences reaches in this case, and best-fit spectral parameters are at $\sigma$ from each other."
" The pile-up thus still has a small influence on the EPIC-pn data taken in the large window mode. and those data should thus be considered with caution,"," The pile-up thus still has a small influence on the EPIC-pn data taken in the large window mode, and those data should thus be considered with caution."
 A final check was made on the data from Rev. 1620. which ylelds the most discrepant (2.57) position forPuppis.. if we trust the pipeline processing.," A final check was made on the data from Rev. 1620, which yields the most discrepant (2.5”) position for, if we trust the pipeline processing."
 We first derived the position of ffrom the PATTERN=0 data using the SAS task edetectcliain. and then extracted the spectra using a circular region centered on that position.," We first derived the position of from the $PATTERN$ =0 data using the SAS task $edetectchain$, and then extracted the spectra using a circular region centered on that position."
 We compared these spectra to those extracted on PATTERN=0 data using a circular region centered on the Hipparcos position ofPuppis., We compared these spectra to those extracted on $PATTERN$ =0 data using a circular region centered on the Hipparcos position of.
. Both sets of spectra appear identical in Xspec: the small centroiding errors have thus no impact on the spectra as long as a circular region is used., Both sets of spectra appear identical in Xspec: the small centroiding errors have thus no impact on the spectra as long as a circular region is used.
 We therefore extracted lightcurves and spectra of iun a circular region centered on the Hipparcos position of the target., We therefore extracted lightcurves and spectra of in a circular region centered on the Hipparcos position of the target.
 We used only the PATT ERNZO event files., We used only the $PATTERN$ =0 event files.
 While this is not necessary for the small window mode. it ensures a homogeneous data reduction.," While this is not necessary for the small window mode, it ensures a homogeneous data reduction."
 Xspec v12.6.0 was used to fit spectra. and our own software to analyze EPIC lighteurves.," Xspec v12.6.0 was used to fit spectra, and our own software to analyze EPIC lightcurves."
 Note that spectra were grouped using the new SAS task specgroup., Note that spectra were grouped using the new SAS task $specgroup$.
" It enables to reduce the oversampling. which may “cause problems during spectral fitting because the spectral bins are then not completely independent"" (excerpt from SAS 10.0.0. online documentation)."," It enables to reduce the oversampling, which may “cause problems during spectral fitting because the spectral bins are then not completely independent” (excerpt from SAS 10.0.0, online documentation)."
 We choose an oversampling factor of 5. ensuring that no spectral bin is narrower than 1/5 of the full width half maximum resolution at the central photon energy of the bin.," We choose an oversampling factor of 5, ensuring that no spectral bin is narrower than 1/5 of the full width half maximum resolution at the central photon energy of the bin."
 Note that. while providing more statistically correct data. this process dramatically reduces the number of spectral bins.," Note that, while providing more statistically correct data, this process dramatically reduces the number of spectral bins."
 The data were also grouped to ensure that a minimum signal-to-noise of 3 was reached in each spectral bin of the background-corrected spectra., The data were also grouped to ensure that a minimum signal-to-noise of 3 was reached in each spectral bin of the background-corrected spectra.
 The RGS datasets were also reduced in à standard way with SAS v10.0.0., The RGS datasets were also reduced in a standard way with SAS v10.0.0.
 Many new. important RGS features were modified in that version (e.g. the spectral binning in wavelength rather than in dispersion angle units).," Many new, important RGS features were modified in that version (e.g. the spectral binning in wavelength rather than in dispersion angle units)."
 This ensures a better calibration of our datasets., This ensures a better calibration of our datasets.
 It also solved the calibration problems (wavelength shift and reduced flux) found when using earlier versions of the SAS for the two observations where wwas placed off-axis., It also solved the calibration problems (wavelength shift and reduced flux) found when using earlier versions of the SAS for the two observations where was placed off-axis.
 Note that the data were extracted using the proposal position of the source. which ts the same in all observations but the first two (Revs.," Note that the data were extracted using the proposal position of the source, which is the same in all observations but the first two (Revs."
 0091 and 0156. shift of 0.0002* in both RA and DEC) - this small position shift has no impact on the derived RGS spectra.," 0091 and 0156, shift of $^{\circ}$ in both RA and DEC) - this small position shift has no impact on the derived RGS spectra."
 When detected. flares were discarded using resfilter’.," When detected, flares were discarded using $rgsfilter$."
 The tasks and then provided unbinned source and background spectra. as well as response matrices for each order (1.2) and each instrument (1.2).," The tasks and then provided unbinned source and background spectra, as well as response matrices for each order (1,2) and each instrument (1,2)."
 A final. combined spectrum was also calculated using all 18 RGS datasets and the task rgscombine.," A final, combined spectrum was also calculated using all 18 RGS datasets and the task $rgscombine$."
 The background files and matrix responses were attached to the source spectra using the new SAS task specgroup. which we also use to ensure an oversampling factor of maximum 5 (see above).," The background files and matrix responses were attached to the source spectra using the new SAS task $specgroup$, which we also use to ensure an oversampling factor of maximum 5 (see above)."
 Fluxed spectra combining both RGS instruments and both orders were obtained using the task rgsf/uxer., Fluxed spectra combining both RGS instruments and both orders were obtained using the task $rgsfluxer$.
 Note that a correction for off-axis angles ts applied to ensure that the fluxes are real photon fluxes and not simply recorded count rates (1:89. the arf response matrix is fully taken into account)., Note that a correction for off-axis angles is applied to ensure that the fluxes are real photon fluxes and not simply recorded count rates (i.e. the arf response matrix is fully taken into account).
 The spectra of one revolution were sampled to get 1500 spectral bins. while the spectrum combining the 18 datasets was calculated to get 3000 spectral bins.," The spectra of one revolution were sampled to get 1500 spectral bins, while the spectrum combining the 18 datasets was calculated to get 3000 spectral bins."
 This ensures an oversampling factor of about 3 and 6 for the former and latter cases. respectively.," This ensures an oversampling factor of about 3 and 6 for the former and latter cases, respectively."
 Two caveats should be noted., Two caveats should be noted.
 First. the rmf matrix is not fully taken into account by rgsfIuxer and there is no instrumental width correction.," First, the rmf matrix is not fully taken into account by $rgsfluxer$ and there is no instrumental width correction."
 This needs to be accounted for when modelling the spectra (see Paper Hl)., This needs to be accounted for when modelling the spectra (see Paper III).
 Second. there are known small wavelength shifts in RGS spectra. apparently depending on the Sun aspect angle.," Second, there are known small wavelength shifts in RGS spectra, apparently depending on the Sun aspect angle."
 However. no sign of such an effect is detected in our dataset when we use SAS v10 (they would appear as small spectral variations. and there are none. see Paper II).," However, no sign of such an effect is detected in our dataset when we use SAS v10 (they would appear as small spectral variations, and there are none, see Paper II)."
 With data of such high quality. the error bars on the spectra are very small. and it is therefore very difficult to get a formally acceptable fit.," With data of such high quality, the error bars on the spectra are very small, and it is therefore very difficult to get a formally acceptable fit."
 We thus avoided to try to get a perfect fit (Le. vo 1). which is actually impossible to get without going into unrealistic. overcomplicated models (e.g. 10 components fits with independent. free abundances) when all instruments agree.," We thus avoided to try to get a perfect fit (i.e. $\chi^2\sim1$ ), which is actually impossible to get without going into unrealistic, overcomplicated models (e.g. 10 components fits with independent, free abundances) when all instruments agree."
 Rather. we have tried to get a fit as simple as possible which ts at the same time as realistic as possible and as close as possible to the spectral data.," Rather, we have tried to get a fit as simple as possible which is at the same time as realistic as possible and as close as possible to the spectral data."
bv the All-Skv. Monitor (ASAI) onboard the Rossi X-ray Timing Explorer (RATE).,by the All-Sky Monitor (ASM) onboard the Rossi X-ray Timing Explorer (RXTE).
 After ihe ASAI collected sufficient amounts of data. several time-frequency analysis methods such as the wavelet. transform (Riboéetal.2001).. dynamic power spectrum (Clarksonetal. 2003a).. and sliding Lomb-Scargle periodogram (Trowbridgeetal.2007) were applied to ihe lisht curve to investigate the variations in the superorbital period of SAIC X-1.," After the ASM collected sufficient amounts of data, several time-frequency analysis methods such as the wavelet transform \citep{Ribo2001}, dynamic power spectrum \citep{Clarkson2003}, and sliding Lomb-Scargle periodogram \citep{Trowbridge2007} were applied to the light curve to investigate the variations in the superorbital period of SMC X-1."
 The mechanism of superorbital modulations in SAIC Ντ. similar to those in Her N-1. LMC X-4. and (νο X-2. ave interpreted by a warped and tilted accretion disk 1993).," The mechanism of superorbital modulations in SMC X-1, similar to those in Her X-1, LMC X-4, and Cyg X-2, are interpreted by a warped and tilted accretion disk \citep{Wojdowski1998}."
. When the disk precesses. it obscures our line of view to the central X-ray source.," When the disk precesses, it obscures our line of view to the central X-ray source."
 llowever. (he mechanisms that cause (lie precession period to change with Gime are still unknown.," However, the mechanisms that cause the precession period to change with time are still unknown."
 Ribóοἱal.(2001) analvzed the first ~1700 d of the ASM light curve and found that the superorbital period first decreases [rom 60 d (o ~45 d and subsequently increases to ~GO d. By fitting the maximum points of the wavelet spectrum will a sinusoidal curve. thev found that the modulation period varies wilh a period of 142148 d. analvzed the first ~2200 d of the ASAI light eurve along with all the Burst aud Transient Source Experiment (DATSE) data using the dvnamie power spectral technique.," \citet{Ribo2001} analyzed the first $\sim 1700$ d of the ASM light curve and found that the superorbital period first decreases from $\sim60$ d to $\sim45$ d and subsequently increases to $\sim60$ d. By fitting the maximum points of the wavelet spectrum with a sinusoidal curve, they found that the modulation period varies with a period of $1421\pm8$ d. \citet{Clarkson2003} analyzed the first $\sim2200$ d of the ASM light curve along with all the Burst and Transient Source Experiment (BATSE) data using the dynamic power spectral technique."
 The dynamic power spectra of both the RNTE and BATSE light curves show that the superorbital modulation period changes between GO d and ~45 d. The superorbital period appears (ο vary in a 1600 d time scale., The dynamic power spectra of both the RXTE and BATSE light curves show that the superorbital modulation period changes between $\sim 60$ d and $\sim 45$ d. The superorbital period appears to vary in a $\sim 1600$ d time scale.
 Irowbridgeοἱal.(2007) analyzed the first 4000 d of ASAI data using the sliding Lomb-Searele periodogram and the the evele length counting method. and the results indicated that the evele length variation is perhaps not as smooth as Chat obtained from the ονπας power spectrum.," \citet{Trowbridge2007} analyzed the first $\sim$ 4000 d of ASM data using the sliding Lomb-Scargle periodogram and the the cycle length counting method, and the results indicated that the cycle length variation is perhaps not as smooth as that obtained from the dynamic power spectrum."
 The evele length variation possibly contains a shorter periodicity in addition to the ~1600 d period., The cycle length variation possibly contains a shorter periodicity in addition to the $\sim1600$ d period.
 llowever. due to the constraints of the window size in both the dvnamie power spectrum and the sliding Lomb-Scarele periodogram. it is dillicull to further analvze the periodicities of the superorbital period variation.," However, due to the constraints of the window size in both the dynamic power spectrum and the sliding Lomb-Scargle periodogram, it is difficult to further analyze the periodicities of the superorbital period variation."
simular BCD sample aud remeasured the 1.lCdIz fluxes using the images from NWSS aud FIRST. providing radio detections for 2 sources and better upper ΜΗ. for another five sources that overlaps with the galaxies in our sample.,"similar BCD sample and remeasured the 1.4GHz fluxes using the images from NVSS and FIRST, providing radio detections for 2 sources and better upper limits for another five sources that overlaps with the galaxies in our sample."
 Our final sample consists of 28 ealaxies. all of which haveSpitzer mid-IR μια and/or you flux neasurements.," Our final sample consists of 28 galaxies, all of which have mid-IR $\mu$ m and/or $\mu$ m flux measurements."
 Among these galaxies. 23 sources have ΓΕ] radio continue data and 5 have measured upper limits.," Among these galaxies, 23 sources have 1.4GHz radio continuum data and 5 have measured upper limits."
" pau aud pan fluxes are available or 16 sources and 3 more are detected only at μα, We list the photometry of our sources in Table 1..", $\mu$ m and $\mu$ m fluxes are available for 16 sources and 3 more are detected only at $\mu$ m. We list the photometry of our sources in Table \ref{tab1}.
" The uncertainty in the 22,4n or {μια photometry is less han", The uncertainty in the $\mu$ m or $\mu$ m photometry is less than.
 TheZRAS μι aud san fluxes typically ave less than a few percent of error. but could eo up to 415i for some of the fainter sources.," The $\mu$ m and $\mu$ m fluxes typically have less than a few percent of error, but could go up to $\sim$ for some of the fainter sources."
 The rius noise level for NVSS is «0.5 |l, The rms noise level for NVSS is $\sim$ $^{-1}$.
 The sensitivity aud efficiency of theSpitzer Space Telescope has allowed us to probe the correlation between the iid-IB. and radio Iuniünosities for a large number of ealaxies., The sensitivity and efficiency of the Space Telescope has allowed us to probe the correlation between the mid-IR and radio luminosities for a large number of galaxies.
 Using the jnu aud jin MOPS imaging data of the First Look Survey. Appletonetal.(2001) have demonustrated the first direct evidence for the universality of the mid-IR/radio aud FIR/radio correlation to z~l.," Using the $\mu$ m and $\mu$ m MIPS imaging data of the First Look Survey, \citet{Appleton04} have demonstrated the first direct evidence for the universality of the mid-IR/radio and FIR/radio correlation to $\sim$ 1."
 Wuetal.(2005). have also studied the mid-IR/racdio correlation in a smuuple of star-forming galaxies aac found that both the Syan and jan huumnostyv are clearly correlated with the 1.1€1IIz. radio IDunuinosity., \citet{Wu05} have also studied the mid-IR/radio correlation in a sample of star-forming galaxies and found that both the $\mu$ m and $\mu$ m luminosity are clearly correlated with the 1.4GHz radio luminosity.
" ""Their sample iucluded only a few (3) dwarf galaxies aux suggested that there may be a slope change for chwart ealaxies. which could be due to the lower dust-to-gas ratios and lower metallicities of the dwarts."," Their sample included only a few (3) dwarf galaxies and suggested that there may be a slope change for dwarf galaxies, which could be due to the lower dust-to-gas ratios and lower metallicities of the dwarfs."
 A dotaile analysis of the spatial distribution of the infrared to racio correlation usingSpitzer data on a sample of nearby late type spiral galaxies has been performed by Murphyοal.(20062.b). in which the authors found that the ratio of the mid-IR and FIR to radio eiuission docs vary from one region to the other.," A detailed analysis of the spatial distribution of the infrared to radio correlation using data on a sample of nearby late type spiral galaxies has been performed by \citet{Murphy06a,Murphy06b} in which the authors found that the ratio of the mid-IR and FIR to radio emission does vary from one region to the other."
 Iu general though. the dispersion is all aud the ratio decreases with surface brightucss and galactocentric radius.," In general though, the dispersion is small and the ratio decreases with surface brightness and galactocentric radius."
 Using the data from this suple of low ictallicity dwarf galaxies we plot in Fie., Using the data from this sample of low metallicity dwarf galaxies we plot in Fig.
 1 the radio huuinositv of the sample as a function of the FIR. hunuinosity., 1 the radio luminosity of the sample as a function of the FIR luminosity.
 The huuinosities of the sources we study span nearly l orders of magnitudes. but the correlation between the FIR aud the radio is remarkably tight.," The luminosities of the sources we study span nearly 4 orders of magnitudes, but the correlation between the FIR and the radio is remarkably tight."
 The scatter in the ratio of the FIR aud radio bhuuinositfies is 0.23 dex. ie. less than a factor of 2.," The scatter in the ratio of the FIR and radio luminosities is 0.23 dex, i.e. less than a factor of 2."
 We perfouned a least-squares bisector fit to the data and found: teu «log(Lgis/L.110.75.," We performed a least-squares bisector fit to the data and found: $_{\rm 1.4GHz}$ $^{-1}$ $\times$ $_{\rm
 FIR}$ $_\odot$ )+10.75."
 For logicolmparison. we have 1j|21.09imcluded in the plot the starburst ealaxies from Draudletal.(2006) marked with trianeles. which has au identical slope (within 16).," For comparison, we have included in the plot the starburst galaxies from \citet{Brandl06} marked with triangles, which has an identical slope (within $\sigma$ )."
 This slope of L.09+0.07 for the cawart galaxy data agrees well with the slope of 1.102:0.014 found by Bell(2003). for a sample of 162 ealaxies. as well as the slope of 1110.02 for the infrared selected sources from theFRAS Bright Calaxy Sample (BCS) (Coudonetal.1991).," This slope of $\pm$ 0.07 for the dwarf galaxy data agrees well with the slope of $\pm$ 0.04 found by \citet{Bell03} for a sample of 162 galaxies, as well as the slope of $\pm$ 0.02 for the infrared selected sources from the Bright Galaxy Sample (BGS) \citep{Condon91}."
. This is also in agreement with ITopliusetal.(2002) and would suggest that elobally our DCDs have a very similar FIR/racio correlation to normal galaxies.," This is also in agreement with \citet{Hopkins02}
 and would suggest that globally our BCDs have a very similar FIR/radio correlation to normal galaxies."
" Another wav to parameterize the IR/radio correlation is to calculate the ratio of FIR to radio Iuniuositv gery according to the ITelouetal.(1985). formula as well as the qo, following the definition of Appletonetal.(2001) (2001).", Another way to parameterize the IR/radio correlation is to calculate the ratio of FIR to radio luminosity $q_{\rm FIR}$ according to the \citet{Helou85} formula as well as the $q_{24}$ following the definition of \citet{Appleton04} .
".. We plot the do, values for our saunple asa function of the lan lIuminositv of the ealaxies iu Fig.", We plot the $q_{24}$ values for our sample asa function of the $\mu$ m luminosity of the galaxies in Fig.
 2., 2.
 For this sample we fud that gry;= 2.1z:0.2. cousisteut with the value of normal galaxies of qgqji= 2.30.2 found by Concou (1992).," For this sample we find that $q_{\rm FIR}=$ $\pm$ 0.2, consistent with the value of normal galaxies of $q_{\rm FIR}=$ $\pm$ 0.2 found by \citet{Condon92}."
". When using the mid-IR {μια data. we fiud that qa, =L340.1 (sce Table 2))."," When using the mid-IR $\mu$ m data, we find that $q_{24}=$ $\pm$ 0.4 (see Table \ref{tab2}) )."
" The standard deviation in do, is ~twice that of 4gig.", The standard deviation in $q_{24}$ is $\sim$ twice that of $q_{\rm FIR}$.
 This is not unexpected eiven that the spectrum of star forming galaxies shows substantially larger variations in spectral slope in the nud-IR (seeBrandletal.2006) than in the FIR (Daleetal. 2006)., This is not unexpected given that the spectrum of star forming galaxies shows substantially larger variations in spectral slope in the mid-IR \citep[see][]{Brandl06} than in the FIR \citep{Dale06}.
". A sinall change in the geometry of the cnutting regions would affect Ελ μια). and thus qo, ratio mich inore than the FIR enüssiou aud qeu.", A small change in the geometry of the emitting regions would affect $_{\nu}$ $\mu$ m) and thus $_{24}$ ratio much more than the FIR emission and $_{\rm FIR}$.
" A similar result has also been noticed bv Appletonctal.(2001). and Murphyetal.(2006a) who found a larger dispersion i qo, as compared to οτι” and suggested that this is probably due to a larecr intrinsic dispersion iu the IR/radio correlation at shorter waveleugths.", A similar result has also been noticed by \citet{Appleton04} and \citet{Murphy06a} who found a larger dispersion in $_{24}$ as compared to $_{70}$ and suggested that this is probably due to a larger intrinsic dispersion in the IR/radio correlation at shorter wavelengths.
" Taterestinely. the Appletonetal.(2001) values forWwe qo,= 0.510.255 or the k-corrected qo, of LET are soniewhat sinaller than our results. though consistent within 26."," Interestingly, the \citet{Appleton04} values for $_{24}=$ $\pm$ 0.28 or the k-corrected $_{24}$ of $\pm$ 0.23 are somewhat smaller than our results, though consistent within $\sigma$ ."
" Oue possible explanation is that the dust temperature of low hnuuiuositv chwart ealaxies tends to peak at shorter wavelength than normal star forming galaxies. which would result iu au elevated 2|, based hIuninositv."," One possible explanation is that the dust temperature of low luminosity dwarf galaxies tends to peak at shorter wavelength than normal star forming galaxies, which would result in an elevated $\mu$m based luminosity."
 An alternative explanation is, An alternative explanation is
bent and wound close to the photon sphere.,bent and wound close to the photon sphere.
 As a result. a number of relativistic images appear around the photon sphere.," As a result, a number of relativistic images appear around the photon sphere."
 However. it has been shown that there is no photon sphere (Dev&Sen2003) in Ellis wormhole lensing.," However, it has been shown that there is no photon sphere \citep{dey08} in Ellis wormhole lensing."
 Therefore. (here is no contribution of relativistic images {ο the magnification in Ellis wormhole lensing.," Therefore, there is no contribution of relativistic images to the magnification in Ellis wormhole lensing."
 We thus conclude that the weak-lield hypothesis is a good approximation unless (he throat radius is comparable to the galactic distance., We thus conclude that the weak-field hypothesis is a good approximation unless the throat radius is comparable to the galactic distance.
" The probability of a microlensing event to occur lor a star is expressed by the optical depth 7: where n(D,) is the number density of wormboles as a fhunction of the line of sight.", The probability of a microlensing event to occur for a star is expressed by the optical depth $\tau$ : where $n(D_L)$ is the number density of wormholes as a function of the line of sight.
 Here we simply assume that ας) is constant (n(D)= ny The event rate expected for a source star Dis calculated as There is no reliable prediction of the number densitv of wormholes., Here we simply assume that $n(D_L)$ is constant $n(D_L) = n$ ): The event rate expected for a source star $\Gamma$ is calculated as There is no reliable prediction of the number density of wormholes.
 Several authors (Ixrasnikov2000:Lobo.2007) have speculated Chat wormnholesare verv common in the," Several authors \citep{kra00, lob09} have speculated that wormholesare very common in the"
Fig.,Fig.
 2 shows the average Spectral Energy Distribution (SED) of MIPS-detected LBGs and BX/BM galaxies., 2 shows the average Spectral Energy Distribution (SED) of MIPS-detected LBGs and $/$ BM galaxies.
" The SEDs have been constructed using available ‘averaged’UGRViJK, IRAC, MIPS and SPIRE flux measurements."," The SEDs have been constructed using available `averaged', IRAC, MIPS and SPIRE flux measurements."
 For the MIPS-LBGs we also use the 1.1 mm Aztec measurement from Magdis et al (2010b)., For the MIPS-LBGs we also use the 1.1 mm Aztec measurement from Magdis et al (2010b).
" We fit the optical/near-infrared part with model SEDs generated using the Bruzual Charlot (2003, BC03) code, while the mid-to-far infrared part is fit using Chary Elbaz (2001, CE01) template SEDs."," We fit the $/$ near-infrared part with model SEDs generated using the Bruzual Charlot (2003, BC03) code, while the mid-to-far infrared part is fit using Chary Elbaz (2001, CE01) template SEDs."
" In brief, we use BC03 and construct stellar population models with a Salpeter IMF and constant star formation rate, which has been shown (van Dokkum et al."," In brief, we use BC03 and construct stellar population models with a Salpeter IMF and constant star formation rate, which has been shown (van Dokkum et al."
" 2004, Rigopoulou et al."," 2004, Rigopoulou et al."
" 2006, Lai et al."," 2006, Lai et al."
 2007) to provide an adequate description of the properties of high redshift galaxies with ongoing star formation., 2007) to provide an adequate description of the properties of high redshift galaxies with ongoing star formation.
" Age, stellar mass, dust reddening and star formation rates are then derived from the model fits."," Age, stellar mass, dust reddening and star formation rates are then derived from the model fits."
" It is beyond the scope of the present work to discuss these results, a detailed analysis of the properties of the stellar population inSpitzer detected LBGs can be found in e.g. Rigopoulou et al. ("," It is beyond the scope of the present work to discuss these results, a detailed analysis of the properties of the stellar population in detected LBGs can be found in e.g. Rigopoulou et al. ("
"2006), Magdis et al. (","2006), Magdis et al. ("
"20102, for LBGs) and Reddy et al. (","2010a, for LBGs) and Reddy et al. ("
"2006, for BX/BMs).","2006, for $/$ BMs)."
" It is however, worth noting the differences in the optical part of the SED with the BX/BM galaxies showing a much ‘bluer’ SED."," It is however, worth noting the differences in the optical part of the SED with the $/$ BM galaxies showing a much `bluer' SED."
" We fit the far-IR/submm part with templates from the CEOI library, with the best-fit templates rendering mean (Lin) values of 2.8(3-0.6)x101? Lo for MIPS-LBGs and 1.5(+0.5)x10! Le for MIPS-BX/BMs."," We fit the $/$ submm part with templates from the CE01 library, with the best-fit templates rendering mean $\langle L_{\rm IR} \rangle$ values of $\pm$ $ \times 10^{12}$ $_{\odot}$ for MIPS-LBGs and $\pm$ $ \times 10^{11}$ $_{\odot}$ for $/$ BMs."
 The derived averaged Lim for LBGs is typical of those seen in Ultraluminous Infrared Galaxies (ULIRGs)., The derived averaged $L_{\rm IR}$ for LBGs is typical of those seen in Ultraluminous Infrared Galaxies (ULIRGs).
" Using the infrared luminosities we derive average Star Formation Rates (SFRs) of 296Mo/yr and 245M /yr, for the z ~3 LBGs and BX/BM galaxies, respectively."," Using the infrared luminosities we derive average Star Formation Rates $\langle SFRs \rangle$ of $_{\odot}/$ yr and $_{\odot}/$ yr, for the $z\sim$ 3 LBGs and $/$ BM galaxies, respectively."
 The <SFR> derived from the IR for LBGs is in agreement with the radio SFR estimate (280--85M. , The $<$ $>$ derived from the IR for LBGs is in agreement with the radio SFR estimate $\pm$ $_{\odot}$ 
RV.,RV.
 In Fig., In Fig.
" 2 the histogram of the radial velocity distribution of these stars is shown, together with a Gaussian fit with «Vaga >=32.90 km s! and ao=0.73 km s!."," \ref{fig:vrad_distribution_sel} the histogram of the radial velocity distribution of these stars is shown, together with a Gaussian fit with $<V_{\rm rad}>$ =32.90 km $^{-1}$ and a $\sigma=0.73$ km $^{-1}$."
" In Table A4 the RV values are listed for the stars of the final sample, while in Table A5 the values of the single RV measurements are given for the stars we discarded."," In Table \ref{tab:targets_selected} the RV values are listed for the stars of the final sample, while in Table \ref{tab:nonmembers} the values of the single RV measurements are given for the stars we discarded."
 In Fig., In Fig.
 3 we show the enlarged portion of the colour-magnitude diagram CMD containing the original sample; in this Figure the discarded and the retained stars are indicated with different colours., \ref{fig:cmdenlarged} we show the enlarged portion of the colour-magnitude diagram CMD containing the original sample; in this Figure the discarded and the retained stars are indicated with different colours.
" Many of the discarded stars tend to occupy the brighter side of the main sequence, where binaries are indeed expected to be present."," Many of the discarded stars tend to occupy the brighter side of the main sequence, where binaries are indeed expected to be present."
" On the other hand, our procedure still leaves several stars which are apparently above the photometric main sequence."," On the other hand, our procedure still leaves several stars which are apparently above the photometric main sequence."
 This is because the radial velocity measurements are not of superb quality and because the time span by the observations is of only 18 days., This is because the radial velocity measurements are not of superb quality and because the time span by the observations is of only 18 days.
 Long period binaries will not be discovered by our three radial velocity observations., Long period binaries will not be discovered by our three radial velocity observations.
 We shall see as seven stars clearly stand up also in the Magnitude — Temperature diagram (see Fig. 5)), We shall see as seven stars clearly stand up also in the Magnitude – Temperature diagram (see Fig. \ref{fig:Teff_V}) )
 and they are best candidates for binaries of similar mass., and they are best candidates for binaries of similar mass.
" We have kept them in the sample, and we anticipate that their presence does not influence our analysis or conclusions."," We have kept them in the sample, and we anticipate that their presence does not influence our analysis or conclusions."
" Given that our targets are on the main sequence of a cluster of solar metallicity and age, the critical astrophysical parameters for the selection of the best solar analogues is the effective temperature."," Given that our targets are on the main sequence of a cluster of solar metallicity and age, the critical astrophysical parameters for the selection of the best solar analogues is the effective temperature."
 We have used two spectroscopic methods to compute the stellar effective temperature: the line-depth ratios and the Ha wings., We have used two spectroscopic methods to compute the stellar effective temperature: the line-depth ratios and the $\alpha$ wings.
" To calibrate these methods we have used a grid of synthetic spectra, computed with SYNTHE from a grid of ID LTE model atmospheres computed with version 9 of the ATLAS code (Kurucz 1993a,b)) in its Linux version (Sbordoneetal.2004;Sbordone 2005))"," To calibrate these methods we have used a grid of synthetic spectra, computed with SYNTHE from a grid of 1D LTE model atmospheres computed with version 9 of the ATLAS code \citealt{Kurucz1993a,Kurucz1993b}) ) in its Linux version \citealt{Sbord2004,Sbord2005}) )."
" All the models have been computed with the “NEW” Opacity Distribution Functions (Castelli&Kurucz 2003)) which are based on solar abundances from Grevesse&Sauval(1998) with mmicro-turbulence, a mixing-length parameter oof 1.25 and no overshooting."," All the models have been computed with the “NEW” Opacity Distribution Functions \citealt{CK03}) ) which are based on solar abundances from \cite{Grevesse1998} with micro-turbulence, a mixing-length parameter of 1.25 and no overshooting."
" The grid of synthetic spectra covers the temperature range 5450-6300 K with [Fe/H]=0, log g=4.4377, €=1 km s! and was degraded to the resolution of the FLAMES/GIRAFFE spectra."," The grid of synthetic spectra covers the temperature range 5450–6300 K with [Fe/H]=0, $\log g$ =4.4377, $\xi$ =1 km $^{-1}$ and was degraded to the resolution of the FLAMES/GIRAFFE spectra."
 We stress that for both methods these models are used to quantify the difference between the stellar spectra and the solar spectrum., We stress that for both methods these models are used to quantify the difference between the stellar spectra and the solar spectrum.
" Zero point shifts are most likely present, due, for instance, to limitations in the atmospheric models or to not perfect treatment of the Ha lines."," Zero point shifts are most likely present, due, for instance, to limitations in the atmospheric models or to not perfect treatment of the $\alpha$ lines."
" While these inaccuracies will reflect in a wrong temperature for the Sun, the difference between the stars and the Sun will be much less affected."," While these inaccuracies will reflect in a wrong temperature for the Sun, the difference between the stars and the Sun will be much less affected."
" It has been demonstrated that for stars with B—V=0.41.5 line-depth ratios (LDRs) are a powerful temperature indicator, capable to resolve temperature differences lower than 10 K (Gray&Johanson1991;Catalanoetal.2002;Biazzo 2007))."," It has been demonstrated that for stars with $B-V=0.4-1.5$ line-depth ratios (LDRs) are a powerful temperature indicator, capable to resolve temperature differences lower than 10 K \citealt{GrayJoha1991,Catalano2002,Biazzo2007}) )."
" Since our stars are within this B—V range, we have applied the LDR method to the members previously selected by radial velocity measurements (see Table A4))."," Since our stars are within this $B-V$ range, we have applied the LDR method to the members previously selected by radial velocity measurements (see Table \ref{tab:targets_selected}) )."
 To convert the line-depth ratios of our stars into effective temperature we need to calibrate a temperature scale for the measured LDRs., To convert the line-depth ratios of our stars into effective temperature we need to calibrate a temperature scale for the measured LDRs.
" To this purpose we have considered an initial sample of about 100 lines of iron group elements (which are usually temperature sensitive) present in the spectral range covered by our observations, from which we selected lines with the following characteristics: weak (to avoid saturation effects),"," To this purpose we have considered an initial sample of about 100 lines of iron group elements (which are usually temperature sensitive) present in the spectral range covered by our observations, from which we selected lines with the following characteristics: weak (to avoid saturation effects),"
of the E-modes into B-modes is still low.,of the E-modes into B-modes is still low.
" At higher multipoles however, the bandwidth smearing effect and loss of coherence would be a real issue for bolometric interferometers while in a heterodyne interferometer, the separation into small bands would prevent the sensitivity from dropping."," At higher multipoles however, the bandwidth smearing effect and loss of coherence would be a real issue for bolometric interferometers while in a heterodyne interferometer, the separation into small bands would prevent the sensitivity from dropping."
 The main remaining question is whether the gain in terms of systematic effects is worth the price of this sensitivity reduction if one builds an interferometer instead of an imager., The main remaining question is whether the gain in terms of systematic effects is worth the price of this sensitivity reduction if one builds an interferometer instead of an imager.
" In terms of optics for instance, an interferometer directly observes the sky."," In terms of optics for instance, an interferometer directly observes the sky."
" The primary beam is therefore only set by that of the horns, while in an imager, the telescope (mirror or lenses) produces sidelobes inducing poorly predictible ground pickup that often prevent one from reaching the nominal sensitivity."," The primary beam is therefore only set by that of the horns, while in an imager, the telescope (mirror or lenses) produces sidelobes inducing poorly predictible ground pickup that often prevent one from reaching the nominal sensitivity."
 An interferometer is also completely insensitive to spatially uniform polarized signals that vary with time such as polarized atmospheric contamination., An interferometer is also completely insensitive to spatially uniform polarized signals that vary with time such as polarized atmospheric contamination.
 These could also prevent an imager from reaching its nominal sensitivity by adding some spread in the noise., These could also prevent an imager from reaching its nominal sensitivity by adding some spread in the noise.
 These examples mitigate the statistical sensitivity loss of an interferometer with respect to an imager., These examples mitigate the statistical sensitivity loss of an interferometer with respect to an imager.
" The differences in terms of systematic effects between imagers and bolometric and heterodyne interferometers are not obvious and deserve a detailed quantitative study in continuation of the work done by (Bunn, 2007)).", The differences in terms of systematic effects between imagers and bolometric and heterodyne interferometers are not obvious and deserve a detailed quantitative study in continuation of the work done by \cite{bunn}) ).
NGC 1068. they do so iu a regulu. haxinonicallv related or “quantized”. manner.,"NGC 1068, they do so in a regular, harmonically related or ""quantized"", manner."
 The quantization is in units of Az = 0.05. decreasing outwards. as the triplet size increases. (in redshift steps of z = Az. 2Az. 3Az....) frou a Πακπια redshift of Zuican = 1.25 near the central poiut. to a uiunimadl Of μαι = 0.5 im the most remote triplet.," The quantization is in units of $\Delta$ z = 0.05, decreasing outwards, as the triplet size increases, (in redshift steps of z = $\Delta$ z, $\Delta$ z, $\Delta$ z,...) from a maximum redshift of $_{\rm mean}$ = 1.25 near the central point, to a minimum of $_{\rm mean}$ = 0.5 in the most remote triplet."
 This can be seen nore clearly i Fie., This can be seen more clearly in Fig.
 7. where the μμ values have been plotted (filed circles) versus the mean angular displacement of the sources iu cach triplet roni NGC 1068 (col 6 in Table 2).," 7, where the $_{\rm mean}$ values have been plotted (filled circles) versus the mean angular displacement of the sources in each triplet from NGC 1068 (col 6 in Table 2)."
 The ocations of perfecectly quantized redshifts are imcdicated by he horizoital clasje lines and show good agreement with he ineasur‘ed values., The locations of perfectly quantized redshifts are indicated by the horizontal dashed lines and show good agreement with the measured values.
 The “quantized” redshifts are simply fhe mean of two neasured values.," The ""quantized"" redshifts are simply the mean of two measured values."
 Also inuplici oeji this figure. is he indicatioi that fwre ds a ndsjsune triplet with Zucan PCCIshit near 0.95., Also implicit in this figure is the indication that there is a missing triplet with $_{\rm mean}$ redshift near 0.95.
 Although its redshitt Las rot been measred. source 9 is located at aporoxmnuatev the correct aueular distance (22/6) ;m NGCI16S to be part of such a triplet.," Although its redshift has not been measured, source 9 is located at approximately the correct angular distance $22\farcm6$ ) from NGC 1068 to be part of such a triplet."
 Siice. the singlet and pair imidpoiut iu cach riplet are separating from one another at different aneles to fi6 bo- (eiven by 3). the mca shifts of the 1| pairs would be expected to lave superiniposed Doppler componeuts.," Since the singlet and pair midpoint in each triplet are separating from one another at different angles to the l-o-s (given by $\beta$ ), the mean redshifts of the 4 pairs would be expected to have superimposed Doppler components."
" It is thus 10 clear vet whether the quantization is a real quantization di an intrinsic reshift couponcut. or simply due to some fortuitous colmbination of Doppler aud intrinsic redshifts,"," It is thus not clear yet whether the quantization is a real quantization in an intrinsic redshift component, or simply due to some fortuitous combination of Doppler and intrinsic redshifts."
 Oltainine redshifts for sources 10 and 12 should help to clarify this., Obtaining redshifts for sources 10 and 12 should help to clarify this.
 In Fig., In Fig.
 7. as the recshifts decrease 1u what appear to be quautized steps. the distance to the relevaut triplets increases iu a simular manner such that the relation between increasing aneular displacement aud miean-paidr redshift varies sinoothly and approximately linearily.," 7, as the redshifts decrease in what appear to be quantized steps, the distance to the relevant triplets increases in a similar manner such that the relation between increasing angular displacement and mean-pair redshift varies smoothly and approximately linearily."
 A simular quasi-linear relation was found previously for sources near NGC 3516 (Chuctal.1998)., A similar quasi-linear relation was found previously for sources near NGC 3516 \citep{chu98}.
. Note that this quautization appears in only onc conponent (Zuge) of the measured recshifts., Note that this quantization appears in only one component $_{\rm mean}$ ) of the measured redshifts.
" Since the systemic redshift of NCC LOGS is known. the quantized couponeuts cannot contain a large ""uuknown cosmokigical component."," Since the systemic redshift of NGC 1068 is known, the quantized components cannot contain a large unknown cosmological component."
 This appareut quantization is tlrerefore unlikely to be related to the periodicity revorted in other studies (BurbidecandNapicr2001:Ikarlssou1977) where the entire measured redshifts lave been usec.," This apparent quantization is therefore unlikely to be related to the periodicity reported in other studies \citep{bur01,kar77} where the entire measured redshifts have been used."
 e) Is there a fifth triplet?, e) Is there a fifth triplet?
 The imauuer in which the singlet redsfts vary with distance from NCC 1OS cannot be determined vet because the redshifts of singlets 10 aud 12 have not been imeastred., The manner in which the singlet redshifts vary with distance from NGC 1068 cannot be determined yet because the redshifts of singlets 10 and 12 have not been measured.
 In Fig., In Fig.
"B. 7 he nuüssug triplet postulated above has Όσσα abeled E and ds predicted to be located at an augular distance of ~22"" from NCC 1068.", 7 the missing triplet postulated above has been labeled E and is predicted to be located at an angular distance of $\sim22\arcmin$ from NGC 1068.
 This mforiation. aud position information fron he other four triplets. is Όσιο examined iui an attempt to piupoiut on the sky the most likely ocatious of the two sources required to conrete viplet E (Bell. 2001. iu preparation).," This information, and position information from the other four triplets, is being examined in an attempt to pinpoint on the sky the most likely locations of the two sources required to complete triplet E (Bell, 2001, in preparation)."
 f) Proper motions of the ejected sources., f) Proper motions of the ejected sources.
 The age of triplet A can be used together with, The age of triplet A can be used together with
"Throughout this paper we adopt the ? cosmological parameters used in our simulations to convert redshifts to distances: (0,,.Oy.O4.n;σε.) = (0.26. 0.044. 0.74. 0.95. 0.71. 0.72).","Throughout this paper we adopt the \citet{spergel/etal:2007} cosmological parameters used in our simulations to convert redshifts to distances: $\Omega_m, \Omega_b, \Omega_{\Lambda}, n_s, \sigma_8, h$ ) = (0.26, 0.044, 0.74, 0.95, 0.77, 0.72)."
 All distances and separations are in comoving coordinates., All distances and separations are in comoving coordinates.
 The goal of Chis section is to measure (he eroup multiplicity hunction of a subsample of the SDSS LRGs., The goal of this section is to measure the group multiplicity function of a subsample of the SDSS LRGs.
" For a complete spectroscopic sample covering the full skv (or in a periodic simulation box). our method is as follows: In 2.2 we present the technical details of accounting for the facts that the SDSS has boundaries aud holes. that (he spectroscopic sample of LRGs is incomplete. ancl that the SDSS cannot simultaneously take spectra of two objects separated by <55"". so that regions ol (he sky observed only once spectroscopically may have missing close pairs of LRGs."," For a complete spectroscopic sample covering the full sky (or in a periodic simulation box), our method is as follows: In \ref{data} we present the technical details of accounting for the facts that the SDSS has boundaries and holes, that the spectroscopic sample of LRGs is incomplete, and that the SDSS cannot simultaneously take spectra of two objects separated by $< 55''$, so that regions of the sky observed only once spectroscopically may have missing close pairs of LRGs."
 We use the LRGs from the SDSS imaging sample to supplement the spectroscopic sample (??)..," We use the LRGs from the SDSS imaging sample to supplement the spectroscopic sample \citep{blanton/etal:2005, adelman-mccarthy/etal:2007}."
 We identify potential pairs from the imaging sample and calibrate (is step using pairs of objects from the spectroscopic sample., We identify potential pairs from the imaging sample and calibrate this step using pairs of objects from the spectroscopic sample.
 Since most nonisolated LRGs are in groups of 2 and candidate pairs from (he imaging sample neiehboring more (han one LAG ave highly likely to be eroup menbers. we apply the small correction for false LRG pair detections to INejc(ni=1).," Since most nonisolated LRGs are in groups of 2 and candidate pairs from the imaging sample neighboring more than one LRG are highly likely to be group members, we apply the small correction for false LRG pair detections to $N_{CiC}(n_{sat} = 1)$."
" A more complex scheme involving corrections at each ny, would not have enough statistics to calibrate on the spectroscopic sample.", A more complex scheme involving corrections at each $n_{sat}$ would not have enough statistics to calibrate on the spectroscopic sample.
 We eliminate [from our sample LRGs close to the survey boundary. though they are allowed to be grouped wilh LRGs away from (he boundary.," We eliminate from our sample LRGs close to the survey boundary, though they are allowed to be grouped with LRGs away from the boundary."
 This ensures (hat our multiplicity AieGon is not biased due to unobserved LRGs outside the boundary., This ensures that our multiplicity function is not biased due to unobserved LRGs outside the boundary.
 However. since the bright star masks are numerous and individually. verv small. (his approach is not practical for dealing with objects near bright star masks.," However, since the bright star masks are numerous and individually very small, this approach is not practical for dealing with objects near bright star masks."
 Instead. we adjust Nga) bv estimating (the probability that there is an LRG covered by each bright star mask. and then computing the change in ο) if there were one.," Instead, we adjust $N(n_{sat})$ by estimating the probability that there is an LRG covered by each bright star mask, and then computing the change in $N(n_{sat})$ if there were one."
 Table 1. shows that, Table \ref{table:sdssmulttable} shows that
Although both samples are not exactly on the same scale at the high metallicity end (ef section 5.1.4 of Paper ID. the mean metallicities for the metal-poor population ts extremely similar in both samples (see also the decomposition of the full Baade's Window sample on the same metallicity scale at the end of section 3.1).,"Although both samples are not exactly on the same scale at the high metallicity end (cf section 5.1.4 of Paper II), the mean metallicities for the metal-poor population is extremely similar in both samples (see also the decomposition of the full Baade's Window sample on the same metallicity scale at the end of section 3.1)."
 At b—6* the Wilks’ test allows to keep a solution with 3 components presented in Table 5 rather tha the 2 components one., At $b=-6\degr$ the Wilks' test allows to keep a solution with 3 components presented in Table \ref{tab:b6decomp} rather than the 2 components one.
 Population A and population B coule correspond to the population A and B observed in Baade's Window., Population A and population B could correspond to the population A and B observed in Baade's Window.
 The mean metallicities are coherent although their spread is smaller., The mean metallicities are coherent although their spread is smaller.
 The radial velocity dispersions of population A are identical., The radial velocity dispersions of population A are identical.
 The radial velocity dispersion of population B decreases at b.=—67 as expected for a bar-like kinematic behaviour (see Fig., The radial velocity dispersion of population B decreases at $b=-6\degr$ as expected for a bar-like kinematic behaviour (see Fig.
 8 and associated text)., \ref{fig:compFuxSigVr} and associated text).
 Population C represents only 6+2% of the sample with a low mea metallicity of —1.1+0.1 dex and a high velocity dispersion of 127426 km/s. which could therefore be associated to the halo.," Population C represents only $6\pm2$ of the sample with a low mean metallicity of $-1.1\pm0.1$ dex and a high velocity dispersion of $\pm$ 26 km/s, which could therefore be associated to the halo."
 The Besangoon model prediction of only 0.440.1% of halo star in the sample could therefore have been underestimated., The Besançoon model prediction of only $0.4\pm0.1$ of halo star in the sample could therefore have been underestimated.
 This population could also have been hidden in the population A at b=—4°., This population could also have been hidden in the population A at $b=-4\degr$.
" Selecting all stars with [Fe/H]<—0.9 in our 3 fields we obtain 16 stars with a radial velocity dispersion of c,=116x2} km/s. which is coherent withthe solar neighbourhooc velocity dispersions measured in this metallicity range (e.g. 2 cp~11010 km/s) containing both thick dise. and halo stars."," Selecting all stars with $<-0.9$ in our 3 fields we obtain 16 stars with a radial velocity dispersion of $\sigma_r= 116\pm21$ km/s, which is coherent withthe solar neighbourhood velocity dispersions measured in this metallicity range (e.g. \cite{ChibaBeers00}: $\sigma_U \sim 110 \pm 10$ km/s) containing both thick disc and halo stars."
 We are not in a position to clearly associate this population either to the halo or to a metal-poor thick disc ii our inner galactic samples., We are not in a position to clearly associate this population either to the halo or to a metal-poor thick disc in our inner galactic samples.
 SEMMUL did not converge on the b=—12° field due to the smaller number of stars and the higher contamination with thin. thick dises and halo stars expected 11 this field.," SEMMUL did not converge on the $b=-12\degr$ field due to the smaller number of stars and the higher contamination with thin, thick discs and halo stars expected in this field."
 At b.=—12? the metal rich component present at b=-4° and b=—-6° seems to have fully disappeared., At $b=-12\degr$ the metal rich component present at $b=-4\degr$ and $b=-6\degr$ seems to have fully disappeared.
 The metal rich velocity part of the velocity dispersion corresponds toa disce like component (see next section) while the metal poor part shows a velocity dispersion still coherent with the metal poor population of b.=—4* and b= -6*. although we cannot distinguish a spheroid and a thick dise contribution.," The metal rich velocity part of the velocity dispersion corresponds to a disc like component (see next section) while the metal poor part shows a velocity dispersion still coherent with the metal poor population of $b=-4\degr$ and $b=-6\degr$ , although we cannot distinguish a spheroid and a thick disc contribution."
found. for samples including lower luminosity objects (Fan et al.,found for samples including lower luminosity objects (Fan et al.
 2001a)., 2001a).
 The Sloan Digital Sky Survey Data Release 3 (SDSS. Stoughton ct al.," The Sloan Digital Sky Survey Data Release 3 (SDSS, Stoughton et al."
 2002: DRS Abazajian et al., 2002; DR3 Abazajian et al.
 2005) provides moderately deep CCD imaging in five bands igqriz covering ~ 5282 dee?, 2005) provides moderately deep CCD imaging in five bands $ugriz$ covering $\sim$ 5282 $^2$.
 We present and ciseuss a new sample of high- racio-sclected QSO. candidates in a 1378.5. deg? area of overlap between. FIRST and SDSS DI in the north Galactic cap., We present and discuss a new sample of high-redshift radio-selected QSO candidates in a 1378.5 $^2$ area of overlap between FIRST and SDSS DR3 in the north Galactic cap.
 The selection criteria are: 1) £x19.1 and starlike in ΑΝ d) σεν6g;c1 ην. di) optical separation less than 1.5 aresee. iv) colour οLE>2 (including ο non-detections) and v) starlike in SDSS.," The selection criteria are: i) $E \le 19.1$ and starlike in APM, ii) $S_{\rm 1.4 \ GHz}
\ge 1$ mJy, iii) radio-optical separation less than 1.5 arcsec, iv) colour $O - E \ge 2$ (including $O$ non-detections) and v) starlike in SDSS."
 ‘This new sample. with a wider colour range. has several advantages.," This new sample, with a wider colour range, has several advantages."
 Firstly. the SDSS photometric catalogue provides reliable morphological classification of the sources. allowing us to readily eliminate the galaxies. classecl as starlike in APAL POSS-I. Secondly. SDSS provides  3 A--resolution spectra and spectroscopic classifications of many objects. particularly those selected. as QSO candidates on the basis of their ngriz colours. or as counterparts of FIRST sources.," Firstly, the SDSS photometric catalogue provides reliable morphological classification of the sources, allowing us to readily eliminate the galaxies classed as starlike in APM POSS-I. Secondly, SDSS provides $\sim$ 3 -resolution spectra and spectroscopic classifications of many objects, particularly those selected as QSO candidates on the basis of their $ugriz$ colours, or as counterparts of FIRST sources."
 We continue to use the APAL catalogue for colour selection since our previous work showed a high cllicicney and completeness in theselection of z>3.85 QSOs using O E3and with the new limit O—£c2 we can check that no z> 3.7 QSOs have O—£2« 3., We continue to use the APM catalogue for colour selection since our previous work showed a high efficiency and completeness in theselection of $z > 3.85$ QSOs using $O-E \ge 3$ and with the new limit $O-E \ge 2$ we can check that no $z >$ 3.7 QSOs have $O-E <$ 3.
 The paper is structured. as follows., The paper is structured as follows.
 In. Section 2 we present the sample anc the status of the. spectroscopic Classification., In Section 2 we present the sample and the status of the spectroscopic classification.
 Section 3 reports optical spectroscopy of part of the sample., Section 3 reports optical spectroscopy of part of the sample.
 Lhe spectroscopic classification of the sample as QSOs. narrow emission line galaxies or stars is presented in Section 4.1.," The spectroscopic classification of the sample as QSOs, narrow emission line galaxies or stars is presented in Section 4.1."
 Vhe distribution of optical magnitudes and O—E colours is discussed in Section 4.2., The distribution of optical magnitudes and $O-E$ colours is discussed in Section 4.2.
 In Section. 4.3 the sample is compared: with previous racio-selected QSO samples from the literature. in terms of the selection criteria and the resulting QSO redshift distribution.," In Section 4.3 the sample is compared with previous radio-selected QSO samples from the literature, in terms of the selection criteria and the resulting QSO redshift distribution."
 In Section 4.4 we comment brielly on the spectra of seven QSOs exhibiting strong blueshiftecl broad. absorption lines (BALs) ancl we analyse the fraction of BAL QSOs in the sample., In Section 4.4 we comment briefly on the spectra of seven QSOs exhibiting strong blueshifted broad absorption lines (BALs) and we analyse the fraction of BAL QSOs in the sample.
 Section 4.5 is devoted to the peculiar QSO FIRST 1413|4505., Section 4.5 is devoted to the peculiar QSO FIRST 1413+4505.
 In Section 4.6 we compute the absolute magnitudes. &-corrections and radio luminosities of the LO QSOs with -&i3.7.or and we discuss the completeness of a sub-sample of seven of them.," In Section 4.6 we compute the absolute magnitudes, $k$ -corrections and radio luminosities of the 10 QSOs with $z \ge 3.7$, and we discuss the completeness of a sub-sample of seven of them."
 In Section 5 we use this sample to calculate the space density of QSOs., In Section 5 we use this sample to calculate the space density of QSOs.
 Section 6 summarizes our conclusions., Section 6 summarizes our conclusions.
 The sample was selected [roni the 1378.5-deg?. area celine: in Table 1., The sample was selected from the $^2$ area defined in Table 1.
 “Phis area includes most of the region covered by SDSS DRS in the north Galactic cap. which is also coverec by the FIRST survey and by the APM catalogue of POSS-1.," This area includes most of the region covered by SDSS DR3 in the north Galactic cap, which is also covered by the FIRST survey and by the APM catalogue of POSS-I."
" The FIRST survey includes 122463 sources in this area with 5,4cuz (peak) zc 1 my. of which 113 have APA Lc 19.1. οLo 2.0 and lie within 1.5 arcsec of starlike objects in SDSS DIU."," The FIRST survey includes 122463 sources in this area with $S_{\rm 1.4 \ GHz}$ (peak) $\ge$ 1 mJy, of which 113 have APM $E <$ 19.1, $O-E >$ 2.0 and lie within 1.5 arcsec of starlike objects in SDSS DR3."
 Eighteen of these were undetected in APM POSS-I O but were detected in APS POSS-I O anc had APS(O—E) «2. and were therefore removed from the sample.," Eighteen of these were undetected in APM POSS-I $O$ but were detected in APS POSS-I $O$ and had $O-E$ ) $ < 2$, and were therefore removed from the sample."
 For the source FIRST 134015619 we found a large dilference. between. the SDSS and. APAL magnitudes (r=23.51 versus £=18.99) and the source was eliminate from the sample after confirming with SDSS that the APA counterpart is a blend., For the source FIRST 1340+5619 we found a large difference between the SDSS and APM magnitudes $r=23.51$ versus $E=18.99$ ) and the source was eliminated from the sample after confirming with SDSS that the APM counterpart is a blend.
 The final sample thus includes. 94 candidate high-redshift radio QSOs., The final sample thus includes 94 candidate high-redshift radio QSOs.
 Of these 94 candidates. spectra were first obtained for seven in papers b and . for six in the literature (founc using the NASA Extragalactic Database - NIZD) and for 41 by SDSS DRS (which also reobserved nine of the thirteen previously discovered).," Of these 94 candidates, spectra were first obtained for seven in papers I and II, for six in the literature (found using the NASA Extragalactic Database - NED) and for 41 by SDSS DR3 (which also reobserved nine of the thirteen previously discovered)."
" In Section 3 we present (PNG optica spectroscopy of 13 Grandonmly selected) of the 40 remaining candidates. and we classify an SDSS DIU source given spectral class ""unknown! in the Sloan survey."," In Section 3 we present TNG optical spectroscopy of 13 (randomly selected) of the 40 remaining candidates, and we classify an SDSS DR3 source given spectral class `unknown' in the Sloan survey."
 Subsequen to these observations. SDSS DIt4 (2005 June 30) reporte spectroscopy of LO of the remaining 27 candidates Cane also of four of those observed here).," Subsequent to these observations, SDSS DR4 (2005 June 30) reported spectroscopy of 10 of the remaining 27 candidates (and also of four of those observed here)."
 In total. 78 of the 94 candidates (83 per cent) are now spectroscopically elassifiec (sec Section 4.1).," In total, 78 of the 94 candidates (83 per cent) are now spectroscopically classified (see Section 4.1)."
 Spectra of 11 candidates (indicated in column 9 of Table 2) were obtained with the Telescopio Nazionale Galileo (LNG) on 2005 March 10 and 11 using the DOLORES (Device Optimized for LOw ItSolution) spectrograph in long-slit mode., Spectra of 11 candidates (indicated in column 9 of Table 2) were obtained with the Telescopio Nazionale Galileo (TNG) on 2005 March 10 and 11 using the DOLORES (Device Optimized for LOw RESolution) spectrograph in long-slit mode.
 The LR-B erism was used. viekding a wavelength range 3000 — SSOO aand dispersion 2.9 pixel+.," The LR-B grism was used, yielding a wavelength range 3000 – 8800 and dispersion 2.9 $^{-1}$."
 The detector was a thinned. back-illuminated Loral CCD with 15/3 [nupixels.," The detector was a thinned, back-illuminated Loral CCD with $\mu$ m pixels."
 Exposure1 times were typically1 900 s. ‘Three spectrophotometric standard. stars were observed. in order to calibrate the instrumental spectral response., Exposure times were typically 900 s. Three spectrophotometric standard stars were observed in order to calibrate the instrumental spectral response.
 The seeing was z Ls aresec. and the width of 10 slit was set to 2 aresec. vielding a spectral resolution of 23 X.. as measured from sky lines.," The seeing was $\approx$ 1.8 arcsec, and the width of the slit was set to 2 arcsec, yielding a spectral resolution of 23 , as measured from sky lines."
 Standard: cata reduction was carried out using the package., Standard data reduction was carried out using the package.
 Are-lamp exposures were used for the wavelength calibration. ancl vielded solutions with rms residuals « 1.5," Arc-lamp exposures were used for the wavelength calibration, and yielded solutions with rms residuals $<$ 1.5"
The 3-hour period may alternatively be due to the rotation of2MO036+18.,The 3-hour period may alternatively be due to the rotation of.
. Using the inferred radius of 0.09 R.. the equatorial velocity is 37 km s7!. leading to an inclination angle of 248°. or nearly pole-on.," Using the inferred radius of 0.09 $_\odot$, the equatorial velocity is 37 km $^{-1}$, leading to an inclination angle of $24\pm 8^\circ$, or nearly pole-on."
 The azimuthal ortentation. o. of the axis of rotation relative to the line of sight is not known.," The azimuthal orientation, $\phi$, of the axis of rotation relative to the line of sight is not known."
 However. given the low inclination and the large covering fraction. the radio-emitting region has to be located at a low latitude for most o values.," However, given the low inclination and the large covering fraction, the radio-emitting region has to be located at a low latitude for most $\phi$ values."
 This is because a location near the pole would make the bulk of the region visible throughout the rotation period. resulting in a nearly constant flux level.," This is because a location near the pole would make the bulk of the region visible throughout the rotation period, resulting in a nearly constant flux level."
 As can be seen from Figure 2. the flux drops to nearly zero between the peaks indicating the emission region is fully occulted for about half the rotation period., As can be seen from Figure \ref{fig:c} the flux drops to nearly zero between the peaks indicating the emission region is fully occulted for about half the rotation period.
 We note also that the low inclination may explain the constant sign of the circular polarization since for most values ofo only one of the hemispheres ts visible., We note also that the low inclination may explain the constant sign of the circular polarization since for most values of $\phi$ only one of the hemispheres is visible.
 If the 3-hour period is related to the rotation of tthen we can estimate the Rossby number. Ro=P/7.. which is relevant for dynamo models.," If the 3-hour period is related to the rotation of then we can estimate the Rossby number, $Ro=P/\tau_c$, which is relevant for dynamo models."
" The convective turnover time for iis 7=MR""/L)?0.8 yr indicating that Roz4.4«1077."," The convective turnover time for is $\tau_c=
(MR^2/L)^{1/3}\approx 0.8$ yr indicating that $Ro\approx 4.4\times
10^{-4}$."
 At such low values early M dwarfs exhibit saturated X-ray emission. Lyπρι~107 (Pizzolatoeraf2003).," At such low values early M dwarfs exhibit saturated X-ray emission, $L_X/L_{\rm bol}\sim 10^{-3}$ \citep{pmm+03}."
.. Clearly. the conditions in L dwarfs are sufficiently different that a low Rossby number does not result in X-ray emission.," Clearly, the conditions in L dwarfs are sufficiently different that a low Rossby number does not result in X-ray emission."
 We also note that the large value of 7. may explain the stability of the radio emission (flux and circular polarization) over a period of about three years., We also note that the large value of $\tau_c$ may explain the stability of the radio emission (flux and circular polarization) over a period of about three years.
 Finally. it is possible that the variable persistent emission is in fact a series of periodic flares with a rate of occurrence of about 0.33 hr!.," Finally, it is possible that the variable persistent emission is in fact a series of periodic flares with a rate of occurrence of about 0.33 $^{-1}$."
" The frequency of strongflares. similar to the one detected by BOL with a flux of about 720 sry. is <0.04 hr""."," The frequency of strongflares, similar to the one detected by B01 with a flux of about 720 $\mu$ Jy, is $\lesssim 0.04$ $^{-1}$."
 This pattern is similar to the case of the M9 dwarf 22065 for which strong Ha flares have an occurrence rate of <0.03 he! while weak flares occur at a rate as high as about 0.5 he! (Martín&Ardila2001)., This pattern is similar to the case of the M9 dwarf 2065 for which strong $\alpha$ flares have an occurrence rate of $\lesssim 0.03$ $^{-1}$ while weak flares occur at a rate as high as about 0.5 $^{-1}$ \citep{ma01}.
. If this is the case here. then the timescale to build up magnetic stresses is about 3 hours. while the lifetime of the electrons. corresponding to the width of the flares. is about | hour.," If this is the case here, then the timescale to build up magnetic stresses is about 3 hours, while the lifetime of the electrons, corresponding to the width of the flares, is about 1 hour."
 In this scenario it ts possible that the periodic weak flares are inefficient at heating the corona and chromosphere. thus explaining the lack of accompanying X-ray and Ha emission.," In this scenario it is possible that the periodic weak flares are inefficient at heating the corona and chromosphere, thus explaining the lack of accompanying X-ray and $\alpha$ emission."
 The more rare strong flares. on the other hand. with an impulsive energy release a few times larger may be accompanied by efficient heating.," The more rare strong flares, on the other hand, with an impulsive energy release a few times larger may be accompanied by efficient heating."
" However. we note that the time-integrated energy release in the putative flares observed here. E~vL,£4«1078 erg. is not so different from the integrated energy release inferred for the stronger flare (BO2)."," However, we note that the time-integrated energy release in the putative flares observed here, $E\sim \nu L_\nu t\sim 4\times 10^{26}$ erg, is not so different from the integrated energy release inferred for the stronger flare (B02)."
 The growing sample of late-M and L dwarfs observed in the radio and X-rays allows to investigate the patterns of activity às a function of relevant physical parameters., The growing sample of late-M and L dwarfs observed in the radio and X-rays allows to investigate the patterns of activity as a function of relevant physical parameters.
 Ha observations of late-M and L dwarfs indicate a significant drop in the level of emission compared to the bolometric luminosity (Figure 14))., $\alpha$ observations of late-M and L dwarfs indicate a significant drop in the level of emission compared to the bolometric luminosity (Figure \ref{fig:activity_sptype}) ).
" Objects earlier than about M6 tend to have a ratio Lj,/Ly; 1n the range of 107 to 107."," Objects earlier than about M6 tend to have a ratio $L_{H\alpha}/L_{\rm
bol}$ in the range of $10^{-4}$ to $10^{-3}$."
 However. by spectral type LO this ratio is typically lower than 107.," However, by spectral type L0 this ratio is typically lower than $10^{-5}$."
 The few late-M and L dwarfs which exhibit flares (~1%: Liebertetof. 2003)) appear to have levels of emission similar to those of early M dwarfs. with only asingle source exceeding that level to date.," The few late-M and L dwarfs which exhibit flares $\sim 1\%$; \citealt{lkc+03}) ) appear to have levels of emission similar to those of early M dwarfs, with only a single source exceeding that level to date."
 Thus. Ha emission drops significantly with a decreased surface temperature (later spectral type). and flares exceeding the saturation value are extremely rare.," Thus, $\alpha$ emission drops significantly with a decreased surface temperature (later spectral type), and flares exceeding the saturation value are extremely rare."
" Similarly. the ratio Ly/Lpo, drops from a saturation value of about 10? in mid-M dwarfs to less than 107 for the few late- dwarfs observed to date (e.g.. Flemingetaf.1993;Rut-ledgeetαἰ.2000:Martín&Bouy2002))."," Similarly, the ratio $L_X/L_{\rm bol}$ drops from a saturation value of about $10^{-2.5}$ in mid-M dwarfs to less than $10^{-4}$ for the few late-M dwarfs observed to date (e.g., \citealt{fgs+93,rbm+00,mb02}) )."
 With the exception of a recently-discovered M9 dwarf for which a ratio of about 0.1 was measured during a flare (Hambaryanefαἰ.2004).. the flares observed in other late-M dwarfs do not exceed the saturation level observed in the mid-M dwarfs (Stelzer2004).," With the exception of a recently-discovered M9 dwarf for which a ratio of about 0.1 was measured during a flare \citep{hss+04}, the flares observed in other late-M dwarfs do not exceed the saturation level observed in the mid-M dwarfs \citep{ste04}."
. Thus. in both Ho and X-rays the saturation effect appears to be real. with flares replacing persistent emission in late-M and L dwarfs.," Thus, in both $\alpha$ and X-rays the saturation effect appears to be real, with flares replacing persistent emission in late-M and L dwarfs."
 aappears to follow the general trend with upper limits on the X- and Ho emission that are about three orders of magnitude below the respective saturation values., appears to follow the general trend with upper limits on the X-ray and $\alpha$ emission that are about three orders of magnitude below the respective saturation values.
 On the other hand. the radio activity in late-M and L dwarfs. and in particular2M0036+18.. exhibits a completely different trend.," On the other hand, the radio activity in late-M and L dwarfs, and in particular, exhibits a completely different trend."
" The observed ratio Lr/Lpo, is about 105 for persistent emission from early and mid-M dwarfs.", The observed ratio $L_R/L_{\rm bol}$ is about $10^{-8}$ for persistent emission from early and mid-M dwarfs.
 However. the level of radio emission from late-M and L dwarfs exceeds this value by about an order of magnitude. while Hares are brighter by at least two orders of magnitude.," However, the level of radio emission from late-M and L dwarfs exceeds this value by about an order of magnitude, while flares are brighter by at least two orders of magnitude."
 iin particular has a ratio Le/Lpe of about 10797., in particular has a ratio $L_R/L_{\rm bol}$ of about $10^{-6.2}$.
 Thus. while a significantly smaller number of objects have been observed in the radio compared to Ha. single detection 1s brighter than the value observed in. mid-M dwarfs.," Thus, while a significantly smaller number of objects have been observed in the radio compared to $\alpha$, single detection is brighter than the value observed in mid-M dwarfs."
 Moreover. the fraction of detected objects is much higher. ~40%.," Moreover, the fraction of detected objects is much higher, $\sim
40\%$."
 This suggests that a saturation effect does not exists in the radio band., This suggests that a saturation effect does not exists in the radio band.
 Moreover. the radio luminosity does not seem to decrease with surface temperature.," Moreover, the radio luminosity does not seem to decrease with surface temperature."
 Another important trend. proposed by BO2. is that rapic rotators exhibit stronger radio activity.," Another important trend, proposed by B02, is that rapid rotators exhibit stronger radio activity."
 This is shown i Figure 15.., This is shown in Figure \ref{fig:activity_vel}.
 In Ho there is a clear rotation-activity relation i dwarfs earlier than M7. but this relation clearly breaks down imn late-M and L dwarfs.," In $\alpha$ there is a clear rotation-activity relation in dwarfs earlier than M7, but this relation clearly breaks down in late-M and L dwarfs."
 A similar relation exists in X-rays. but it again breaks down in late-M and L dwarfs.," A similar relation exists in X-rays, but it again breaks down in late-M and L dwarfs."
 In the radio band. οἱ the other hand. the rotation-activity relation appears to extend to late-M and L dwarfs.," In the radio band, on the other hand, the rotation-activity relation appears to extend to late-M and L dwarfs."
 The only objects for which the ratio of radio to bolometric luminosity is lower than the detected objects are those with vsiné<<10 km s., The only objects for which the ratio of radio to bolometric luminosity is lower than the detected objects are those with $v{\rm sin}i\lesssim 10$ km $^{-1}$.
 Thus. it appears that the efficiency of magnetic field generation and dissipation in fact does not decrease in late-M and L dwarfs. but the efficiency of chromospheric and coronal emission does.," Thus, it appears that the efficiency of magnetic field generation and dissipation in fact does not decrease in late-M and L dwarfs, but the efficiency of chromospheric and coronal emission does."
 Similarly. the rotation activity relation appears to still hold. but the activity 1s manifested in radio emission instead of Ha or X-rays.," Similarly, the rotation activity relation appears to still hold, but the activity is manifested in radio emission instead of $\alpha$ or X-rays."
 The detection of rare Ho. and X- flares does suggest that chromospheric and coronal heating may still occur. but this may require large reconnection events. possibly accompanied by large radio flares.," The detection of rare $\alpha$ and X-ray flares does suggest that chromospheric and coronal heating may still occur, but this may require large reconnection events, possibly accompanied by large radio flares."
 The simultaneous. multi-wavelength observations presented in this paper provide unparalleled insight into the magnetic field properties of an L dwarf.," The simultaneous, multi-wavelength observations presented in this paper provide unparalleled insight into the magnetic field properties of an L dwarf."
 Most importantly. these observations directly confirm that radio activity is more prevalent in late-M and L dwarfs compared to Ha and X-ray activity.," Most importantly, these observations directly confirm that radio activity is more prevalent in late-M and L dwarfs compared to $\alpha$ and X-ray activity."
 The detection rate in the radio may be as high as 40%. while in Hoy it ts at most a few percent.," The detection rate in the radio may be as high as $40\%$ , while in $\alpha$ it is at most a few percent."
 Only a few objects later than M7 have been observed in the X-rays. but persistent emission is clearly rare.," Only a few objects later than M7 have been observed in the X-rays, but persistent emission is clearly rare."
 A detailed analysis of the radio emission along with the lack of detectable X-ray and Ho emission give rise to the following, A detailed analysis of the radio emission along with the lack of detectable X-ray and $\alpha$ emission give rise to the following
to account for the cillerence of a static ancl dynamical atmosphere.,to account for the difference of a static and dynamical atmosphere.
 Because of the lack of dynamical mocel atmospheres. al ∖↓⊔⊔↓⊔≱∖∣∣⋡∢⋅∠⇂⋖⋅↿∢⋅↓⋅⊔∐⊔∢⊾∠⇂⋖⋅⊔↓↓≻↓↓⋅⊔∼⋜↧∐∙∖⇁⇂↓⋅∪⊔↓ . - ⋅ the quantities derived from photometry.," Because of the lack of dynamical model atmospheres, $a^{\rm (dyn)}$ must be determined empirically from the quantities derived from photometry."
 Ίσα. (3)), Eq. \ref{1.100}) )
" is transformed to the definition equation (1)) of QSAA if ür/ür=ϐ and als=Q aU=ᾖ, because these simplifications imply Orfol= Rand gAR)=GMR > ie. the UNA is valid."," is transformed to the definition equation \ref{1.102}) ) of QSAA if $\partial v/\partial r = 0$ and $a^{\rm (tang)}=0$, $a^{\rm (dyn)}=0$ , because these simplifications imply $\partial v/\partial t={\ddot R}$ and $g_{\rm s}(R)=G{\cal M}R^{-2}$ , i.e. the UAA is valid."
 As a dynamical equation. this simplified equation of motion is the basis of mass determination in BW analyses (e.g. Liu&Janes1990.. Cacciari.Clementini&Buser 1989.. ete).," As a dynamical equation, this simplified equation of motion is the basis of mass determination in BW analyses (e.g. \citealt{liuj1}, \citealt{cacc1}, etc)."
 If the sound: velocityis assumed as an upper limit for rp and ο. aENTE is some O.lms 7.," If the sound velocityis assumed as an upper limit for $v_{\theta}$ and $v_{\phi}$, $a^{\rm (tang)}$ is some $0.1\mbox{ms}^{-2}$ ."
 Ht is negligible in comparison with the other components of the acceleration at anv rA., It is negligible in comparison with the other components of the acceleration at any $r \la R$.
 Evpical values are ος 10ms.7 during the ↓≻⇂⇂↓≻⋜↧⇂↕⋖≱↓↥≼⇍∙∖⇁≼⇍⇂∢⊾∪⇂⋅⋜⋯∐∐⊳∖⇂⋜⊔⋅∖∖⋎↓↕⊲↓↓⋖⋅∙↙∣∣⇁≼⇍⋜⋯⋖⋅⇀∖≼∙⋖⋅∢⋅∠⇂↓∪∪⊔↓⊳∖ when the atmosphere is in the state of maximal compression., Typical values are $g_{\rm s} \la 10{\rm ms}^{-2}$ during the pulsation cycle of an RR star while $g_{\rm e}$ can exceed $100{\rm ms}^{-2}$ when the atmosphere is in the state of maximal compression.
 ‘Thus. αPus! aub be neglected and is a periodic function of / with zero points.," Thus, $a^{\rm (tang)}$ will be neglected and is a periodic function of $t$ with zero points."
 Neglecting the r-dependence. of the temperature. 7. assuming a constant density upper boundary condition pI()]. an integration of (4)) over the interval or.A] with the equation of state of a perfect gas gives the approximate density stratification of the model atmosphere: where ho(H.0)=pgGOEÉERATGIDCR.i are the reciprocal barometric scale height at /7. the universal eas constant. and average molecular mass. respectively.," Neglecting the $r$ -dependence of the temperature $T$, assuming a constant density upper boundary condition $\varrho[R(t)]$, an integration of \ref{2.100}) ) over the interval $[r,R]$ with the equation of state of a perfect gas gives the approximate density stratification of the model atmosphere: where $h_0(R,t)=\mu g_e(t)/{\cal R}T(R),{\cal R},\mu$ are the reciprocal barometric scale height at $R$, the universal gas constant, and average molecular mass, respectively."
 Essentially we use the static model atmospheres to measure —o‘OpOr). the atmospheric pulsation is driven by g(r./)αενα] . ancl the thermal processes are represented by the variable hy.," Essentially we use the static model atmospheres to measure $-\varrho^{-1}(\partial p/\partial r)$, the atmospheric pulsation is driven by $g(r,t)-a^{\rm (dyn)}$ , and the thermal processes are represented by the variable $h_0$."
 This concition formulates that the dynamical excess of acceleration in the upper photosphere is smaller than the Ag error o£ g and the same error is assumed for οἱ|refOr.," This condition formulates that the dynamical excess of acceleration in the upper photosphere is smaller than the $\Delta g$ error of $g$ and the same error is assumed for $\partial v/\partial t+v\partial v/\partial
r$."
 The combination of (5)) ancl (3)) allows to estimate the constant and O(d) terms of a7vul7., The combination of \ref{2.120}) ) and \ref{1.100}) ) allows to estimate the constant and ${\rm O}(d)$ terms of $a^{\rm (dyn)}$.
 However. the satisfaction of Condition IHE can be checked afterwards when the terms with other powers of @ were determined.," However, the satisfaction of Condition II can be checked afterwards when the terms with other powers of $d$ were determined."
 “Phe quotient qir.0)-TEEg characterizes the degree of excellence of the QSAA.," The quotient $q(r,t)=a^{\rm (dyn)}/g$ characterizes the degree of excellence of the QSAA."
 Phe QSAA is exact from. the hvdrodynamic point of view if q= 0., The QSAA is exact from the hydrodynamic point of view if $q=0$ .
 The perfect. spherical svnunetry means that in Euler picture (Pringle&Wing2007) we have to use (3))-(7)) with ai(Gs41)=0 and the mass conservation Law must be taken into account.," The perfect spherical symmetry means that in Euler picture \citep{prin1} we have to use \ref{1.100}) \ref{2.110}) ) with $a^{\rm (tang)}(r,t)=0$ and the mass conservation law must be taken into account."
" With assumption (7)) the analvtic solution of (8)) is where a,=ChthyOL). ay=R|Ca,3h,1). ay=.2Caoh,1 7.οξέα1phy. CH1."," With assumption \ref{2.110}) ) the analytic solution of \ref{1.101}) ) is where $a_1=Ch_0^{-1}(\partial h_0/\partial t)$, $a_0={\dot
 R}+Ca_1(R-3h_0^{-1})$, $a_{-1}=2Ca_0h_0^{-1}$, $a_{-2}=Ca_{-1}h_0^{-1}$, $C=1$."
 C was introduced: to incorporate UA by talking aU(p)20 and €=0.," $C$ was introduced to incorporate UAA by taking $a^{\rm (dyn)}(R,t)=0$ and $C=0$."
 Themain term in p ds @ and the convergence of (9)) is excellent because rhx1., Themain term in $v$ is ${\dot R}$ and the convergence of \ref{2.302}) ) is excellent because $rh_0\gg 1$.
 After dillerentiations of e. hy. and ?. taking r=2. the Following form is convenient for a numerical solution: where Md.0)—galt)(Ov1)(ο)aUauem C.," After differentiations of $v$, $h_0$ , and $\vartheta$, taking $r=R$, the following form is convenient for a numerical solution: where ${\cal M}(d,t)=[g_{\rm e}(t) -(\partial v/\partial t)
-v(\partial v/\partial r) -a^{\rm (dyn)}(R,t)]R^2/ G$ ."
 |t can. be solved i£ there exist. two or more f intervals satislving Condition Land a7(24) is identical for them.," It can be solved if there exist two or more $t$ intervals satisfying Condition I and $a^{\rm (dyn)}(R,t)$ is identical for them."
 The assumed differences aD!(2.44)aleeryoxO and the satisfaction. of Condition LL must be verified afterwards.," The assumed differences $a^{\rm (dyn)}(R,t_1)-a^{\rm (dyn)}(R,t_2)\approx 0$ and the satisfaction of Condition II must be verified afterwards."
 Iq. (10)), Eq. \ref{107a}) )
 must be solved. for different phase pairs (1.2) and the values d and © must be omitted from the final averaging if αιοgn.yoo) ds violated.," must be solved for different phase pairs $(\varphi_1,\varphi_2)$ and the values $d$ and ${\cal M}$ must be omitted from the final averaging if $a^{\rm (dyn)}(R,\varphi_1)\approx a^{\rm (dyn)}(R,\varphi_2)$ is violated."
 'l'o account lor the dillerence of the local and. elfective temperature. the boundary temperature Z(/2)=2τς of a grav atmospheric model at 7=0 (Alihalas1978). will be used in fo(feof).," To account for the difference of the local and effective temperature, the boundary temperature $T(R)=2^{-0.25}T_{\rm e}$ of a gray atmospheric model at $\tau=0$ \citep{miha1} will be used in $h_0(R,t)$."
 Hydrogen and helium are mainlv. neutral at the boundary temperature of II stars., Hydrogen and helium are mainly neutral at the boundary temperature of RR stars.
 Consequently. assuming a typical chemical composition means that so—1.8 is the appropriate choice in Pou(4.," Consequently, assuming a typical chemical composition means that $\mu=1.3$ is the appropriate choice in $h_0(R,t)$ ."
 1). Awcza(2002)derived good quality light curves of. the lutab star SU Dra by homogenizing 45 vearsof CDVCURL: observations (see his Table 7)., \citet{barc0} derived good quality light curves of the RRab star SU Dra by homogenizing 45 yearsof $UBV(RI)_C$ observations (see his Table 7).
 The method described in the previous sectionwill now be applied., The method described in the previous sectionwill now be applied.
 Instead. of smoothing the light curves. binned data will beused to keep the results as close as possible to observed quantities.," Instead of smoothing the light curves, binned data will beused to keep the results as close as possible to observed quantities."
 Lhe30 colour index pairs containing at least one C were obtained. [rom CLT=f2Bq|pVULU.ReLB.MIB de}.," The30 colour index pairs containing at least one $U$ were obtained from ${\rm CI}_i\vert_{i=1}^{i=10}=
\{U-2B+V,U-V,U-R_C,U-I_C,B-V,B-R_C,B-I_C,V-R_C,V-I_C,R_C-I_C\}$ ."
 Taking Al] 1.60. νιVW)=0.015 (Liu&Janes1990). the basic quantities," Taking $[M]=-1.60$ , $E(B-V)=0.015$ \citep{liuj1} the basic quantities"
 has0.1iubeen The iron abundance of the ICAL is not determined by cluster simmiations. and so has to be assumed. based on data frou observed clusters.," 0.1in The iron abundance of the ICM is not determined by cluster simulations, and so has to be assumed based on data from observed clusters."
" ναι the small field of view of the NRS on-boardAstro-E2, we are probing regions fairly close to the clusters core C50.2royy)."," Given the small field of view of the XRS on-board, we are probing regions fairly close to the cluster's core $\lta 0.2\;
r_{200}$ )."
" The iron abundance in cluster cores is typically higher than the average over the whole cluster. due to the frequent preseuce of ceutral metal abundance gradients im cold core (""cooliug How’) clusters (μον et al."," The iron abundance in cluster cores is typically higher than the average over the whole cluster, due to the frequent presence of central metal abundance gradients in cold core (“cooling flow”) clusters (Ulmer et al."
 1987: White ct al., 1987; White et al.
 1991). which constitute the majority of clusters.," 1994), which constitute the majority of clusters."
 Asstuine an average central mon abundance of 0.27 solar for non- core clusters aud 0.17 for cold core svsteiis (DeCiraucdi et al., Assuming an average central iron abundance of $0.27$ solar for non-cold core clusters and $0.47$ for cold core systems (DeGrandi et al.
 2001). and recognizing that το90% of ai X-ray flux-Iuuited sample of clusters have cold cores (Edge. Stewart Fabian 1992). we obtain an average central iron abunudauce of ~0.L relative to the solar plhotospleric values of Auclers Grevesse (1989).," 2004), and recognizing that $70-90\%$ of an X-ray flux-limited sample of clusters have cold cores (Edge, Stewart Fabian 1992), we obtain an average central iron abundance of $\sim$ 0.4 relative to the solar photospheric values of Anders Grevesse (1989)."
 This is the value asstuned throughout this work., This is the value assumed throughout this work.
 One of the main obstacles in measuring velocities of the intracluster eas with curent spectrometers ds the inability to accurately calibrate the temporal and spatial fluctuations of instrumental euin. the conversion of pulse-height iuto photon οποιον (Dupke Breeman 200110).," One of the main obstacles in measuring velocities of the intracluster gas with current spectrometers is the inability to accurately calibrate the temporal and spatial fluctuations of instrumental gain, the conversion of pulse-height into photon energy (Dupke Bregman 2001b)."
 For the case of CCDs the iutrachip (positional dependent) eain changes are related to the charge trausfer inefficieucy. which evolves with time.," For the case of CCDs the intrachip (positional dependent) gain changes are related to the charge transfer inefficiency, which evolves with time."
 Often these fiuctuations are ou the same order as the velocities one is tryvime to measure., Often these fluctuations are on the same order as the velocities one is trying to measure.
 Although this is not an issue for the NRS calorimeter. sanall diifts in the temperature of the detector heat siuk can cause elobal gain variations.," Although this is not an issue for the XRS calorimeter, small drifts in the temperature of the detector heat sink can cause global gain variations."
 The absolute precision of the gain variations in the NRS is expected to be ~12 eV (50100kms 4) at 6 keV and will have to be inserted iuto the uncertainties of the velocity mecasurcments describe here (Figueroa. personal conumnuuncatiou. 2001).," The absolute precision of the gain variations in the XRS is expected to be $\sims 1-2$ eV $50-100 \kms$ ) at 6 keV and will have to be inserted into the uncertainties of the velocity measurements described here (Figueroa, personal communication, 2004)."
 These uucertainties were not incorporated iun our evaluations of velocity gradients. but they do not affect the overal results derived in this work.," These uncertainties were not incorporated in our evaluations of velocity gradients, but they do not affect the overall results derived in this work."
 Similarly. for the purpose of this paper we do uot fake a background spectrum.," Similarly, for the purpose of this paper we do not fake a background spectrum."
 The main source of background for the NRS is due to energv deposition in the detector by euergetie protous., The main source of background for the XRS is due to energy deposition in the detector by energetic protons.
 However. eiven the small detector size and the distribution of the proton cherey spectrum. its contribution iu the frequencies of interest (<< 10 keV) are expected to be negligible (<< 5%)) for our mock observations.," However, given the small detector size and the distribution of the proton energy spectrum, its contribution in the frequencies of interest $<$ 10 keV) are expected to be negligible $<$ ) for our mock observations."
 Finally. iu this paper we assume that the intracluster plasma is optically thin.," Finally, in this paper we assume that the intracluster plasma is optically thin."
 Calfanov. Suuvaev Cliurazov (1987) have poiuted out that this nav not always be true for the Fe Ίνα liue.," Gil'fanov, Sunyaev Churazov (1987) have pointed out that this may not always be true for the Fe $\alpha$ line."
 Resonant scattering (the absorption and inunuediate re-cuuission of an Fe Ίνα line photon by au Fe jon) can be significant iu the cores of clusters., Resonant scattering (the absorption and immediate re-emission of an Fe $\alpha$ line photon by an Fe ion) can be significant in the cores of clusters.
 If so. the lueσα cuuission from the cluster core is reduced because of the scattering of photous out of the field of view.," If so, the line emission from the cluster core is reduced because of the scattering of photons out of the field of view."
 Early evidence for resonant scattering (Moleudi et al., Early evidence for resonant scattering (Molendi et al.
 1998) secu in the Perseus cluster has proven ambiguous. allowing other possible interpretations such as an overabundance of Ni due the ceutral SN Ia ejecta dominance (Dupke Arnaud 20013.," 1998) seen in the Perseus cluster has proven ambiguous, allowing other possible interpretations such as an overabundance of Ni due the central SN Ia ejecta dominance (Dupke Arnaud 2001)."
NALAF observations of the Perseus cluster have not exhibited resonant scattering in the Fe Ίνα linc. and the absence of observed scattering has been used as evidence for the presence of turbulent (non-therimial) velocities iu clusters cores (Castaldello Moleucdi 20041. Churazov et al.," observations of the Perseus cluster have not exhibited resonant scattering in the Fe $\alpha$ line, and the absence of observed scattering has been used as evidence for the presence of turbulent (non-thermal) velocities in clusters' cores (Gastaldello Molendi 2004, Churazov et al."
 2001)., 2004).
 -ü.2iu To quantify the potential inpact of resonant scattering. we compute the optical depth for the central poiutiug of each cluster projection used in this study (see Appendix A or details).," -0.2in To quantify the potential impact of resonant scattering, we compute the optical depth for the central pointing of each cluster projection used in this study (see Appendix \ref{sec:appa} for details)."
 To be conservative. we calculate the optical depth through the eutire cluster.," To be conservative, we calculate the optical depth through the entire cluster."
 The results are shown in Fieure 3.., The results are shown in Figure \ref{fig:tau}.
 When the non-thermal velocities seen in the cnussion line are iucluded. the average optical depth for 1 FoNNY Ko line in the FOV of the central pointing is 0.59. with standard deviation 0.3L.," When the non-thermal velocities seen in the emission line are included, the average optical depth for the FeXXV $\alpha$ line in the FOV of the central pointing is $0.59$, with standard deviation $0.34$."
 The results of Churazov ct al. (, The results of Churazov et al. (
2001) and Calfanov et al. (,2004) and Gil'fanov et al. (
1987) show iat for 7~0.6. the flux and the linewidth near the cluster core will show approximately a reduction.,"1987) show that for $\tau \sim 0.6$, the flux and the linewidth near the cluster core will show approximately a reduction."
 Additionally. it is possible that some deformation of the ine profile may occur (Cülfauov et al.," Additionally, it is possible that some deformation of the line profile may occur (Gil'fanov et al."
 1987)., 1987).
 None of τοσο effects would introduce errors in the Lue centroids we recover., None of these effects would introduce errors in the line centroids we recover.
 The reduction in linewidth would affect our ine broadening study. but. as our results will show. a reduction is not prolibitive.," The reduction in linewidth would affect our line broadening study, but, as our results will show, a reduction is not prohibitive."
 We first investigate the LOS velocity structure of the siuulated clusters by creating one set of four mock exposures filed around the most bound position of the relevant cluster. as shown by boxes l-1 iu Figue |," We first investigate the LOS velocity structure of the simulated clusters by creating one set of four mock exposures tiled around the most bound position of the relevant cluster, as shown by boxes 1-4 in Figure \ref{fig:tiles}."
 The mock exposures result iu four separate Poissou-uoise-added spectra., The mock exposures result in four separate Poisson-noise-added spectra.
" We rebin these spectra to ensure at least 15 counts iu cach bin using the routine from the FTOOLS We fit the grouped spectra to find the position of the centroid velocity ce; of the Fe ἵνα line in nuage ἐ,", We rebin these spectra to ensure at least 15 counts in each bin using the routine from the FTOOLS We fit the grouped spectra to find the position of the centroid velocity $v_i$ of the Fe $\alpha$ line in image $i$.
 The spectra ave fit in NSPEC to a BAPEC model., The spectra are fit in XSPEC to a BAPEC model.
 BAPEC is a thermally broadened APEC that allows for additional Gaussian velocity broadeniug uuder a parameter σι., BAPEC is a thermally broadened APEC that allows for additional Gaussian velocity broadening under a parameter $\sigma_v$.
 The resulting centroid positions are differenced to define velocity gradients οσοοἱ among the six distinct pairs of the four pointings., The resulting centroid positions are differenced to define velocity gradients $v_{ij} \se |v_i -v_j|$ among the six distinct pairs of the four pointings.
 Thus. from our 1.836 independent cluster projections we eeuecrate 7.311 spectra which are used to eive 11.016 velocity differences.," Thus, from our 1,836 independent cluster projections we generate 7,344 spectra which are used to give 11,016 velocity differences."
 Au advantage of working with simulated clusters is our complete kuowledee of the underlvine velocity structure of the eas., An advantage of working with simulated clusters is our complete knowledge of the underlying velocity structure of the gas.
 In an isothermal cluster. the redshitt of the line ceuter could be recovered from the deusity-squared. weighted velocity (heuceforth called the ENGweiehted velocity) of the emitting eas.," In an isothermal cluster, the redshift of the line center could be recovered from the density-squared weighted velocity (henceforth called the EM-weighted velocity) of the emitting gas."
 Our simmlated. clusters are uot isothermal. but it lappeus that the Fe Ίνα line flux is uot stronely temperature dependent within the teiiperature range expected for hot ICAL eas.," Our simulated clusters are not isothermal, but it happens that the Fe $\alpha$ line flux is not strongly temperature dependent within the temperature range expected for hot ICM gas."
 The peal flux varies by only a factor of 2 over a temperature, The peak flux varies by only a factor of 2 over a temperature
source couuts iu the soft baud (0.83.5 keV) obtained by EExteuded Mecimu Seusitivitv Survey (EMSS: Cüola 11990) is about 223 times smaller than that in the hard baud (210 keV) obtained by the ffluctuation analvsis (Butcher 11997) when we assume a power-law photon iudex of 1.7.,source counts in the soft band (0.3–3.5 keV) obtained by Extended Medium Sensitivity Survey (EMSS; Gioia 1990) is about 2–3 times smaller than that in the hard band (2–10 keV) obtained by the fluctuation analysis (Butcher 1997) when we assume a power-law photon index of 1.7.
" The ssatellite (Tanaka. Πποπο, Πο 1991). launched in 1993 February. was expected to change this situation."," The satellite (Tanaka, Inoue, Holt 1994), launched in 1993 February, was expected to change this situation."
 It is the first iniaeius satellite capable of study of the N-rav ud above 2 keV with a sensitivity up to several LOtlore ((210 keV) and covers the wide euergv. band from 0.5 to LO keV. which allows us o directly compare results of the cucrey bands below aud above 2 keV with sinele detectors. hence acconipauied with much less uncertainties than previous studies.," It is the first imaging satellite capable of study of the X-ray band above 2 keV with a sensitivity up to several $ 10^{-14}$ (2–10 keV) and covers the wide energy band from 0.5 to 10 keV, which allows us to directly compare results of the energy bands below and above 2 keV with single detectors, hence accompanied with much less uncertainties than previous studies."
 Dv taking these advantages. several X-ray surveys lave been performed with tto reveal the nature of hard X-ray populations: the LLaree Sky Survey (LSS: Ueda 11998). the DDeep Sky Survey (DSS: Ogasaka 11998: Ixhisaki 1999 for the Lockinan Hole). the MMedciuni-Seusitivity Survey CAMSS or the CUS catalog project: Ueda 11997. Takahashi 1998. Ueda 1999b: see also Cagnoni. Della Coca. Maccacaro 1998 and Della Ceca 1999). a survev of ldeep fields (Ceorgantopoulos 1997: Bovle 1998). aud so on.," By taking these advantages, several X-ray surveys have been performed with to reveal the nature of hard X-ray populations: the Large Sky Survey (LSS; Ueda 1998), the Deep Sky Survey (DSS; Ogasaka 1998; Ishisaki 1999 for the Lockman Hole), the Medium-Sensitivity Survey (AMSS or the GIS catalog project: Ueda 1997, Takahashi 1998, Ueda 1999b; see also Cagnoni, Della Ceca, Maccacaro 1998 and Della Ceca 1999), a survey of deep fields (Georgantopoulos 1997; Boyle 1998), and so on."
 The sensitivity lnüts aud survey area are summarized in Table 1., The sensitivity limits and survey area are summarized in Table 1.
 Iu this paper. we preseut main results of the ssurvevs. focusing on the LSS 2). the Lockman Hole deep survey 3). and the AMSS Εν," In this paper, we present main results of the surveys, focusing on the LSS 2), the Lockman Hole deep survey 3), and the AMSS 4)."
 5. we stummarize these results and discuss their inplicatious for the origin of the CNB.," In 5, we summarize these results and discuss their implications for the origin of the CXB."
with increasing wavelength.,with increasing wavelength.
 In fact the modulation is clearly recognizable only in DB. V. while in 1t E the magnitude is practically constant.," In fact the modulation is clearly recognizable only in B V, while in R I the magnitude is practically constant."
 This fact also follows from the PS of each filters data set (Fig. 3))., This fact also follows from the PS of each filter's data set (Fig. \ref{BVRIps}) ).
 Since D V are varving in phase. the modulation of the B.N colour has only a small amplitude.," Since B V are varying in phase, the modulation of the $-$ V colour has only a small amplitude."
 In Figure 6 one can see that the folded D LC. the filter in which the signal is strongest. is not a perfect sinusoid.," In Figure \ref{harmony} one can see that the folded B LC, the filter in which the signal is strongest, is not a perfect sinusoid."
 “Phis xutern of a Hlat-top LC is also apparent in the folded V. and Clear LCs. dthough not as vivid as in D. In order to verily he significance of the distorted. sinusoid. we calculated the yhase of each filkter's second harmonic.," This pattern of a flat-top LC is also apparent in the folded V and Clear LCs, although not as vivid as in B. In order to verify the significance of the distorted sinusoid, we calculated the phase of each filter's second harmonic."
 Γον all fell within 1.13 of the evele of one another., They all fell within 0.13 of the cycle of one another.
 For random deviation from a sine function. the probability of obtaining such a crowding of he phase values isA.," For random deviation from a sine function, the probability of obtaining such a crowding of the phase values is."
. Fhus the non harmonic structure of the periodic LC seems to be systematic and significant., Thus the non harmonic structure of the periodic LC seems to be systematic and significant.
 We did not perform any spectral measurements on the WO spectra and. restricted: ourselves only to line identification and general properties of the spectrum due to the poor resolution. S/N ratio and. inability to resolve 10 from B417 and the nebulositv.," We did not perform any spectral measurements on the WO spectra and restricted ourselves only to line identification and general properties of the spectrum due to the poor resolution, S/N ratio and inability to resolve B416 from B417 and the nebulosity."
 The APO spectra on the other hand. were clean of large light contributions from external sources (except for several faint stars around. D416) and we were able to separate D416 from the surrounding nebula., The APO spectra on the other hand were clean of large light contributions from external sources (except for several faint stars around B416) and we were able to separate B416 from the surrounding nebula.
 The three sets of APO spectra were analysecl using the IRAP task to measure line and continuum properties., The three sets of APO spectra were analysed using the IRAF task to measure line and continuum properties.
 The relatively. high. resolution of these spectra. enabled: us o deblend sets of nearby lines and to detect. variations in ine profiles., The relatively high resolution of these spectra enabled us to deblend sets of nearby lines and to detect variations in line profiles.
 Table 2. lists the mean measured. properties of some of the major identified lines in the spectra. which are oesented in Figures 7 S..," Table \ref{lines} lists the mean measured properties of some of the major identified lines in the spectra, which are presented in Figures \ref{APOred} \ref{APOblue}."
 The line properties such as Dux. EW ete.," The line properties such as flux, EW etc."
 were averaged among the three spectra ancl the errors in the measured »xoperties represent the SLD of the mean., were averaged among the three spectra and the errors in the measured properties represent the STD of the mean.
 In this sense an error larger than about (a rough estimation of the uncertainty in flux calibration) in a measured property in one of the major strong lines of the star indicates a possible variation between the three epochs., In this sense an error larger than about (a rough estimation of the uncertainty in flux calibration) in a measured property in one of the major strong lines of the star indicates a possible variation between the three epochs.
 In the nebular lines the errors are much larger cue to lower fluxes and cillerences in the subtraction parameters used to extract the lince-ol-sight nebula between the three epochs (see 2.2))., In the nebular lines the errors are much larger due to lower fluxes and differences in the subtraction parameters used to extract the line-of-sight nebula between the three epochs (see \ref{2.2}) ).
 The spectrum of B46 in general resembles that of a blue hvpergiant. or more specifically an LBV in quiescence. as suggested in Massey et al. (," The spectrum of B416 in general resembles that of a blue hypergiant, or more specifically an LBV in quiescence, as suggested in Massey et al. ("
1996).,1996).
 He shows mainly Balmer. He1 Fe and Fe n] emission (B416 is also. identilied in Corral (1996) as object S145. where it is classified. as a compact LE region).," It shows mainly Balmer, He, Fe and [Fe ] emission (B416 is also identified in Corral (1996) as object S145 where it is classified as a compact H region)."
 The main spectral feature is the strong Ilo. emission. which can be evaluated. from Table as having the luminosity of —1700 L.; (using a distance of SOO0kpe to M33). and an [EN of -l0G+2A. which is consistent with the -109.1 vyalue measured for this star bv Calzetti et al. (," The main spectral feature is the strong $\alpha$ emission, which can be evaluated from Table \ref{lines} as having the luminosity of $\sim$ 1700 $_{\sun}$ (using a distance of 800kpc to M33) and an EW of $\pm$ 2, which is consistent with the -109.1 value measured for this star by Calzetti et al. ("
1995).,1995).
 ‘Phe measured wavelengths of, The measured wavelengths of
the y ancl z axes of our coordinate svstem such (hat (he magnetic field lies in the (x.v) plane.,"the $y$ and $z$ axes of our coordinate system such that the magnetic field lies in the (x,y) plane."
 The input parameters to our model simulation describing the external radiation field are (he integral of ca(vr.Q) over all Lrequencies: the blackbody temperature 7. and the radial extent παω.," The input parameters to our model simulation describing the external radiation field are the integral of $u_{\rm ext} (\nu, r, \Omega)$ over all frequencies: the blackbody temperature $T$, and the radial extent $R_{\rm ext}$."
 We have used the Monte-Carlo code developed by Roustazadeh&Bollcher(2010)., We have used the Monte-Carlo code developed by \cite{rb10}.
. This code evaluates 25 absorption ancl pair production using the full analytical solution {ο the pair production spectrum of Bolicher&Schlickeiser(1997) under the assumption tha the produced electron and positron travel along (he direction of propagation of (he incoming -rav., This code evaluates $\gamma\gamma$ absorption and pair production using the full analytical solution to the pair production spectrum of \cite{bs97} under the assumption that the produced electron and positron travel along the direction of propagation of the incoming $\gamma$ -ray.
 The trajectories of the particles are Followed in full 2-D geometry., The trajectories of the particles are followed in full 3-D geometry.
 Compton scattering is evaluated. using the head-on approximation. assuming that the scattered photon travels along the direction of motion of the electron/positron at the time of scattering.," Compton scattering is evaluated using the head-on approximation, assuming that the scattered photon travels along the direction of motion of the electron/positron at the time of scattering."
 While the Compton enerev loss to the elecron is properly accounted Dor. we neglect the recoil effect on the travel direction of the electron.," While the Compton energy loss to the elecron is properly accounted for, we neglect the recoil effect on the travel direction of the electron."
 In order to improve the statistics of the otherwise very few highest-energv photons. we introduce a statistical weight. «e. inversely proportional to the square of: (he photon energy. w=21/e.," In order to improve the statistics of the otherwise very few highest-energy photons, we introduce a statistical weight, $w$ , inversely proportional to the square of the photon energy, $w = 1/\epsilon^2$."
 -Where e=Lb. σα., Where $\epsilon = \frac{E_{\gamma}}{m_e c^2}$ .
 To save CPU (ime. we precaleulate tables for the absorption opacity A... Compton scaltering length: Ay. aud Compton cross section for each photon energy. electron energy and interaction angle before the start of the actual simulation.," To save CPU time, we precalculate tables for the absorption opacity $\kappa_{\gamma\gamma}$, Compton scattering length $\lambda_{\rm IC}$, and Compton cross section for each photon energy, electron energy and interaction angle before the start of the actual simulation."
 The simulation produces a photon event file. logging the energy. statistical weight. aud travel direction of every photon that escapes the region bounded by. Rog.," The simulation produces a photon event file, logging the energy, statistical weight, and travel direction of every photon that escapes the region bounded by $R_{\rm ext}$."
 A separate routine is (hen used (o extract augle-dependent. photon spectra with. arbitrary angular and energy. binning [rom the log files., A separate post-processing routine is then used to extract angle-dependent photon spectra with arbitrary angular and energy binning from the log files.
 For comparison with our previous study on mono-energetic radiation fields. we conduct a similar parameter study as presented in (Houstazadeh&Boltcher2010).. investigating the effects of parameter variations on the resulting angle-dependent photon spectra.," For comparison with our previous study on mono-energetic radiation fields, we conduct a similar parameter study as presented in \citep{rb10}, investigating the effects of parameter variations on the resulting angle-dependent photon spectra."
" Standarcl parameters for most simulations in our parameter study are: a magnetic field of B=1 mG.oriented at an angle of 05=5° with respect to the X axis (D,=I mG. D,=0.1 mG): anexternal radiationenereyv density of te.=10? erg . extended over a region of radius"," Standard parameters for most simulations in our parameter study are: a magnetic field of $B = 1$ mG,oriented at an angle of $\theta_B = 5^o$ with respect to the X axis $B_x = 1$ mG, $B_y = 0.1$ mG); anexternal radiationenergy density of $u_{\rm ext} = 10^{-5}$ erg $^{-3}$ extended over a region of radius"
results noted above. E assumed the bulk of the CV population to have the kinematics of Galactic INQ giants as tabulated in Mibalas&Binney(1981). (p. 123). with ίσισιoy)=(31.21.16) kms,"results noted above, I assumed the bulk of the CV population to have the kinematics of Galactic K0 giants as tabulated in \citet{mihalbinn} (p. 423), with $(\sigma_U,\sigma_V,\sigma_W) = (31, 21, 16)$ km $^{-1}$."
 HI added to this a high-velocity tail with a normalization of 0.05 times the bulk population. with (100. 75.50) kms 4. sunilar to the subcdwarks.," I added to this a high-velocity tail with a normalization of 0.05 times the bulk population, with (100, 75, 50) km $^{-1}$, similar to the subdwarfs."
 Finally. E added a lower-velocity core. with 0.2 times the normalization. with (21. 13. 10) kins !. similar to FO cdwarfs.," Finally, I added a lower-velocity core, with 0.2 times the normalization, with (24, 13, 10) km $^{-1}$, similar to F0 dwarfs."
 This composite probability density was evaluated over a range of hypothetical true parallaxes., This composite probability density was evaluated over a range of hypothetical true parallaxes.
 The proper motion was adjusted to the local standard of rest at each parallax before the probability deusity was evaluated., The proper motion was adjusted to the local standard of rest at each parallax before the probability density was evaluated.
 The uncertainty in the proper motion determination (equ., The uncertainty in the proper motion determination (eqn.
 27 in Suith1987b)) was ignored., 27 in \citealt{smith87b}) ) was ignored.
 The apparent maguitudes of mauy of these systems cau also be used to coustrain the distance., The apparent magnitudes of many of these systems can also be used to constrain the distance.
" Warner(1987). (also. Warner 1995)) showed th: the absolute magnitudes of cdwarl novae at uaxinui light. corrected for iuclinatiou. are strongly correlated with P,4,."," \citet{warn87} (also \citealt{warn}) ) showed that the absolute magnitudes of dwarf novae at maximum light, corrected for inclination, are strongly correlated with $P_{\rm orb}$."
" For Poy,«2 hr. the uaxinuim iagnitude generally occurs iu superoutburst. which is about a magnitude brighter than iormal outburst: I assumed this to be the case."," For $P_{\rm orb} < 2$ hr, the maximum magnitude generally occurs in superoutburst, which is about a magnitude brighter than normal outburst; I assumed this to be the case."
 The optical colors of dwarf uovae in outburst are alrly close to zero. so the distinction between ips and ny is ignored.," The optical colors of dwarf novae in outburst are fairly close to zero, so the distinction between $m_{\rm
pg}$ and $m_{\rm V}$ is ignored."
 The orbital inclinatious lor nost of the program objects are uncertain. so the inclination correction is not known.," The orbital inclinations for most of the program objects are uncertain, so the inclination correction is not known."
 To account or this and unexpected scatter in the relation. and to avokl casstunine what we are trying to »ove.the —jagnitudes were assumed to follow a very broad CGatussian.," To account for this and unexpected scatter in the relation, and to avoid `assuming what we are trying to prove', the magnitudes were assumed to follow a very broad Gaussian."
 Sometimes other constraints on the dist:ice were available (e.g.. from detectious of secoudary stars). and again relatively broad o»obability distributions were assumed to avoid steering the estimate too much.," Sometimes other constraints on the distance were available (e.g., from detections of secondary stars), and again relatively broad probability distributions were assumed to avoid steering the estimate too much."
 The notes on individual stars detail the adopted absolute magnitude coustraits., The notes on individual stars detail the adopted absolute magnitude constraints.
 The parallax. proper motion. and maenitudee information was combined as follows.," The parallax, proper motion, and magnitude information was combined as follows."
 Bayes’ tlieoreiu states. in general teris where P? represents a probability. A the hypothesis. D tle data (in this case. the observed parallax). aud { the prior information about the problem (corstraints derived [rom the proper motion aud inagnitude. auc asstumptious such as the uormal distribution of errors).," Bayes' theorem states, in general terms where $P$ represents a probability, $H$ the hypothesis, $D$ the data (in this case, the observed parallax), and $I$ the prior information about the problem (constraints derived from the proper motion and magnitude, and assumptions such as the normal distribution of errors)."
 lu this case H is a hypothesized true parallax. e.g. VY Aqr las a true j»arallax of 8.2 mas’. aud we are asking; lor the likelihood that A is correct given 2 (the measured parallax auc its estimated uucertainty) aud J (the proper motion aud the assumptious about the space velocity. the apparent maguitucde aud the asstuuptious about the plausible range of absolute maguitudes. and the assumed normal distribution of the experimental error).," In this case $H$ is a hypothesized true parallax, e.g. `VY Aqr has a true parallax of 8.2 mas', and we are asking for the likelihood that $H$ is correct given $D$ (the measured parallax and its estimated uncertainty) and $I$ (the proper motion and the assumptions about the space velocity, the apparent magnitude and the assumptions about the plausible range of absolute magnitudes, and the assumed normal distribution of the experimental error)."
 The true parallax can be auy positive number. so we ruu the computatious for a range of parallaxes [rom uear 0 up to large values (i.e.. we vary H ). creating a coutinuous probability density.," The true parallax can be any positive number, so we run the computations for a range of parallaxes from near 0 up to large values (i.e., we vary $H$ ), creating a continuous probability density."
 This continuous probability density for PH|[DI) is exactly what we want: the relative likelihood of each parallax. all the information we have available. including our measurement.," This continuous probability density for $P(H|DI)$ is exactly what we want: the relative likelihood of each parallax, all the information we have available, including our measurement."
 The first factor on the right is the probability of obtailing our measured parallax. [or the given true parallax aud the prior information.," The first factor on the right is the probability of obtaining our measured parallax, for the given true parallax and the prior information."
 Ouce the true parallax has been fixed. tlie assumption of," Once the true parallax has been fixed, the assumption of"
excess luminosity of Lx=6.2x10 ss7! and a duration of 9000 s. the flare released about Ex=5.6107 erg only at soft X-ray energies.,"excess luminosity of $L_{\rm X}= 6.2 \times 10^{30}$ $^{-1}$ and a duration of 9000 s, the flare released about $E_{\rm X}= 5.6 \times 10^{34}$ erg only at soft X-ray energies."
 The flare emission is dominated by extremely hot plasma and even the average temperature is found to be around 45 MK., The flare emission is dominated by extremely hot plasma and even the average temperature is found to be around 45 MK.
 In the PN spectrum the prominent emission line complex from He-like is clearly visible at 6.7 keV. consistent with a hot plasma component at temperatures of around 60—770 MK as deduced from the spectral nodel.," In the PN spectrum the prominent emission line complex from He-like is clearly visible at 6.7 keV, consistent with a hot plasma component at temperatures of around 70 MK as deduced from the spectral model."
 Weak excess emission is present at the low energy tail of this feature. that could be due to fluorescent emission from neutral iron from the KKe line at 6.4 keV. However. formally the line is not detected at the contribution to the X-ray flux is with an equivalent width of x100 eV at best marginal.," Weak excess emission is present at the low energy tail of this feature, that could be due to fluorescent emission from neutral iron from the $\alpha$ line at 6.4 keV. However, formally the line is not detected at the contribution to the X-ray flux is with an equivalent width of $\lesssim 100$ eV at best marginal."
 No significant changes in the absorption component are detected: in contrast. the metallicity of the flare plasma is found to be significantly higher than those of the quasi-quiescent phase.," No significant changes in the absorption component are detected; in contrast, the metallicity of the flare plasma is found to be significantly higher than those of the quasi-quiescent phase."
 While the absolute scale is rather poorly constrained. we find a distinct metallicity trend. most noticeable when further separating the flare into individual time intervals as discussed in refsf..," While the absolute scale is rather poorly constrained, we find a distinct metallicity trend, most noticeable when further separating the flare into individual time intervals as discussed in \\ref{sf}."
 To study the evolution of the flare and its underlying structure in greater detail. we perform a time resolved analysis for each of the six time segments as shown in the upper panel of refflp..," To study the evolution of the flare and its underlying structure in greater detail, we perform a time resolved analysis for each of the six time segments as shown in the upper panel of \\ref{flp}."
 To derive the properties of the flare plasma (1-6). we added plasma components to the quasi-quiescent (QQ) model derived for the PN data.," To derive the properties of the flare plasma (1-6), we added plasma components to the quasi-quiescent (QQ) model derived for the PN data."
 Here we use models with one or two temperature components and with or without metallicity as a free parameter to describe the flare plasma., Here we use models with one or two temperature components and with or without metallicity as a free parameter to describe the flare plasma.
 The models lead to overall consistent results. see for example the time-evolution of the plasma properties presented in the middle panel of refflp for models with solar metallicity.," The models lead to overall consistent results, see for example the time-evolution of the plasma properties presented in the middle panel of \\ref{flp} for models with solar metallicity."
 The temperature given for the 2-T model is the emission measure weighted temperature., The temperature given for the 2-T model is the emission measure weighted temperature.
 The flare reaches around maximum an X-ray ∣ ⇂↴∣⋪≣∶↔↾∣⇈∏⊖⋋⋋∪↑↴⋅≏↧⇂↴⋂∐↾∣⋂∶↔⊺∠≻∖⇁∶∍∟⊰⊜∣⋪∶↔⊺⋋⋋≓⋅≏↧⋯↿⋅≏⋯⋅≏↧∖⇁⊜∣⋪∐⊱⊺⊜ plasma temperature of around 80 MK. shortly after the impulsive heating during the rise phase (bin 1). even roughly 100 MK plasma is present.," The flare reaches around maximum an X-ray brightness of about $\log L_{\rm X}=31.5$ $^{-1}$ and an average plasma temperature of around 80 MK, shortly after the impulsive heating during the rise phase (bin 1), even roughly 100 MK plasma is present."
 The maximum temperature leads the maximum emission measure and thus in general the flare shows the temporal evolution of a typical coronal flare. Le. a clock-wise turn in the temperature vs. density plane (seee.g.?)..," The maximum temperature leads the maximum emission measure and thus in general the flare shows the temporal evolution of a typical coronal flare, i.e. a clock-wise turn in the temperature vs. density plane \citep[see e.g.][]{gue04}."
 Here the square-root of the emission measure is used as a density proxy. thus a fixed spatial extend of the X- emission region is assumed.," Here the square-root of the emission measure is used as a density proxy, thus a fixed spatial extend of the X-ray emission region is assumed."
 While the initial decay ts apparently steep. but only poorly constrained (bins 2-3).4 the flare evolution trajectory shows some flattening (bins 3-6) of the decay path.," While the initial decay is apparently steep, but only poorly constrained (bins 2-3), the flare evolution trajectory shows some flattening (bins 3-6) of the decay path."
 The origin of this behavior is unknown: 1t could be due to minor re-heating events in the same structure or produced by additional reconnection events in other magnetic structures. that again. might be triggered by the large flare.," The origin of this behavior is unknown; it could be due to minor re-heating events in the same structure or produced by additional reconnection events in other magnetic structures, that again might be triggered by the large flare."
 Another possibility is a change in size of the emitting region over the duration of the event., Another possibility is a change in size of the emitting region over the duration of the event.
 Assuming a single. loop-like emitting structure we can derive its size by adopting the formalism described in ? and parameters applicable to XMM/EPIC as given in ?..," Assuming a single, loop-like emitting structure we can derive its size by adopting the formalism described in \cite{rea97} and parameters applicable to XMM/EPIC as given in \cite{rea07}."
 By using the maximum plasma temperature and the flare decay time obtained from the light curve one can determine the size of the flaring structure. that Is treated as an instantaneously heated coronal loop.," By using the maximum plasma temperature and the flare decay time obtained from the light curve one can determine the size of the flaring structure, that is treated as an instantaneously heated coronal loop."
" Basically one finds that Lxtj.VT. with {, being the loop half-length. μις the maximum plasma temperature and r the decay time from the X-ray light curve and an additional scaling factor."," Basically one finds that $L\propto \tau_{lc} \sqrt{T_{\rm max}}$, with $L$ being the loop half-length, $T_{\rm max}$ the maximum plasma temperature and $\tau$ the decay time from the X-ray light curve and an additional scaling factor."
 This factor depends on the amount of additional heating during the decay. thus one needs to estimate the Importance of sustained heating.," This factor depends on the amount of additional heating during the decay, thus one needs to estimate the importance of sustained heating."
 This is done by fitting the decay path derived from the one temperature model with a straight line to about about one tenth of the maximum intensity., This is done by fitting the decay path derived from the one temperature model with a straight line to about about one tenth of the maximum intensity.
 This corresponds to bins 55 in the middle panel of refflp.. where we derive a slope of ~| for the decay and thus moderate re-heating has to be taken into account.," This corresponds to bins 5 in the middle panel of \\ref{flp}, where we derive a slope of $\sim 1$ for the decay and thus moderate re-heating has to be taken into account."
 Adopting the appropriate modeling parameter. Le. Του=65 MK. an e-folding decay time of r=1200 s (measured to 1/10 of the maximum) and a correction factor for a slope of one. we obtain L=1.8κιο) em (0.1 2.) for the loop half-length.," Adopting the appropriate modeling parameter, i.e. $T_{\rm obs} =65$ MK, an e-folding decay time of $\tau = 1200$ s (measured to 1/10 of the maximum) and a correction factor for a slope of one, we obtain $L = 1.8 \times 10^{10}$ cm $0.1~R_{*}$ ) for the loop half-length."
 Since à coronal loop model might not be perfectly appropriate for IQ Aur and possible heating at later phases of the flare is only poorly constrained. we use as à conservative approach our most extreme values. i.e. 100MMK plasma temperature and I400ss decay time and ignore ongoing heating.," Since a coronal loop model might not be perfectly appropriate for IQ Aur and possible heating at later phases of the flare is only poorly constrained, we use as a conservative approach our most extreme values, i.e. MK plasma temperature and s decay time and ignore ongoing heating."
 This method results in an upper limit for the size and we derive Lx3.510 em (0.2 R.) as maximum extent of the structure., This method results in an upper limit for the size and we derive $L \lesssim 3.5 \times 10^{10}$ cm $0.2~R_{*}$ ) as maximum extent of the structure.
 While the flaring structure might be much smaller. even the upper limit corresponds to a moderately sized structure when compared to the stellar dimensions of IQ Aur.," While the flaring structure might be much smaller, even the upper limit corresponds to a moderately sized structure when compared to the stellar dimensions of IQ Aur."
 This finding excludes very large magnetic structures. e.g. from the global dipole. as origin of the flare.," This finding excludes very large magnetic structures, e.g. from the global dipole, as origin of the flare."
 The IQ Aur flare is accompanied by a significant change in the metallicity of the X-ray emitting plasma as illustrated in the lower panel of refflp where we plot the results obtained from a 1-T model with metallicity as free parameter., The IQ Aur flare is accompanied by a significant change in the metallicity of the X-ray emitting plasma as illustrated in the lower panel of \\ref{flp} where we plot the results obtained from a 1-T model with metallicity as free parameter.
 During the impulsive phase of the flare as traced by the plasma temperature. more metal rich material is heated to X-ray temperatures.," During the impulsive phase of the flare as traced by the plasma temperature, more metal rich material is heated to X-ray temperatures."
 The freshly heated plasma from the large flare exhibits solar or even slightly, The freshly heated plasma from the large flare exhibits solar or even slightly
(Leroyctal.2007:Ixomugiet2011).. although whether this is due to a general deficit of molecular gas or simply a low CO-to-Lls ratio remains unclear.,"\citep{leroy07,komugi11}, although whether this is due to a general deficit of molecular gas or simply a low $_{2}$ ratio remains unclear."
 The authors would. like to thank and Whrumbolz for stimulating discussions regarding the role that molecular gas plays in star formation., The authors would like to thank and Krumholz for stimulating discussions regarding the role that molecular gas plays in star formation.
" ""They also thank SSchave and the referee. BBlack. for their. feedback on an earlier. draft of this paper."," They also thank Schaye and the referee, Black, for their feedback on an earlier draft of this paper."
 The authors acknowledge financial support from. the Landesstiftung Daden-M'ürertemberg. via. their. program International Collaboration| LL(grant. P-LS-SPIL/1S). from the German Bundesministerium fürr Bildung und Forschung via the ASTIONIZE. project STAR. FORALAVE (grant 05A00VILA). from the DFC under erants IXL1358/4 and IXL1358/5. and from a Frontier grant. of Heidelberg University sponsored by the German Excellence Initiative.," The authors acknowledge financial support from the Landesstiftung Baden-Würrrtemberg via their program International Collaboration II (grant P-LS-SPII/18), from the German Bundesministerium fürr Bildung und Forschung via the ASTRONET project STAR FORMAT (grant 05A09VHA), from the DFG under grants KL1358/4 and KL1358/5, and from a Frontier grant of Heidelberg University sponsored by the German Excellence Initiative."
 The simulations reported on in this paper were primarily performed: using the cluster at. the University of Ileidelberg. which is funded in part by theDEG via LEnini-Noether gerant. BA 3706. and. the authors would. like to thank BBanerjec. DBerentzen. GCirichdis ancl AlMarceus for the hard work that they have done to keep dvolob functioning.," The simulations reported on in this paper were primarily performed using the cluster at the University of Heidelberg, which is funded in part by theDFG via Emmy-Noether grant BA 3706, and the authors would like to thank Banerjee, Berentzen, Girichdis and Marcus for the hard work that they have done to keep functioning."
The region of reasonable parameter space for the putative jet is bracketed by M_s/v\.320.04—6 as derived from modeling of 11998bw radio light curves (Li&Chevalier1999:Waxman2004b) decay... and by ερ=0.0020.25 as inferred for most cosmological GRBs.,"The region of reasonable parameter space for the putative jet is bracketed by $\dot{M}_{-5}/v_{w,3}=0.04 - 6$ as derived from modeling of 1998bw radio light curves \citep{lc99,wax04} , and by $\epsilon_B=0.002-0.25$ as inferred for most cosmological GRBs."
 The summed constraints— (from ¢=11 to 2049 days) rule out the majority of this region., The summed constraints (from $t=11$ to $2049$ days) rule out the majority of this region.
 The search for an off-axis jet can similarly be carried out toward other local (d;<100 Mpe) Type Ibe SNe for which there are late-time (f~1—20 year) radio observations., The search for an off-axis jet can similarly be carried out toward other local $d_L < 100$ Mpc) Type Ibc SNe for which there are late-time $t\sim 1-20$ year) radio observations.
 Eight type Ibe SNe were taken from Bergeretal...(2003a):: seven upper limits 22001B.. 2200161. 2200lef. 22001]. 2200Lis. 22002). 22002bm). and a detection (SN 2002ap) at fz:0.5 years.," Eight type Ibc SNe were taken from \citet{bkf+03}: seven upper limits 2001B, 2001ci, 2001ef, 2001ej, 2001is, 2002J, 2002bm), and a detection (SN 2002ap) at $t\approx 0.5$ years."
 We have supplemented this sample with Very Large archival observations of 11983N.. II984L. 11985F. 11990B. 119941. and 11997X taken at [cm16.7.2.9.7.].0.8.8.0 and 2.4 years. respectively.," We have supplemented this sample with Very Large archival observations of 1983N, 1984L, 1985F, 1990B, 1994I, and 1997X taken at $t\approx 16.7,~2.9,~7.1,~0.8,~8.0$ and $2.4$ years, respectively."
 In addition. we include the recent detection of 22001em at [oz2.4 years (Stockdaleetaf.2004) — the only SN within this sample for which there are no early time radio observations.," In addition, we include the recent detection of 2001em at $t\approx 2.4$ years \citep{svs+04} – the only SN within this sample for which there are no early time radio observations."
 The SNe data are plotted in Figure 1.., The SNe data are plotted in Figure \ref{fig:lum_limits}.
" With the exception of 22001611. all of the late-time observations are significantly fainter than the ήςνι=1 off-axis jet prediction and nine SNe constrain M.5/v,.30.1."," With the exception of 2001em, all of the late-time observations are significantly fainter than the $\dot{M}_{-5}/v_{w,3}=1$ off-axis jet prediction and nine SNe constrain $\dot{M}_{-5}/v_{w,3}\le 0.1$."
 Using the same method applied in refsec:sn98bw.. we derive constraints for these 15. SNe.," Using the same method applied in \\ref{sec:sn98bw}, we derive constraints for these 15 SNe."
 The resulting contours are plotted with 11998bw. in Figure 3.., The resulting contours are plotted with 1998bw in Figure \ref{fig:mdot_epsb}.
 For SNe with later observations. the constraints on eg values improve. as demonstrated by the contours for 119941 and 11983N. The more robust constraints are provided by the faintest luminosity limits given by the nearest type Ic supernovae. 119941. 22002ap. 11983N and 11985F. all at <8 Mpe.," For SNe with later observations, the constraints on $\epsilon_B$ values improve, as demonstrated by the contours for 1994I and 1983N. The more robust constraints are provided by the faintest luminosity limits given by the nearest type Ic supernovae, 1994I, 2002ap, 1983N and 1985F, all at $\lesssim 8$ Mpc."
" Additional constraints are derived from the SNe for which we detect late-time emission. 119941 and22002ap*.. since they show evidence that the radio luminosity is decaying with respect to earlier lighteurve measurements (see Berger.Kulkarni&Chevalier 2002)) and are thus inconsistent with off-axis jet emission,"," Additional constraints are derived from the SNe for which we detect late-time emission, 1994I and, since they show evidence that the radio luminosity is decaying with respect to earlier lightcurve measurements (see \citealt{bkc02}) ) and are thus inconsistent with off-axis jet emission."
 The region of reasonable parameter space for an off-axis jet is roughly bracketed by M_s/vy3a2A.0.1—1 and eg=0.002—0.25 as typically inferred from broadband modeling of cosmological GRBs (Panaitescu&Kumar2002:Chevalie," The region of reasonable parameter space for an off-axis jet is roughly bracketed by $\dot{M}_{-5}/v_{w,3}=A_*=0.1-1$ and $\epsilon_B=0.002-0.25$ as typically inferred from broadband modeling of cosmological GRBs \citep{pk02,clf03,yhs+03}."
r.Li&Fransson2003:Yostetαἱ. 2003).. Figure 3. shows that four of the SNe (19941. 1985Ε. 1997X. 1990B) have contours that independently rule out this entire parameter space while eleven place significant constraints.," Figure \ref{fig:mdot_epsb} shows that four of the SNe (1994I, 1985F, 1997X, 1990B) have contours that independently rule out this entire parameter space while eleven place significant constraints."
 The late-time radio emission of SN22001em is consistent with a typical GRB jet only if the mass loss rate is M.sva%| (for egz0.2).," The late-time radio emission of 2001em is consistent with a typical GRB jet only if the mass loss rate is $\dot{M}_{-5}/v_{w,3}\approx 1$ (for $\epsilon_B\approx 0.2$ )."
" Continued radio monitoring of 22001em will reveal the origin of the emission, which may simply result from a late peaking radio SN within a dense circumstellar medium (Williamsefaf...2002)."," Continued radio monitoring of 2001em will reveal the origin of the emission, which may simply result from a late peaking radio SN within a dense circumstellar medium \citep{wpv+02}."
. Late-time observations of 11998bw (f—5.6 yrs) have allowed us to test the hypothesis that 9980425 was a standard GRB viewed far away from the jet axis., Late-time observations of 1998bw $t\simeq$ 5.6 yrs) have allowed us to test the hypothesis that 980425 was a standard GRB viewed far away from the jet axis.
 Our measured upper limits at 1384 and 2368 MHz are consistent with the continued power-law decay of the SN emission., Our measured upper limits at 1384 and 2368 MHz are consistent with the continued power-law decay of the SN emission.
 These limits imply an off-axis jet is only plausible if the normalized mass loss rate of the progenitor star Is Mνι0.04 (for eg20.1).," These limits imply an off-axis jet is only plausible if the normalized mass loss rate of the progenitor star is $\dot{M}_{-5}/v_{w,3} \le 0.04$ (for $\epsilon_B \ge 0.1$ )."
 This is ~20—200 times smaller than the observed mass loss rates for local Wolf-Rayet stars (Cappa.Goss&vanderHucht2003) and is below the range typically observed in GRBs., This is $\sim 20-200$ times smaller than the observed mass loss rates for local Wolf-Rayet stars \citep{cgv03} and is below the range typically observed in GRBs.
 Larger mass loss rates are possible but only if the energy fraction in magnetic fields Is low(i.e. ej<107).," Larger mass loss rates are possible but only if the energy fraction in magnetic fields is low, $\epsilon_B \lesssim 10^{-3}$ )."
 Even tighter constraints are derived for the off-axis jet model when we examine a larger sample of local Ibe SNe., Even tighter constraints are derived for the off-axis jet model when we examine a larger sample of local Ibc SNe.
 The low lummosity limits derived. for this sample require values of Μιςνι~0.0120.1 or ep<107 which are below values for typical GRBs.," The low luminosity limits derived for this sample require values of $\dot{M}_{-5}/v_{w,3}\approx 0.01-0.1$ or $\epsilon_B \lesssim 10^{-3}$ which are below values for typical GRBs."
 The absence of any late-time radio emission can therefore be used to put a limit on the fraction of core- SNe that produce collimated. relativistic outflows.," The absence of any late-time radio emission can therefore be used to put a limit on the fraction of core-collapse SNe that produce collimated, relativistic outflows."
 Our results imply that off-axis jets from nearby SNe are rare (<6%)) with the possible exception that the radio emission from 22001911 is due to a GRB jet., Our results imply that off-axis jets from nearby SNe are rare $\lesssim $ ) with the possible exception that the radio emission from 2001em is due to a GRB jet.
 This conclusion complements the findings of Bergereral. who constrained the GRB/SN fraction through a radio survey of local Ibe SNe at early time., This conclusion complements the findings of \citet{bkf+03} who constrained the GRB/SN fraction through a radio survey of local Ibc SNe at early time.
 Bergerefal...(20033) used early. bright radio emission as a proxy for relativistic ejecta. as in the case for 11998bw.," \citet{bkf+03} used early, bright radio emission as a proxy for relativistic ejecta, as in the case for 1998bw."
 After studying 33 local SNe with detection limits 10? times fainter than 11998bw. Bergeretal.(2003a) found no evidence for relativistic ejecta in any of the SNe observed. thereby constraining the GRB/SN fraction to <3%.," After studying 33 local SNe with detection limits $10^3$ times fainter than 1998bw, \citet{bkf+03} found no evidence for relativistic ejecta in any of the SNe observed, thereby constraining the GRB/SN fraction to $\lesssim 3\%$."
 Taken together. these results support a view that 1998bw was a rare and unusually energetic SN — distinct. from local SNe and GRBs.," Taken together, these results support a view that 1998bw was a rare and unusually energetic SN – distinct from local SNe and GRBs."
 In this scenario. the characteristics of 9980425 are not dictated by the observer's viewing angle. but rather by the properties of its central engine.," In this scenario, the characteristics of 980425 are not dictated by the observer's viewing angle, but rather by the properties of its central engine."
" 11998bw was an engine-driven explosion (Li&Chevalier 1999).. in which of the kinetic energy (10°"" erg) was coupled to mildly (Dz 2) relativistic ejecta (Kulkarnietal... 1998)... while a mere was detected in the ultra- (Dz 100) flow."," 1998bw was an engine-driven explosion \citep{lc99}, in which of the kinetic energy $\sim 10^{50}$ erg) was coupled to mildly $\Gamma\approx 2$ ) relativistic ejecta \citep{kfw+98}, while a mere was detected in the ultra-relativistic $\Gamma\approx 100$ ) flow."
 In contrast. GRBs couplemost of their energy to relativistic 7-rays.," In contrast, GRBs couplemost of their energy to relativistic $\gamma$ -rays."
 The observed diversity of cosmic explosions (SNe. X-ray flashes. and GRBs) may therefore be explained with a standard energy yield. but with a varying fraction of that energy given to relativistic ejecta (Bergeretαἱ. 2003b)..," The observed diversity of cosmic explosions (SNe, X-ray flashes, and GRBs) may therefore be explained with a standard energy yield, but with a varying fraction of that energy given to relativistic ejecta \citep{bkp+03}. ."
 We thank Edo Berger. Sarah Yost and Eli Waxman for helpful discussions.," We thank Edo Berger, Sarah Yost and Eli Waxman for helpful discussions."
 AMS ts supported by the NSFGRFP., AMS is supported by the NSFGRFP.
some planet candidates (likely to be identified by future transit searches) by detecting the orbital interactions of the planets. similar to the methods used for confirming the planets around PSIRI257212 (Rasio 1992: Malhotra 1993) and D5111620-26 (Ford et 22000).,"some planet candidates (likely to be identified by future transit searches) by detecting the orbital interactions of the planets, similar to the methods used for confirming the planets around PSR1257+12 (Rasio 1992; Malhotra 1993) and PSR1620-26 (Ford et 2000)."
 This could prove particularly. valuable for planet candidates (hat have small masses and/or orbit [aint stars. so that radial velocity confirmation is impractical (e.g.. most of the 16 transiting planet candidates orbiting faint stars recently published by Sahlu et 22006. and the many transiting planets expected to be found by future space missions).," This could prove particularly valuable for planet candidates that have small masses and/or orbit faint stars, so that radial velocity confirmation is impractical (e.g., most of the 16 transiting planet candidates orbiting faint stars recently published by Sahu et 2006, and the many transiting planets expected to be found by future space missions)."
 It would be extremely exciting to detect a£ransiting Earth-mass planet., It would be extremely exciting to detect a Earth-mass planet.
" Such a detection would enable follow-up observations to study the physical properties of the planet. such as the planets radius and density (Brown et 22001. Sato et 22005. Charbonneau οἱ 22006). the atmospheric composition (Charbonneau et 22002. Demine et 22005. Dozorgnia et 22006). and possibly even ""resolve"" surface/atmospheric features (Ford et 22001: Harrington οἱ 22006: Gaiclos et 22006)."," Such a detection would enable follow-up observations to study the physical properties of the planet, such as the planet's radius and density (Brown et 2001, Sato et 2005, Charbonneau et 2006), the atmospheric composition (Charbonneau et 2002, Deming et 2005, Bozorgnia et 2006), and possibly even “resolve” surface/atmospheric features (Ford et 2001; Harrington et 2006; Gaidos et 2006)."
 We have demonstrated that a stb-Earth-mass Trojan planet could also result in a transit timing signal that can be readily measured with ground based observatories., We have demonstrated that a sub-Earth-mass Trojan planet could also result in a transit timing signal that can be readily measured with ground based observatories.
 Since the orbital planes are likely nearly aligned. the fact that a eiant planet is already. known (o transit the star increases (he odds that other planets orbiting that star will also transit. (Holman Aburray 2005).," Since the orbital planes are likely nearly aligned, the fact that a giant planet is already known to transit the star increases the odds that other planets orbiting that star will also transit (Holman Murray 2005)."
 Thus. the transit ümine method is particularly good al searching for (ransiline Earth-like Trojan planets that would enable extremely. interesting follow-up observations.," Thus, the transit timing method is particularly good at searching for transiting Earth-like Trojan planets that would enable extremely interesting follow-up observations."
 Our technique could be applied to search for terrestrial-1iass Trojans of giant planets orbiting in (he habitable zone of their stars (Schwarz et 22005). particularly for low mass stars where the habitable zone ean be ~0.015 AU aaway [rom the star.," Our technique could be applied to search for terrestrial-mass Trojans of giant planets orbiting in the habitable zone of their stars (Schwarz et 2005), particularly for low mass stars where the habitable zone can be $\simeq~0.015$ AU away from the star."
 Once transitting terrestrial mass planets are cliscoverec. (his technique could be extended (o search for extrasolar Trojans with asteroid-like masses.," Once transitting terrestrial mass planets are discovered, this technique could be extended to search for extrasolar Trojans with asteroid-like masses."
 While 822 and previous work have emphasized (he sensitivity of transit timing observations. we eaulion that solving the inverse problem of determine planet properües [rom transil üiming observations is likely to pose a significant challenge and be more difficult. than interpreting other (vpes of extrasolar planet observations.," While 2 and previous work have emphasized the sensitivity of transit timing observations, we caution that solving the inverse problem of determing planet properties from transit timing observations is likely to pose a significant challenge and be more difficult than interpreting other types of extrasolar planet observations."
 For example. in the radial velocity method. the dominant periodictv in the observed Gane series is reaclily identified with the orbital period of a massive companion and the amplitude of the variations is proportional {ο ihe mass of the companion (xonacki Alaciejewski 1999).," For example, in the radial velocity method, the dominant periodicity in the observed time series is readily identified with the orbital period of a massive companion and the amplitude of the variations is proportional to the mass of the companion (Konacki Maciejewski 1999)."
 However. in LIV data. the dominant periodidiv could be due to any one of several physical effects. (see 22). including the reflex motion of the star due to the second planet (with a period equal to the orbital period of the second planet). the long-term mutual gravitational perturbations between the planets (with a period much longer than either orbital period). the short-term eravitational perturbations on the orbit of the transiting giant planet. (on an intermediate," However, in TTV data, the dominant periodicity could be due to any one of several physical effects (see 2), including the reflex motion of the star due to the second planet (with a period equal to the orbital period of the second planet), the long-term mutual gravitational perturbations between the planets (with a period much longer than either orbital period), the short-term gravitational perturbations on the orbit of the transiting giant planet (on an intermediate"
"visual magnitude V«7.3|L.dsin|h|. where b is the galactic latitude (the apparent visual magnitude limit was mace dependent on 6 in order to avoid ο""serving excessive numbers of stars in the Galactic plane) (ESA. 1997).","visual magnitude $V < 7.3 + 1.1{\mathrm sin}|b|$, where $b$ is the galactic latitude (the apparent visual magnitude limit was made dependent on $b$ in order to avoid observing excessive numbers of stars in the Galactic plane) (ESA, 1997)."
 Phis V maenitude limit results in 209 stars., This $V$ magnitude limit results in 209 stars.
 To :wigment the basic sample we increased. the magnitude Imi by 0.9 mag. ic. the limit is ο | 0.9.," To augment the basic sample we increased the magnitude limit by 0.9 mag, i.e. the limit is $V < 7.3 + 1.1{\mathrm sin}|b|$ + $0.9$."
 This resulted in 668 Ix. chiwarfs., This resulted in 668 K dwarfs.
 For all stars the parallax. prxyper motions and DBV colour and apparent V magnitue are available in the Llipparcos catalogue (for more details about the sample and the observations see Ixotoneva and Evnn. 2002).," For all stars the parallax, proper motions and $B - V$ colour and apparent $V$ magnitude are available in the Hipparcos catalogue (for more details about the sample and the observations see Kotoneva and Flynn, 2002)."
 The photometrically determined. metallicities for our. Ix thwarts is only appropriate for single stars., The photometrically determined metallicities for our K dwarfs is only appropriate for single stars.
 Identifving the probable multiple stars was therefore very important., Identifying the probable multiple stars was therefore very important.
 The llipparcos catalog includes a Lag for so called. “probable multiple stars. based on the quality. of the parallax. and proper motion solutions.," The Hipparcos catalog includes a flag for so called “probable multiple stars”, based on the quality of the parallax and proper motion solutions."
" The stars were divided: on this basis into ""probable single stars” ancl “probably multiple stars.", The stars were divided on this basis into “probable single stars” and “probably multiple stars”.
 We show the colour magnitude diagrams of the two eroups in figures 1 and 2., We show the colour magnitude diagrams of the two groups in figures 1 and 2.
 Overlaid on the diagrams is a solar metallicity isochrone (with an age of Ll Gyr) [rom Jimenez. Flynn and Ixotoneya (1998).," Overlaid on the diagrams is a solar metallicity isochrone (with an age of 11 Gyr) from Jimenez, Flynn and Kotoneva (1998)."
 Phe probable single stars are scattered around this line much less than the multiple stars. as one would expect.," The probable single stars are scattered around this line much less than the multiple stars, as one would expect."
 Removing probable multiple stars reduced. the sample to 449 stars (or about 2/3 of the initial sample)., Removing probable multiple stars reduced the sample to 449 stars (or about 2/3 of the initial sample).
 Despite this expedient. a small number of binaries seem to remain in the sample (as covered. in detail in section 4. Fig 10).," Despite this expedient, a small number of binaries seem to remain in the sample (as covered in detail in section 4, Fig 10)."
 After removing also these suspected multiple stars. the final sample consists of 433 stars (Figure 1).," After removing also these suspected multiple stars, the final sample consists of 433 stars (Figure 1)."
 Aletallicities for 213 of the 433 single stars are available fron he sample of Ixotoneva and Flynn (2002)., Metallicities for 213 of the 433 single stars are available from the sample of Kotoneva and Flynn (2002).
 The photometric metallicities are calibrated by a sample of 34 € and Ex ciwarfs or which accurate. spectroscopically determined metallicity abundances. and elective temperatures have been determined with errors of 0.05 dex and z100 Ex. respectively (Flynn ane Morell. 1997).," The photometric metallicities are calibrated by a sample of 34 G and K dwarfs for which accurate, spectroscopically determined metallicity abundances, $_{\mathrm spec}$ and effective temperatures have been determined with errors of 0.05 dex and $\approx 100$ K, respectively (Flynn and Morell, 1997)."
" Photometric metallicities for the 213 stars come either from a method based on Geneva by and Johnson-Cousins 24 colours (Flynn and. Morell. 997). or from a method. based. on Strómmgren m, and Rf colours (Ixotoneva and Flynn. 2002)."," Photometric metallicities for the 213 stars come either from a method based on Geneva $b_1$ and Johnson-Cousins $R - I$ colours (Flynn and Morell, 1997), or from a method based on Strömmgren $m_1$ and $R
- I$ colours (Kotoneva and Flynn, 2002)."
 Both methocls eive metallicities accurate to 0.2 dex (Ixotoneva and Flynn. 2002).," Both methods give metallicities accurate to 0.2 dex (Kotoneva and Flynn, 2002)."
 1n this section. we have followed the method of Jimenez. Fivon and Ixotoneva (1998. hereafter JEIx) to see how well different isochrone sets fit our sample.," In this section we have followed the method of Jimenez, Flynn and Kotoneva (1998, hereafter JFK) to see how well different isochrone sets fit our sample."
 We have tested the Following isochrone sets from the literature:, We have tested the following isochrone sets from the literature:
function of biased tracers.,function of biased tracers.
 While our model for the real space halo correlation function £?) includes non-linear. Lagrangian bias. our perturbation theory calculations for halo infall velocities and dispersions were carried out in standard perturbation theory. where considering non-linear biasing requires the definition of a smoothing scale 1998).," While our model for the real space halo correlation function $\xi^{r}_h(r)$ includes non-linear Lagrangian bias, our perturbation theory calculations for halo infall velocities and dispersions were carried out in standard perturbation theory, where considering non-linear biasing requires the definition of a smoothing scale ."
. also found that the Eulerian local biasing scheme is not very accurate. and it is not equivalent to a Lagrangian biasing scheme1).," also found that the Eulerian local biasing scheme is not very accurate, and it is not equivalent to a Lagrangian biasing scheme."
. For these reasons. we perform our velocity calculations for the simplest model of biased tracers with bias 5. which we identify with the large scale bias we fit to the real space correlation function using LPT.," For these reasons, we perform our velocity calculations for the simplest model of biased tracers with bias $b$, which we identify with the large scale bias we fit to the real space correlation function using LPT."
 The Lagrangian perturbation theory prescription of for describing local biased objects includes both first and second order bias terms. which are related by the peak-background split to. the halo mass function.," The Lagrangian perturbation theory prescription of for describing local biased objects includes both first and second order bias terms, which are related by the peak-background split to the halo mass function."
 This theory provides a good description (accurate at the | per cent level) on scales r>25h !Mpe.," This theory provides a good description (accurate at the 1 per cent level) on scales $r > 25\,h^{-1}$ Mpc."
 We show this explicitly for scales SOA 'Mpe in Fig. 7..," We show this explicitly for scales $r < 80\,h^{-1}$ Mpc in Fig. \ref{fig:xirealPTfit},"
 and note that the fit is good on the full range of scales we study (r.<1807 'Mpe).," and note that the fit is good on the full range of scales we study $r < 180\,h^{-1}$ Mpc)."
 If we tit the LPT prediction to £r) for separations 304 'Mpe «r1804 'Mpe. we find 47=93.104.119 for 99 degrees of freedom and one free parameter (the large-scale halo bias) for the high. low. and HOD halo subsamples. respectively.," If we fit the LPT prediction to $\xi^{r}_h(r)$ for separations $30 \,h^{-1}$ Mpc $<r < 180 \,h^{-1}$ Mpc, we find $\chi^2 = 93, 104, 119$ for 99 degrees of freedom and one free parameter (the large-scale halo bias) for the high, low, and HOD halo subsamples, respectively."
 Fitting instead to linear theory gives yo=212.274.403.," Fitting instead to linear theory gives $\chi^2 = 212, 274, 403$."
 In both cases we use the standard linear theory with the Poisson sampling assumption to derive the covariance matrix (actri)Aetra, In both cases we use the standard linear theory with the Poisson sampling assumption to derive the covariance matrix $\left\langle\Delta \xi(r_i) \Delta \xi(r_j)\right\rangle$.
 To illustrate the amplitude of the second order bias corrections for the halos of interest. we also plot the LPT prediction when > is artificially set to zero (dashed-dot curve) compared with the LPT prediction including nonzero P» (solid curve): the second order bias contribution to £j is quite small for the linear halo bias values we consider.," To illustrate the amplitude of the second order bias corrections for the halos of interest, we also plot the LPT prediction when $b_2$ is artificially set to zero (dashed-dot curve) compared with the LPT prediction including nonzero $b_2$ (solid curve); the second order bias contribution to $\xi^{r}_h$ is quite small for the linear halo bias values we consider."
 A mean (pairwise halo) velocity along the pair separation vector arises from the correlation of the density field with the velocity field: where 5 is the linear halo bias., A mean (pairwise halo) velocity along the pair separation vector arises from the correlation of the density field with the velocity field: where $b$ is the linear halo bias.
 In perturbation theory. the density and velocity fields are written as a sum of terms (0=+03 +...) with the subscript denoting the order of their dependence on the linear density field. 0j(K).," In perturbation theory, the density and velocity fields are written as a sum of terms $\delta = \delta_1 + \delta_2 + \delta_3 + ...$ ), with the subscript denoting the order of their dependence on the linear density field, $\delta_1({\bf k})$ ."
 Up to fourth order in O4(&. there are three distinct corrections to the linear theory expectation vis(7) given in Eq. 7..," Up to fourth order in $\delta_1({\bf k})$, there are three distinct corrections to the linear theory expectation $v_{12}(r)$ given in Eq. \ref{linearinfall},"
" each with a different dependence on bias: The three (oj,;) terms arise from the perturbation theory corrections to P,,. and the bias dependence is the same as the linear theory term. (0,4)."," each with a different dependence on bias: The three $\left\langle\delta_i v_{4-i}\right\rangle$ terms arise from the perturbation theory corrections to $P_{\delta \theta}$, and the bias dependence is the same as the linear theory term, $\left\langle\delta_1 v_1\right\rangle$."
" The terms from three-point correlations i scale with 7. so their contribution will be larger for more (00,highly) biased tracers."," The terms from three-point correlations $\left\langle\delta_i \delta_j v_{4-i-j}\right\rangle$ scale with $b^2$, so their contribution will be larger for more highly biased tracers."
 Note that these terms are exactly the ones evaluated in Appendix B of(2011)., Note that these terms are exactly the ones evaluated in Appendix B of.
 We provide explicit expressions for all of these terms in Appendix A.., We provide explicit expressions for all of these terms in Appendix \ref{ptcalcs}. .
 Finally. the pair-weighting correction. Ε/Γ+τε]. will be larger at a given scale for more biased objects.," Finally, the pair-weighting correction, $1/[1+b^2\xi^r_m(r)]$, will be larger at a given scale for more biased objects."
 The relative contribution for these three corrections is shown in Fig., The relative contribution for these three corrections is shown in Fig.
 8. for b=2., \ref{fig:PTtot} for $b=2$.
" At least for b=2 of interest to BOSS. the two- corrections from P, never dominates. so only including the two-point corrections (as in Eq. 24))"," At least for $b=2$ of interest to BOSS, the two-point corrections from $P_{\delta\theta}$ never dominates, so only including the two-point corrections (as in Eq. \ref{scoccPT}) )"
 will be a poor model for the redshift space power spectrum: we should expect important contributions from the bispectrum as well1)., will be a poor model for the redshift space power spectrum; we should expect important contributions from the bispectrum as well.
. In the left panel of Fig., In the left panel of Fig.
 9 we compare the deviations from linear theory infall velocity predictions measured from our simulations to our perturbation theory calculation., \ref{fig:vstats} we compare the deviations from linear theory infall velocity predictions measured from our simulations to our perturbation theory calculation.
 The expectec ver) depends on the halo bias. and for this we use the firs order bias deduced from fitting the real-space halo clustering to the LPT model of(2008b):: Table |. indicates that the best tit LPT bias can differ at the few percent level from the bes fit linear bias.," The expected $v_{12}(r)$ depends on the halo bias, and for this we use the first order bias deduced from fitting the real-space halo clustering to the LPT model of; Table \ref{table:halos} indicates that the best fit LPT bias can differ at the few percent level from the best fit linear bias."
 At the percent level. Fig.," At the percent level, Fig."
 9. shows that the LPT bias predicts the correct infall velocity amplitude on the larges scales., \ref{fig:vstats} shows that the LPT bias predicts the correct infall velocity amplitude on the largest scales.
" This confirms the common assumption in the literature tha ""velocity bias? is small. at least for halos in the bias range we have studied."," This confirms the common assumption in the literature that “velocity bias” is small, at least for halos in the bias range we have studied."
 Perturbation theory provides a relatively good description of the departure from linear theory., Perturbation theory provides a relatively good description of the departure from linear theory.
 The difference between the simulations and perturbation theory depend on halo bias and agree best for the HOD halo subsample. which Fig.," The difference between the simulations and perturbation theory depend on halo bias and agree best for the HOD halo subsample, which Fig."
 7. indicates is the sample with the smallest second-order bias., \ref{fig:xirealPTfit} indicates is the sample with the smallest second-order bias.
 We note that there is good theoretical motivation to expectthe bias relevant to thematter-velocity cross-correlation to differ from the one inferred from clustering. and have scale-dependence," We note that there is good theoretical motivation to expectthe bias relevant to thematter-velocity cross-correlation to differ from the one inferred from clustering, and have scale-dependence"
explain the observed decline in disc fractions with age.,explain the observed decline in disc fractions with age.
 In effect we are asking whether there is a universal set of disc viscous parameters which can explain the variation in disc lifetime from cluster to cluster purely in terms of the observed spread in X-ray luminosity., In effect we are asking whether there is a universal set of disc viscous parameters which can explain the variation in disc lifetime from cluster to cluster purely in terms of the observed spread in X-ray luminosity.
" As discussed above the basic principle of all photoevaporation models is that discs should evolve viscously, hardly noticing the effects of photoevaporation, until the mass accretion rates in the discs have fallen to a value that is comparable to the"," As discussed above the basic principle of all photoevaporation models is that discs should evolve viscously, hardly noticing the effects of photoevaporation, until the mass accretion rates in the discs have fallen to a value that is comparable to the"
SolIlO/LASCO and radio observations from Wind spacecraft.,SoHO/LASCO and radio observations from Wind spacecraft.
 Chen&Ixrall(2003) have studied. acceleration of three CAIEs using SollO/LASCO observations. and proposed that CALE acceleration occurs in (wo phases. the main phase and (he residual phase.," \citet{Chen.Krall2003} have studied acceleration of three CMEs using SoHO/LASCO observations, and proposed that CME acceleration occurs in two phases, the `main' phase and the `residual' phase."
 While most of the acceleration occurs in the main phase. there lies a second phase of acceleration known as the residual acceleration in (he outer corona.," While most of the acceleration occurs in the main phase, there lies a second phase of acceleration known as the residual acceleration in the outer corona."
 Thev have emploved the magnetic flux rope model to show a relation between the height at (he peak of main acceleration phase. and the footpoint separation of the CME fIux rope.," They have employed the magnetic flux rope model \citep{Chen1989} to show a relation between the height at the peak of main acceleration phase, and the footpoint separation of the CME flux rope."
 In their model. Chen&Ίντα].(2003). have proposed that a change in duration of the flux injection (Ixralletal.2000). determines the strength of the residual acceleration phase.," In their model, \citet{Chen.Krall2003} have proposed that a change in duration of the flux injection \citep{Krall.etal2000} determines the strength of the residual acceleration phase."
 Similarly. Zhang&Dere(2006) have also reported (wo such phases of acceleration based on their study of 50 CMESs observed from SolIlO/LASCO.," Similarly, \citet{Zhang.Dere2006} have also reported two such phases of acceleration based on their study of 50 CMEs observed from SoHO/LASCO."
 All the studies cited above use a single viewpoint to observe the CMESs., All the studies cited above use a single viewpoint to observe the CMEs.
 The results (hen inherently suffer from projection effects of (he transients on to the plane of the sky., The results then inherently suffer from projection effects of the transients on to the plane of the sky.
 In order (o overcome (his. we decided to look al CAIEs Irom the stereoscopic vision of Solar TErrestrial RElations Observatory (8TEDBREO) (Ixaiseretal.2003).," In order to overcome this, we decided to look at CMEs from the stereoscopic vision of Solar TErrestrial RElations Observatory (STEREO) \citep{Kaiser.etal2008}."
. The STEREO spacecralt provide two viewpoints of the prominences and the associated CATES., The STEREO spacecraft provide two viewpoints of the prominences and the associated CMEs.
 We have used a stereoscopic reconstruction technique to determine (he true physical coordinates of a solar feature (Joshi&Srivastava2011)., We have used a stereoscopic reconstruction technique to determine the true physical coordinates of a solar feature \citep{Joshi.Srivastava2011}.
. The stereoscopic reconstruction would allow us {ο observe evolution of the true height of prominences aud CMESs. aud hence (heir {rue velocity and acceleration.," The stereoscopic reconstruction would allow us to observe evolution of the true height of prominences and CMEs, and hence their true velocity and acceleration."
 From this we can examine if the acceleration truly exhibits bimodal profile as the model suggests., From this we can examine if the acceleration truly exhibits bimodal profile as the model suggests.
 This will also give us a clue about the initiation anc, This will also give us a clue about the initiation and
 This will also give us a clue about the initiation ancl, This will also give us a clue about the initiation and
is typically small. with a typical correction value of 1.5.,"is typically small, with a typical correction value of 1.5."
 While it is impossible to make a strong statement about (he shape of the distribution of Beppo--SAN bursts with only the bright bursts. there is still an important result [rom the data.," While it is impossible to make a strong statement about the shape of the distribution of -SAX bursts with only the bright bursts, there is still an important result from the data."
 The data is plotted in Figure 13.., The data is plotted in Figure \ref{fig:NaPSAX}.
 We find that even among the brightest bursts. of bursts with redshifts and of bursts without redshifts are violators.," We find that even among the brightest bursts, of bursts with redshifts and of bursts without redshifts are violators."
 What is particularly provocative about (his result is Chat these are the bursts. and (hus the likely to not be violators.," What is particularly provocative about this result is that these are the bursts, and thus the likely to not be violators."
" It is unlikelv that there are a significant munber of missing, bursts that would be non-violators.", It is unlikely that there are a significant number of `missing' bursts that would be non-violators.
" Any such burst would have to be both bright and have a low Lieu, while süll being sim enough to be missed in the bright burst catalog.", Any such burst would have to be both bright and have a low $E_{peak}$ while still being sim enough to be missed in the bright burst catalog.
 Finally. we provide information as to the average energv ratio of these bursts. but we once again stress that these are only the brightest bursts. and therefore (hese values should be taken with caution.," Finally, we provide information as to the average energy ratio of these bursts, but we once again stress that these are only the brightest bursts, and therefore these values should be taken with caution."
 We therefore feel confident in saving that the Amati relation fails for Beppo--SAX bursts. alihough this statement is not as strong as it is for other detectors because of the sample used.," We therefore feel confident in saying that the Amati relation fails for -SAX bursts, although this statement is not as strong as it is for other detectors because of the sample used."
 Previously. Butler οἱ al. (," Previously, Butler et al. ("
2007) had pointed out that the normalization constant lor the Amati relation was slieghtlv clifferent depending on whether or bursts were used. and this is like noting that. (logDET) has changed.,"2007) had pointed out that the normalization constant for the Amati relation was slightly different depending on whether or bursts were used, and this is like noting that $\langle\log\frac{E^{2.04}_{peak}}{S_{bolo}}\rangle$ has changed."
 Previously. Band Preece (2005) and Goldstein οἱ al. (," Previously, Band Preece (2005) and Goldstein et al. ("
2010) pointed out that >80% of BATSE bursts violate the Amati limit.,2010) pointed out that $>80$ of BATSE bursts violate the Amati limit.
" In this section. we have generalized these analvses. both to looking al many. GRB detector insiyunents and to looking at the (wo dimensional distribution in the Spojo—E,ons diagram."," In this section, we have generalized these analyses, both to looking at many GRB detector instruments and to looking at the two dimensional distribution in the $S_{bolo} - E_{peak,obs}$ diagram."
 All of these data sets give consistent conclusions: (1) The distribution of bursts in the ο varies significantly and greatly [rom satellite-to-satellite. (," All of these data sets give consistent conclusions: (1) The distribution of bursts in the $S_{bolo} - E_{peak,obs}$ diagram varies significantly and greatly from satellite-to-satellite. ("
2) The only data sets {ο pass the generalized Nakar Piran test for the Amati relation are the early heterogeneous sample of bursts with measured spectroscopic redshilts. (,2) The only data sets to pass the generalized Nakar Piran test for the Amati relation are the early heterogeneous sample of bursts with measured spectroscopic redshifts. (
3) The bursts detected by BATSE.ιο... Suzaku. and Konus all have a high fraction (€> 70%) of bursts which violate the Amati limit. with the violations being highly significant ancl by large factors.,"3) The bursts detected by BATSE, Suzaku, and Konus all have a high fraction $\xi>70\%$ ) of bursts which violate the Amati limit, with the violations being highly significant and by large factors."
 That is. the Amati relation fails for bursts from these four satellites. (," That is, the Amati relation fails for bursts from these four satellites. ("
4) The Amati limit is satisfied for the HETE bursts. to the extent that the violator fraction is consistent ihe Amati relation plus normal observational scatter. however. the seatter in the diagram is so large that we conclude that the Amati relation does not satisfactorily apply to the ΗΤΤΙ: data. (,"4) The Amati limit is satisfied for the HETE bursts, to the extent that the violator fraction is consistent the Amati relation plus normal observational scatter, however, the scatter in the $S_{bolo} - E_{peak,obs}$ diagram is so large that we conclude that the Amati relation does not satisfactorily apply to the HETE data. ("
5) We find that no bursts. from any satellite. significantly violate the Ghirlanda limit. (,"5) We find that no bursts, from any satellite, significantly violate the Ghirlanda limit. ("
6) These conclusions are (rue whether we examine only bursts with spectroscopic redshifts or without redshifts.,6) These conclusions are true whether we examine only bursts with spectroscopic redshifts or without redshifts.
GRB explosions are DLA systems.,GRB explosions are DLA systems.
 We have taken from the literature the. EWs of those lines seen in GRB afterglow spectra and. added: these points to the diagrams (figs., We have taken from the literature the EWs of those lines seen in GRB afterglow spectra and added these points to the diagrams (figs.
 24 and 25) presented in Rao Turnshek. (2000)., 24 and 25) presented in Rao Turnshek (2000).
 The results are shown in Fig. 5.., The results are shown in Fig. \ref{fig:ratioMgII}.
 IH is evident that not only the host of (απ 0102222 would qualify as à DLA on the basis of these diagrams. but also the hosts of GRB 970508. CRB 990123. GRB 000926. and possibly of GRB 990510 and GRB 991216 (Metzger et al.," It is evident that not only the host of GRB 0102222 would qualify as a DLA on the basis of these diagrams, but also the hosts of GRB 970508, GRB 990123, GRB 000926, and possibly of GRB 990510 and GRB 991216 (Metzger et al."
 1997: Andersen et al., 1997; Andersen et al.
 1999: Ixulkarni ct al., 1999; Kulkarni et al.
 1999: Castro et al., 1999; Castro et al.
 2001: Vreeswijket al., 2001; Vreeswijk et al.
 2001)., 2001).
 These results reinforce the idea that all GItD host galaxies are DLAs., These results reinforce the idea that all GRB host galaxies are DLAs.
 In addition. Rao Turnshek (2000) conclude. apart from the act that all the systems in their sample with £2582) 007 aare DLAs. that there is no particular correlation between he EWs of 2796 aandAlel 2582 lines.," In addition, Rao Turnshek (2000) conclude, apart from the fact that all the systems in their sample with 2582) $>$ 0.7 are DLAs, that there is no particular correlation between the EWs of 2796 and 2582 lines."
 However. this changes if we adel the values for the GRB mentioned before: there is a clear trend for both EW »ng proportional to each other (see Fig. 6)).," However, this changes if we add the values for the GRB mentioned before: there is a clear trend for both EW being proportional to each other (see Fig. \ref{fig:fig5}) )."
 From the analysis of the optical spectrum of the afterglow of GRB 010222. we conclude that its host galaxy is a DLA absorption line system. with the strongest. column density of neutral hydrogen and the lowest metallicity ever known. and with low dust content.," From the analysis of the optical spectrum of the afterglow of GRB 010222, we conclude that its host galaxy is a DLA absorption line system, with the strongest column density of neutral hydrogen and the lowest metallicity ever known, and with low dust content."
 On the other hand. based on the bright millimetre ancl sub-millimetre constant emission (83.74 £0.53 mv at 350 1112). Frail et al. (," On the other hand, based on the bright millimetre and sub-millimetre constant emission (3.74 0.53 mJy at 350 GHz), Frail et al. ("
2001) conclude that the host galaxy of GRB 010222 is a dusty galaxy with an intense burst of star formation (~500 13) while in the optical and. near-infrared. it is a blue sub-Iuminous ealaxy.,2001) conclude that the host galaxy of GRB 010222 is a dusty galaxy with an intense burst of star formation 500 ) while in the optical and near-infrared it is a blue sub-luminous galaxy.
 There is an apparent contradiction. because Pettini οἱ al. (," There is an apparent contradiction, because Pettini et al. ("
1999). based. on the lack of metallicity evolution [from z=0.5 to 2=3.5. concluded that DLAs do not trace the xulk of star-Forming galaxies at these redshifts.,"1999), based on the lack of metallicity evolution from z=0.5 to z=3.5, concluded that DLAs do not trace the bulk of star-forming galaxies at these redshifts."
 Indeed. low redshift imaging studies (e.g. Le Brun et al.," Indeed, low redshift imaging studies (e.g. Le Brun et al."
 1997). have garown that DLAs represent a wide range of morphological vpes., 1997) have shown that DLAs represent a wide range of morphological types.
 At higher redshifts. the low spin temperatures inferred from the 21 cem absorption line studies. indicate xu DLAs rarely represent massive galaxies anc may »' more analogous to metal-poor dwarfs. (Chengalur. ]xanekar 2000).," At higher redshifts, the low spin temperatures inferred from the 21 cm absorption line studies indicate that DLAs rarely represent massive galaxies and may be more analogous to metal-poor dwarfs (Chengalur Kanekar 2000)."
 On the other hand. the properties (column ensity. profile shapes. velocity. dispersion. and. number of components) of lines seen in the spectra of QSOs seen ab zo ~1.5 are very similar to those observed in jj Carina nebula.," On the other hand, the properties (column density, profile shapes, velocity dispersion and number of components) of lines seen in the spectra of QSOs seen at z 1.5 are very similar to those observed in $\eta$ Carina nebula."
 This means that these lines could be associated with an expansion event. like superbubbles driven by star formation (Danks 1999).," This means that these lines could be associated with an expansion event, like superbubbles driven by star formation (Danks 1999)."
 Llowever. Boned οἱ al. (," However, Bond et al. ("
2001). after comparing the kinematics anc redshift evolution of strong absorbers. conclude that superwinds and star-forming ealaxies can account for a substantial fraction of these. but not for DLAs which show dillerent kinematics and no red-shift evolution.,"2001), after comparing the kinematics and redshift evolution of strong absorbers, conclude that superwinds and star-forming galaxies can account for a substantial fraction of these, but not for DLAs which show different kinematics and no red-shift evolution."
 Furthermore. DLA systems having N(III) =—1077E ," Furthermore, DLA systems having ) $> 10^{22}$ "
uel frequency).,high frequency).
 The analytic properties of D-spliues aud heir transform turus out handy in particular since Tavlor expansions are available when dealing with exponential xofile where the dynamical rauge is large., The analytic properties of B-splines and their transform turns out handy in particular since Taylor expansions are available when dealing with exponential profile where the dynamical range is large.
 Another useful xoperty of B-spline is extrapolation: the correlation of the spline coefficient. induced by the penalty fuuctiou vields an estimate for the behaviour of the profile beyond the ast mneasured point: since the Abel transform requires integration to infinity. this cstimate corrects in part for he tyuncation.," Another useful property of B-spline is extrapolation: the correlation of the spline coefficient induced by the penalty function yields an estimate for the behaviour of the profile beyond the last measured point; since the Abel transform requires integration to infinity, this estimate corrects in part for the truncation."
 Note that au explicit analytic continuation of the model can be added to the spline basis if required., Note that an explicit analytic continuation of the model can be added to the spline basis if required.
 Finally here the requirement is that x is smooth. which is nore strigent than requiring that X (or X7P 3} are smooth.," Finally here the requirement is that $\M{x}$ is smooth, which is more strigent than requiring that $\Sigma$ (or $\Sigma \sigma_{p}^{2}$ ) are smooth."
 Asstuuing that wei have access to discrete ucasnronucnts of+ X and SayD (viaB binningBH as discussed- above). aud that the noiseB inB X aud X67pDi can be consideredB o be Normal we can estimate the error between the neasured profiles aud the non parametric B-spline model where the weight matrix W is the inverse of the covariance matrix of the data (which is diagoual for uncorrelated noise with diagonal elements equal to one over the data variance).," Assuming that we have access to discrete measurements of $\Sigma$ and $\Sigma \sigma_{p}^{2}$ (via binning as discussed above), and that the noise in $\Sigma$ and $\Sigma \sigma_{p}^{2}$ can be considered to be Normal, we can estimate the error between the measured profiles and the non parametric B-spline model as where the weight matrix $\M{W}$ is the inverse of the covariance matrix of the data (which is diagonal for uncorrelated noise with diagonal elements equal to one over the data variance)."
 Lincar penalty functions obey where K is a positive definite matrix., Linear penalty functions obey where $\M{K}$ is a positive definite matrix.
 In practice. we use Kk=Dt-D where D is a finite difference second. order operator Iushort. the solution of (or (1))) is fouud by ninimizing the quantity Q(x}=E(x)|pRox) where L(x) and R(x) are respectively the likelihood aud regularization crus given w and(À6).. x are the (large iunuber) of parameters. and where the Lagrange multiplier HOO0 allows us to tune the level of regularization.," In practice, we use $\M{K} = \T{\M{D}}\cdot \M{D}$ where $\M{D}$ is a finite difference second order operator In short, the solution of (or ) is found by minimizing the quantity $Q(\M{x})=L(\M{x})+\mu\,R(\M{x})$ where $L(\M{x})$ and $R(\M{x})$ are respectively the likelihood and regularization terms given by and, $\M{x}$ are the (large number) of parameters, and where the Lagrange multiplier $\mu>0$ allows us to tune the level of regularization."
 The introduction of the Lagrange uniltiplicr p is formally justified by the fact that we want to minimize Q(x). subject to the constraint that L(x) should be equal to sole value.," The introduction of the Lagrange multiplier $\mu$ is formally justified by the fact that we want to minimize $Q(\M{x})$, subject to the constraint that $L(\M{x})$ should be equal to some value."
 For iustauce. with L(x)=Vx) the problem. is to minimize Q(x) subject to the coustraints that L(x) is in the range NavaracY¥2 Nai)," For instance, with $L(\M{x})=\chi^2(\M{x})$ the problem is to minimize $Q(\M{x})$ subject to the constraints that $L(\M{x})$ is in the range $N_\R{data}\pm\sqrt{2\,N_\R{data}}$ )."
 In practice. the 1ininma of Is:," In practice, the minimum of is:"
core Avcspoirsuegees from 0.03 to O16 kin |. and Acu24 ranges fom about 0.15 to 0.21 kim L,"core $\Delta v_{NT}$ ranges from 0.03 to 0.46 km $^{-1}$ , and $\Delta v_{therm}$ ranges from about 0.15 to 0.24 km $^{-1}$."
 There are rogions in the core. particularly in areas north-northeast of the central source. where the observed width is comparable to the hermal width and the non-thermal width is icelieible.," There are regions in the core, particularly in areas north-northeast of the central source, where the observed width is comparable to the thermal width and the non-thermal width is negligible."
 The ine width in the envelope peaks close to lie source. at a value of 0.18 kan !.," The line width in the envelope peaks close to the source, at a value of 0.48 km $^{-1}$."
 In lis region. thermal broadeniug is not sufficient o account for the observed line widths. aud a conibination of higher eas temperatures and large 1on-theral motions (possibly due to the outflow and an iuner iufallius region. see below) may contribute to the observed. large velocity widths.," In this region, thermal broadening is not sufficient to account for the observed line widths, and a combination of higher gas temperatures and large non-thermal motions (possibly due to the outflow and an inner infalling region, see below) may contribute to the observed large velocity widths."
 The temperature map in Figure 7 shows that the envelope temperature is highest close to. aud soutliwest of the source.," The temperature map in Figure \ref{fig:tkin} shows that the envelope temperature is highest close to, and southwest of the source."
 Thus. the relatively high lue width seen south-southwest of the source (with an average value of about 0.31 ni 8.75) is partly due to the existing higher gas temperatures.," Thus, the relatively high line width seen south-southwest of the source (with an average value of about 0.34 km $^{-1}$ ) is partly due to the existing higher gas temperatures."
 At the southwestern edge of the map. close to where the southern extension connects with the rest of the core. there is a local peak iu the line width with a value of 0.38 kins +.," At the southwestern edge of the map, close to where the southern extension connects with the rest of the core, there is a local peak in the line width with a value of 0.38 km $^{-1}$."
 As discussed above. the southern extension has a clistinct velocity compared to the rest of the envelope.," As discussed above, the southern extension has a distinct velocity compared to the rest of the envelope."
 The unresolved velocity component due to the southern extension cluission results iun appareut wider spectral lines at the southwestern edge of the envelope., The unresolved velocity component due to the southern extension emission results in apparent wider spectral lines at the southwestern edge of the envelope.
 North-nortliwest. aloug the outflow axis. there is a narrow reeion with a relatively large velocity width of about 0:31 ας. 1.," North-northwest, along the outflow axis, there is a narrow region with a relatively large velocity width of about 0.34 km $^{-1}$."
 In this region the temperature is relatively low (see Fig. 7)).," In this region the temperature is relatively low (see Fig. \ref{fig:tkin}) ),"
 aud most of the Hue broadening is due to nou-thermal motions., and most of the line broadening is due to non-thermal motions.
 We speculate that this is duc mostlv to outflow-driven turbulence (see 3.6))., We speculate that this is due mostly to outflow-driven turbulence (see \ref{sec:outflow}) ).
 The high angular resolution observations of Lee et al. (, The high angular resolution observations of Lee et al. (
"2009) detect a flattened structure within -45"" of the central source. with a velocity eracdieut consistent with a Ikeplerian or e,κ(1 rotatiou.","2009) detect a flattened structure within $\sim 0.5''$ of the central source, with a velocity gradient consistent with a Keplerian or $v_{rot} \propto r^{-1}$ rotation."
" Civeu the lengths of their projected baselimes. their observations are insensitive to structures lÓlareer than about 1"". so that if the rotating structure were more extended than what is shown by Lee et al. ("," Given the lengths of their projected baselines, their observations are insensitive to structures larger than about $4\arcsec$, so that if the rotating structure were more extended than what is shown by Lee et al. ("
2009). they would not have detected it.,"2009), they would not have detected it."
 Observations taken with shorter baselines (from Lee et al., Observations taken with shorter baselines (from Lee et al.
" 2007) show that the contiunuun enmission surrounding TT 211-nuu along the major axis (perpendicular to the outflow axis) extends approximately 1"".", 2007) show that the continuum emission surrounding HH 211-mm along the major axis (perpendicular to the outflow axis) extends approximately $4\arcsec$.
" Thus. it could very well be that the iuuer rotating structure surrounding the embedded protostar(s) has a radius of 2 to 37. or about GOO to 800 AU, which is not unusual for low-mass embedded. protostars (ee. Ohashi e al."," Thus, it could very well be that the inner rotating structure surrounding the embedded protostar(s) has a radius of 2 to $3\arcsec$, or about 600 to 800 AU, which is not unusual for low-mass embedded protostars (e.g., Ohashi et al."
 1997: Toegerheijde et al., 1997; Hogerheijde et al.
 1998)., 1998).
 If preseut. such rotating structure would not be resolved iu our svuthesized map. but it would contribute to the observed lue widths near the λα...," If present, such rotating structure would not be resolved in our synthesized map, but it would contribute to the observed line widths near the source."
 To investigate whether the )ossjbilitv of an unresolved Ikepleriau disk near the ceutral source is consistent with our data. we use a simple moclel to examine the effect that such a disk would have on the observations.," To investigate whether the possibility of an unresolved Keplerian disk near the central source is consistent with our data, we use a simple model to examine the effect that such a disk would have on the observations."
 Our tov model disk is assunied to be edge-on aud geometrically thin., Our toy model disk is assumed to be edge-on and geometrically thin.
 For a resolved geometrically thin disk with a Iepleriau velocity profile. the velocity across a disk defined by the polar coordinates + and 0 is (CGuilloteaual. 2006):: where / ds the inclination angle of thedisk (where edee-on disks have /= 907). aud the observed velocity is in a frame that is stationary with," For a resolved geometrically thin disk with a Keplerian velocity profile, the velocity across a disk defined by the polar coordinates $r$ and $\theta$ is \citep{guilloteau2006}: where $i$ is the inclination angle of thedisk (where edge-on disks have $i=90\arcdeg$ ), and the observed velocity is in a frame that is stationary with"
the zodiacal light it is difficult to determine the offset amount [rom a simple inspection of the image.,the zodiacal light it is difficult to determine the offset amount from a simple inspection of the image.
 The quadrant bias is removed. via a procedure based on the bias removal procedure developed by Mark Dickinson as part of the STscl NICMOS team (Mobasheretal.2004b)., The quadrant bias is removed via a procedure based on the bias removal procedure developed by Mark Dickinson as part of the STScI NICMOS team \citep{mob04b}.
. The procedure utilizes the flat field imprint produced on the DC! signal bv the flat. field correction process., The procedure utilizes the flat field imprint produced on the DC signal by the flat field correction process.
 Any flat DC bias will be modulated bv the variations in the flat field., Any flat DC bias will be modulated by the variations in the flat field.
 The process successively subtracts a DC bias from each quadrant before it is [lat [ieldecd. applies the flat [ield and then picks (he bias subtraction that produces the minimum variation in the quadrant.," The process successively subtracts a DC bias from each quadrant before it is flat fielded, applies the flat field and then picks the bias subtraction that produces the minimum variation in the quadrant."
 The variation in the quadrant signal is measured by a gaussian fitting to the histogram Οἱ pixel values in the images., The variation in the quadrant signal is measured by a gaussian fitting to the histogram of pixel values in the images.
 To avoid any residual corner glow [rom (he ampliliers or other quadrant boundary. anomalies only the quadrant region (hat is al least 20 pixels from the «quadrant edges is used to determine (he quadrant bias., To avoid any residual corner glow from the amplifiers or other quadrant boundary anomalies only the quadrant region that is at least 20 pixels from the quadrant edges is used to determine the quadrant bias.
 Dad pixels are also masked oul to prevent them trom dominating the variation signal., Bad pixels are also masked out to prevent them from dominating the variation signal.
 The output of variations from each bias correction is fit by both a second order polynomial and by 5 point smoothing of the output., The output of variations from each bias correction is fit by both a second order polynomial and by 5 point smoothing of the output.
 Ii the cases encountered in the UDF images they are essentially identical., In the cases encountered in the UDF images they are essentially identical.
 The minimnun) variation bias correction is selected as the minimum of the smoothed output., The minimum variation bias correction is selected as the minimum of the smoothed output.
 Both positive and negative bias are subtracted as the bias can have either a positive or negative value., Both positive and negative bias are subtracted as the bias can have either a positive or negative value.
 The bias subtraction used biases between -0.15 and 0.2 adus per second incrementecd in 0.001 adus per second., The bias subtraction used biases between -0.15 and 0.2 adus per second incremented in 0.001 adus per second.
 The procedure returns a warning if any bias corrections do not find a minimum in the provided range of biases., The procedure returns a warning if any bias corrections do not find a minimum in the provided range of biases.
 All of the UDF quadrant images had minimums within the range of biases in the procedure., All of the UDF quadrant images had minimums within the range of biases in the procedure.
 An example of the gaussian width versus subiracted bias is shown in Figure. 2.., An example of the gaussian width versus subtracted bias is shown in Figure. \ref{fig-bias}.
 Visual inspection of the images before ancl after backeround sublraction confirmed that there were no detectable remaining ceuadrant. bias olfsets., Visual inspection of the images before and after background subtraction confirmed that there were no detectable remaining quadrant bias offsets.
 The procedure would be unuecessarily (ime consuming for images where the objects were significantly brighter than the offsets and might not work in images where the width of the pixel signal histograms are dominated bv source variations rather (han noise., The procedure would be unnecessarily time consuming for images where the objects were significantly brighter than the offsets and might not work in images where the width of the pixel signal histograms are dominated by source variations rather than noise.
 Neither is the case for the UDF., Neither is the case for the UDF.
 In reality the IDL code for this procedure actually performs both the flat [ielding and bad pixel correction., In reality the IDL code for this procedure actually performs both the flat fielding and bad pixel correction.
 These procedures. however. are discussed individually in the following sections.," These procedures, however, are discussed individually in the following sections."
 NICAIOS flat fields are created internally., NICMOS flat fields are created internally.
" The ""beam steering mirror” internal to the instrument lies al a optical pupil and is used (o correct the spherical aberration of the LST primary.", The “beam steering mirror” internal to the instrument lies at a optical pupil and is used to correct the spherical aberration of the HST primary.
 It can be illuminated from behind ancl (he reflective coating of the mirror was adjusted to be about 0.01% transmittne producing an ilhuminated pupil for flat fielding., It can be illuminated from behind and the reflective coating of the mirror was adjusted to be about $0.01\%$ transmitting producing an illuminated pupil for flat fielding.
 Flat, Flat
When the ICs must e loaded with a new set of masks. tie operator finds a MIJ on the MHCU.,"When the ICs must be loaded with a new set of masks, the operator finds a MIJ on the MHCU."
 The ICs coming from the telescoye are inserted iu the IC ox. Which is placed ou the robot and he operator runs the [function of the MHSw., The ICs coming from the telescope are inserted in the IC box which is placed on the robot and the operator runs the function of the MHSw.
 Loading the ICs is a 2 step process: liis is accomplished by ΠΠ first theunload (sub)Lunctjon. at the eud of which the no-louger reeclecl masks are returned to the SC while the masks stil needed. are eft in place: the IC aud he SC tables are updated.," Loading the ICs is a 2 step process: this is accomplished by running first the (sub)function, at the end of which the no-longer needed masks are returned to the SC while the masks still needed are left in place; the IC and the SC tables are updated."
 Ther the (sub)function is started. which allows to search the 'equested mask in the SC aud to insert the masks iu [ree IC slots. by appropriately moving the IC pox.," Then the (sub)function is started, which allows to search the requested mask in the SC and to insert the masks in free IC slots, by appropriately moving the IC box."
 The search in the SC occurs manually. moving the bar code scanner uutil the requested mask is located (computer gives au aucliο signal as long as the scanner beam is located in front of the wished task).," The search in the SC occurs manually, moving the bar code scanner until the requested mask is located (computer gives an audio signal as long as the scanner beam is located in front of the wished mask)."
 The bar code is double checked by the scanuer located on the robot stand. before inserting the mask iuto the IC., The bar code is double checked by the scanner located on the robot stand before inserting the mask into the IC.
 Tre SC and IC tables are updated again aud a report (MIT) is uade available to MCS., The SC and IC tables are updated again and a report (MIT) is made available to MCS.
 The report aud the updated tables are propagated to be used for mask insertion management at the iustruuent focal plane., The report and the updated tables are propagated to be used for mask insertion management at the instrument focal plane.
 To make room in he SC [or newly manufactured masks. a MD. must. be issued whenever a mask set is uo longer required. (either the observation has been executed or expired. because of celestial constraints).," To make room in the SC for newly manufactured masks, a MDJ must be issued whenever a mask set is no longer required (either the observation has been executed or expired because of celestial constraints)."
 The function should normally have priority with respect to any other., The function should normally have priority with respect to any other.
 It allows the ope‘ator to search (as described in 6.2) for the masks to be discarded aud to remove them from the SC. thus deleting the files of the mask from the archives. updatiug the SC table aud issuing the füal report.," It allows the operator to search (as described in 6.2) for the masks to be discarded and to remove them from the SC, thus deleting the files of the mask from the archives, updating the SC table and issuing the final report."
Fig.,Fig.
 3 shows the correlation between Ίσα Luminosity and J408-MllIz raclio luminosity., \ref{fig:radir} shows the correlation between $\mu$ m luminosity and 408-MHz radio luminosity.
 Since bya is à local minimum in he typical quasar SED between the ‘bie blue bump! from he accretion disc at shorter wavelengths and the “hig red rump from thermal dust. emission at longer wavelengths. he luminosity here will be less alleeted by the extreme oxoperties of one or other of these bumps.," Since $\mu$ m is a local minimum in the typical quasar SED between the `big blue bump' from the accretion disc at shorter wavelengths and the `big red bump' from thermal dust emission at longer wavelengths, the luminosity here will be less affected by the extreme properties of one or other of these bumps."
 The correlation »etween these two properties is significant at. 2-97% confidence (Lable 3))., The correlation between these two properties is significant at $>97$ confidence (Table \ref{tab:correl}) ).
 A least-squares fit. produces a slope of 0.41+0.20. very similar to the value of 0.4d:0.1 found by Willott et ((1998) for the correlation between D-band and radio luminosity. and gaslightly flatter than the V6dE0.1 result of the original paper on this correlation by Serjeant ct ((1998).," A least-squares fit produces a slope of $0.41\pm0.20$, very similar to the value of $0.4\pm0.1$ found by Willott et (1998) for the correlation between $B$ -band and radio luminosity, and slightly flatter than the $0.6\pm0.1$ result of the original paper on this correlation by Serjeant et (1998)."
 We would expect a similar slope for he Lyanradio correlation since the small scatter in 0544 shown in Fig., We would expect a similar slope for the $\mu$ m–radio correlation since the small scatter in $\alo$ shown in Fig.
 2 ensures that the luminosities at aanel Lyin are very tightly. correlated.," \ref{fig:alpha}
 ensures that the luminosities at and $\mu$ m are very tightly correlated."
 Serjeant et eclaim that the tight opticalradio correlation provides evidence for a direct link between accretion and the fuelling of the radio jets. while Willott et ssugeest a possible cause for the discrepant. slopes.," Serjeant et claim that the tight optical–radio correlation provides evidence for a direct link between accretion and the fuelling of the radio jets, while Willott et suggest a possible cause for the discrepant slopes."
 Finally. we compare our derived optical/near-infrared spectral indices with the degree of radio core dominance.," Finally, we compare our derived optical/near-infrared spectral indices with the degree of radio core dominance."
 We compute the value of the core-dominance. parameter. R (eg. Mine Seheuer 1980). at a rest-frame frequency of (αλ. rather than the more common GGLIz. since 5GCGllIz is the most common frequency at which core measurements are made. and we therefore need. to make only small A-corrections for our objects at z~1.," We compute the value of the core-dominance parameter, $R$ (e.g., Hine Scheuer 1980), at a rest-frame frequency of GHz, rather than the more common GHz, since GHz is the most common frequency at which core measurements are made, and we therefore need to make only small $k$ -corrections for our objects at $z\sim1$."
 As ‘Table 3. shows. neither spectral index is correlated with AR.," As Table \ref{tab:correl}
 shows, neither spectral index is correlated with $R$."
 This tells us that the mechanism responsible for causing the more highlv-inclined: quasars to appear redder in the optical (e.g. Baker Llunsteacd 1995: Baker 1997) is not operating for the most luminous (3C) quasars.," This tells us that the mechanism responsible for causing the more highly-inclined quasars to appear redder in the optical (e.g., Baker Hunstead 1995; Baker 1997) is not operating for the most luminous (3C) quasars."
 We mention two possible causes of this result., We mention two possible causes of this result.
 First. our photometry does not probe wavelengths below AjS0.5primi. so we are not as sensitive to reddening ellecets (or the increased prominence of a UV-bright component. such as from an aceretion disce) as observations which probe shorter wavelengths.," First, our photometry does not probe wavelengths below $\lambda_{\rm rest} \simlt 0.5\,\mu$ m, so we are not as sensitive to reddening effects (or the increased prominence of a UV-bright component, such as from an accretion disc) as observations which probe shorter wavelengths."
 Second. and we believe more importantly. the 3C quasars have extreme radio luminosity. and hence (e.g.. Fig. 3: ," Second, and we believe more importantly, the 3C quasars have extreme radio luminosity, and hence (e.g., Fig. \ref{fig:radir}; ;"
Serjeant et 11998) extreme optical luminosity., Serjeant et 1998) extreme optical luminosity.
 According to the rececding torus model discussed in Section 4.1. the fraction of lightlvy-redcdened quasars should be much lower in this saniple than in samples containing objects with significantly lower racio luminosity.," According to the receding torus model discussed in Section 4.1, the fraction of lightly-reddened quasars should be much lower in this sample than in samples containing objects with significantly lower radio luminosity."
 Comparing our results on 2~1 3€λ objects (SRL. this paper) with lower-redshift θς sources (Llill. Goodrich DePov 1996). and with coeval but less Luminous quasars (Baker 1997). it seems clear that the fraction. of clusty τος quasars declines significantly with increasing racio IuminosiE," Comparing our results on $z \sim 1$ 3C objects (SRL, this paper) with lower-redshift 3C sources (Hill, Goodrich DePoy 1996), and with coeval but less luminous quasars (Baker 1997), it seems clear that the fraction of dusty red quasars declines significantly with increasing radio luminosity."
 We return to the two objects with apparently inverted near-infrarecl spectra. namely 3C 68.1 and ὃς 100.," We return to the two objects with apparently inverted near-infrared spectra, namely 3C 68.1 and 3C 190."
 These inverted spectra cannot result from erroneous corrections {ο the photometry., These inverted spectra cannot result from erroneous corrections to the photometry.
 In the case of 3C 68.1. even if we assume that the host galaxy has a Ravleigh-Jeans spectrum. it would need to be ~20 brighter than the mean A z relation to even produce a flat LR spectrum.," In the case of 3C 68.1, even if we assume that the host galaxy has a Rayleigh-Jeans spectrum, it would need to be $\sim 2\sigma$ brighter than the mean $K$ $z$ relation to even produce a flat IR spectrum."
 Alternatively. the contamination from emission lines to the A-band magnitude would need to be z2554.. which is highly implausible given that the strongest line in the highlv-transmissive part of the filter is the narrow line of S HI] A9532.," Alternatively, the contamination from emission lines to the $K$ -band magnitude would need to be $> 25$, which is highly implausible given that the strongest line in the highly-transmissive part of the filter is the narrow line of [S III] $\lambda$ 9532."
 We note that Ricke. Lebofsky Wisnnikewski (1982) detected 3C 68.1 at L. pum) with a [lux of 0.724EO0.17 mmy. which compares to our upper limit of mm.v at this wavelength (assuming a power law between A ancl L).," We note that Rieke, Lebofsky Wiśnnikewski (1982) detected 3C 68.1 at $L$ $\mu$ m) with a flux of $0.72\pm0.17$ mJy, which compares to our upper limit of mJy at this wavelength (assuming a power law between $K$ and $L'$ )."
 Phe two results are therefore not in conllict if 3€. 68.1 has a true Dux close to our upper limit. but note that 3€ 68.1 appears to be highly variable in the infrared. over a timescale of a lew vears (Rieke et 11982: Stein Sitko 1984).," The two results are therefore not in conflict if 3C 68.1 has a true flux close to our upper limit, but note that 3C 68.1 appears to be highly variable in the infrared over a timescale of a few years (Rieke et 1982; Stein Sitko 1984)."
 If the spectra of 3€ 68.1: and 3€ 7190 are merely intrinsically blue from UV to Ht wavelengths but. heavilv-redcdened. we would expect to find other objects with similar intrinsic spectra. but. unreddened.," If the spectra of 3C 68.1 and 3C 190 are merely intrinsically blue from UV to IR wavelengths but heavily-reddened, we would expect to find other objects with similar intrinsic spectra, but unreddened."
 Pherefore. we suspect there is a link between the red. optical and. blue infrared spectra of these two objects.," Therefore, we suspect there is a link between the red optical and blue infrared spectra of these two objects."
 Scattering by small dust grains seems. plausible since it can produce aspectrum which is, Scattering by small dust grains seems plausible since it can produce aspectrum which is
to it has also cooled.,to it has also cooled.
" In two dimensions, these same cold shells form, but they immediately break up into blobs that are RT unstable."," In two dimensions, these same cold shells form, but they immediately break up into blobs that are RT unstable."
 The cold blobs fall in to the SMBH on a free-fall time in a much more disorganized fashion compared to one dimension., The cold blobs fall in to the SMBH on a free-fall time in a much more disorganized fashion compared to one dimension.
 Figure 7 shows the power spectrum of the Eddington ratio as a function of time for the A2 simulation., Figure \ref{fig:power-spectrum} shows the power spectrum of the Eddington ratio as a function of time for the $A2$ simulation.
" In spite of the complex, detailed implementation of the physics of stellar evolution, atomic cooling, star formation, supernovae, and AGN feedback, the power"," In spite of the complex, detailed implementation of the physics of stellar evolution, atomic cooling, star formation, supernovae, and AGN feedback, the power"
observations not lag far behind the photometric imagine.,observations not lag far behind the photometric imaging.
 I this delay is (oo large. the microlensing event (hat may have amplified a quasar into the [hix-limited sample will have ended. anc the equivalent widths will no longer reflect the fact that the quasar was once lensed.," If this delay is too large, the microlensing event that may have amplified a quasar into the flux-limited sample will have ended, and the equivalent widths will no longer reflect the fact that the quasar was once lensed."
" To estimate the decorrelation timescale. consider the time it takes [or a lens moving with velocity ο (perpendicular to the line of sight) to cover a distance equal to its Einstein radius (2): where (,, is the velocity of the lens perpendicular to the line of sight."," To estimate the decorrelation timescale, consider the time it takes for a lens moving with velocity $v_\perp$ (perpendicular to the line of sight) to cover a distance equal to its Einstein radius \citep{book:sef}: : where $v_{rel}$ is the velocity of the lens perpendicular to the line of sight."
" For a conservative estimate of a 0.001AZ.. lens with a velocity dispersion of 300 kms !. this timescale is only 2.0 vr,"," For a conservative estimate of a $0.001 \Msolar$ lens with a velocity dispersion of $300$ km $^{-1}$, this timescale is only 2.9 yr."
 To address this concern. | queried the SDSS Catalog Archive for the MJD of the spectroscopic and photometric observations of the quasars.," To address this concern, I queried the SDSS Catalog Archive for the MJD of the spectroscopic and photometric observations of the quasars."
 The median time delay between observations was 1.9 vr. with a minimum of just 0.14 vr and a maximum of 2.3 vr: Figure shows a histogram of the time delavs.," The median time delay between observations was 1.9 yr, with a minimum of just 0.14 yr and a maximum of 2.3 yr; Figure \ref{fig:mjd-diffs} shows a histogram of the time delays."
 While these time differences are smaller than the limescale estimate above. μον are a potential source of worry lor the EDR data.," While these time differences are smaller than the timescale estimate above, they are a potential source of worry for the EDR data."
 It is however important to note that an order of magnitude larger lens mass pads the decorrelation timescale by a comfortable Factor of 3., It is however important to note that an order of magnitude larger lens mass pads the decorrelation timescale by a comfortable factor of 3.
 Moreover. because (he quasar survev [lux limit is for (he most part less than the characteristic flux of the luminosity function (see Figure 3)). the amplification bias is correspondinely weaker (han if the (his limit resided on the steep slope of the LF.," Moreover, because the quasar survey flux limit is for the most part less than the characteristic flux of the luminosity function (see Figure \ref{fig:Rz}) ), the amplification bias is correspondingly weaker than if the flux limit resided on the steep power-law slope of the LF."
 The fewer observations there are due to bias. the less the decorrelation between photometric aud spectroscopic measurements matters.," The fewer observations there are due to bias, the less the decorrelation between photometric and spectroscopic measurements matters."
 The first EDR-based SDSS Quasar Catalog is bv no means (he ouly flux-Iimited sample available. nor is it the largest.," The first EDR-based SDSS Quasar Catalog is by no means the only flux-limited sample available, nor is it the largest."
 The Large Bright Quasar Survey (?) set the slanclarcl for the current era of massive spectroscopic survevs. containing approximately 1/4 the number of quasars in the EDR.," The Large Bright Quasar Survey \citep*{lbqs:qsocat} set the standard for the current era of massive spectroscopic surveys, containing approximately 1/4 the number of quasars in the EDR."
 The 2dF collaboration has lar surpassed both with a publicly released spectroscopic catalog of 104 quasars (?).. and will soon release the full catalog of nearly 24.000 The SDSS also very recently announcedits “beta version” of Data Release 1. containing 17.700 quasar spectra with z<2.3.4 ," The 2dF collaboration has far surpassed both with a publicly released spectroscopic catalog of $10^4$ quasars \citep{twodf:10kcat}, and will soon release the full catalog of nearly 24,000 The SDSS also very recently announcedits “beta version” of Data Release 1, containing 17,700 quasar spectra with $z < 2.3$ "
Over the past few decades. aperture-svnthesis observations of radio galaxies have vielded. with increasing precision. maps of the svnchrotron plasma deposited by the jet pair near their extremities (hot spots) as they acvance through the ambient medium.,"Over the past few decades, aperture-synthesis observations of radio galaxies have yielded, with increasing precision, maps of the synchrotron plasma deposited by the jet pair near their extremities (hot spots) as they advance through the ambient medium."
 Already from comparison of the carly high-resolution observations mace at centimetre and metre wavelengths. (Jenkins Scheuer 1976: Cropal-WKrishna Swarup 1977). it was strongly. hinted that the the pair of radio lobes arising from the backtlow from the two hotspots are often separated by an emission gap which is probably real. and not merely ai result. of ractiative losses.," Already from comparison of the early high-resolution observations made at centimetre and metre wavelengths (Jenkins Scheuer 1976; Gopal-Krishna Swarup 1977), it was strongly hinted that the the pair of radio lobes arising from the backflow from the two hotspots are often separated by an emission gap which is probably real, and not merely a result of radiative losses."
 Such emission gaps. particularly the largest of them. will be the focus of the present study.," Such emission gaps, particularly the largest of them, will be the focus of the present study."
" According to one interpretation. the central gap in the ""radio bridge"" arises [rom blocking of the backllowing radio plasma by the denser interstellar medium. (18M) of the parent. elliptical galaxy (e.g... Leahy Williams 1984: Alexander LOST: Wiita Norman 1992)."," According to one interpretation, the central gap in the “radio bridge"" arises from blocking of the backflowing radio plasma by the denser interstellar medium (ISM) of the parent elliptical galaxy (e.g., Leahy Williams 1984; Alexander 1987; Wiita Norman 1992)."
 Llowever. the general applicability of this mechanism. was found to be at odds with the observation that in several powerful radio galaxies at moderate recshifts. the radio gaps have sharp. quasi-linear edges and some of these have wiclths approaching or even exceeding LOO kpe (CGopal-Ixrishna Wiita 1996: Gopal-Ixrishna Nath 1997: Gopal-Ixrishna Wiita 2000. hereafter CIN00).," However, the general applicability of this mechanism was found to be at odds with the observation that in several powerful radio galaxies at moderate redshifts, the radio gaps have sharp, quasi-linear edges and some of these have widths approaching or even exceeding 100 kpc (Gopal-Krishna Wiita 1996; Gopal-Krishna Nath 1997; Gopal-Krishna Wiita 2000, hereafter GKW00)."
" In these papers we argued that such emission. gaps. or superdisks"". can provide consistent. explanations [or some Κον correlations found among radio source properties. such as the Laine-Garrington cllect (Garrington et 11988: Laine LOSS: Ciurington. Conway Leahy 1991). the correlated: radio-optical asvnunetrics (e.g... AleCarthy. van Breugel Ixapahi 1991). the preference. of multiple absorption dips to occur in the Ly-a profiles of high-z racio galaxies with overall sizes below 235 kpc (van Oijk ct 11997: Binette et 22006) and metrewave (lux variability via “superluminal refractive scintillations” (Giopal-Ixrishna 1991: also. Campbell-Wilson Ltunstead 1904: Ferrara"," In these papers we argued that such emission gaps, or “superdisks"", can provide consistent explanations for some key correlations found among radio source properties, such as the Laing-Garrington effect (Garrington et 1988; Laing 1988; Garrington, Conway Leahy 1991), the correlated radio-optical asymmetries (e.g., McCarthy, van Breugel Kapahi 1991), the preference of multiple absorption dips to occur in the $\alpha$ profiles of $z$ radio galaxies with overall sizes below $\sim$ 35 kpc (van Oijk et 1997; Binette et 2006) and metrewave flux variability via “superluminal refractive scintillations"" (Gopal-Krishna 1991; also, Campbell-Wilson Hunstead 1994; Ferrara"
process: most matter Lows in in a strongly anisotropic fashion through filamentary extensions into the neighboring large scale matter distribution.,process: most matter flows in in a strongly anisotropic fashion through filamentary extensions into the neighboring large scale matter distribution.
 As a result. we can expect that many halos will not have settled into a perfect virial state.," As a result, we can expect that many halos will not have settled into a perfect virial state."
 This will certainly be the case for halos that recently sullered a major merger with one or more neighboring clumps., This will certainly be the case for halos that recently suffered a major merger with one or more neighboring clumps.
 The detailed accretion and merging history is a function of the underlying cosmology., The detailed accretion and merging history is a function of the underlying cosmology.
 Low density cosmologics or cosmologies with a high cosmological constant will have frozen their structure formation at early epochs., Low density cosmologies or cosmologies with a high cosmological constant will have frozen their structure formation at early epochs.
 The halos that had formed. by the time of that transition will have had ample time to settle into a perfect. virializedl object., The halos that had formed by the time of that transition will have had ample time to settle into a perfect virialized object.
 Also. there is a dependence on the power spectrum of the corresponding structure formation scenario.," Also, there is a dependence on the power spectrum of the corresponding structure formation scenario."
 Power spectra with a slope n<1.5 (at cluster scales) will imply a more homologous collapse of the cluster sized clumps. less marked by an incessant bombardmoent by smaller clumps.," Power spectra with a slope $n<-1.5$ (at cluster scales) will imply a more homologous collapse of the cluster sized clumps, less marked by an incessant bombardment by smaller clumps."
 Ht may be clear that a more violent life history of a halo will usually be reflected in a substantial deviation from a perfect. virial state., It may be clear that a more violent life history of a halo will usually be reflected in a substantial deviation from a perfect virial state.
 In order to investigate the implications of a dilference in accretion or merging history of halos. we have split the samples of cluster halos in each. of our cosmologies into a and asemple.. Possible dilferences in their virial state should be reflected in the quality of the scaling relations. in particular that of the thickness of the Fundamental Plane.," In order to investigate the implications of a difference in accretion or merging history of halos, we have split the samples of cluster halos in each of our cosmologies into a and a. Possible differences in their virial state should be reflected in the quality of the scaling relations, in particular that of the thickness of the Fundamental Plane."
 The consists of those halos that sulfered a merger with another halo that contained at. least of its mass., The consists of those halos that suffered a merger with another halo that contained at least of its mass.
 Fig., Fig.
 14. shows two examples of halos in the ACDAIF2 cosmology., \ref{fig:merging_accretion} shows two examples of halos in the $\Lambda$ CDMF2 cosmology.
 The top sequence of 4 panels shows he evolution of a quicsencently evolvinghalo. w means of the particledistribution in a 5h+ Mpe box (comoving size) around the cluster core. at z=2.61. z=1.61. z=0.89 and +=0.00.," The top sequence of 4 panels shows the evolution of a quiesencently evolving, by means of the particledistribution in a $5h^{-1}$ Mpc box (comoving size) around the cluster core, at z=2.61, z=1.61, z=0.89 and $z=0.00$."
 Phe circles indicate the location of the LOD identified halos. with the size of the circle proportional ο the radius of the halo (note that the overlap of circles is due to projection of the halo spheres).," The circles indicate the location of the HOP identified halos, with the size of the circle proportional to the radius of the halo (note that the overlap of circles is due to projection of the halo spheres)."
 “Phe lower group of4 xuiels shows the particle distribution at the same redshifts or a halo belonging to the sample., The lower group of 4 panels shows the particle distribution at the same redshifts for a halo belonging to the .
 Ws gradual, Its gradual
we used the unreddened ccolours.,we used the unreddened colours.
 The second step provided the necessary elimination of offsets between the photometries: in each cluster we selected stars belonging to the P component. computed the average offset in colour between oobserved and derived from the colour-temperature relations and used it to define the zero point for the P stars (so that they have zero colour by definition).," The second step provided the necessary elimination of offsets between the photometries: in each cluster we selected stars belonging to the P component, computed the average offset in colour between observed and derived from the colour-temperature relations and used it to define the zero point for the P stars (so that they have zero colour by definition)."
 Only after this normalisation we computed the offsets in colours between the three populations., Only after this normalisation we computed the offsets in colours between the three populations.
 We have 319. 587. and 39 stars in the P. I. E components. respectively.," We have 319, 587, and 39 stars in the P, I, E components, respectively."
 We only considered stars fainter than Mj.=—3.5 because for brighter stars the different sequences for different He are too close and becomeindistinguishable?., We only considered stars fainter than $M_K=-3.5$ because for brighter stars the different sequences for different He are too close and become.
. Results for the 19 GCs are given in Table |.. where we indicate the cluster. the metallicity (Col.," Results for the 19 GCs are given in Table \ref{t:offset}, where we indicate the cluster, the metallicity (Col."
 2. taken from Paper VIID. the number of stars in the P component (Col.," 2, taken from Paper VIII), the number of stars in the P component (Col."
 3) and their average colour (V—K)p and r.m.s. (, 3) and their average colour $(V-K)_P$ and r.m.s. (
Cols.,Cols.
 4. 5). the number of stars in the I component (Col.," 4, 5), the number of stars in the I component (Col."
 6) and their average offset with respect to the P stars. with r.m.s. (," 6) and their average offset with respect to the P stars, with r.m.s. ("
Cols.,Cols.
 7. 8); Cols 9 to 12 show the average metallicities of the P and | components. with their r.m.s..," 7, 8); Cols 9 to 12 show the average metallicities of the P and I components, with their r.m.s.,"
 and Col., and Col.
 13 shows the difference in [Fe/H] (see next Section)., 13 shows the difference in [Fe/H] (see next Section).
 We did not evaluate the corresponding individual cluster values for the E component since the number of E stars is very small and they would not be significant., We did not evaluate the corresponding individual cluster values for the E component since the number of E stars is very small and they would not be significant.
 Figure | shows that differences in colours and in metallicity are well anticorrelated as indicated by the Pearson and Spearman rank coefficients: this is significant at better than the level., Figure \ref{f:dfedvk} shows that differences in colours and in metallicity are well anticorrelated as indicated by the Pearson and Spearman rank coefficients; this is significant at better than the level.
 To see whether there is a difference in colour among the different populations we summed up the P. I. and E components in all clusters and obtained the following weighted ο We see that A(V—K) is of the same order of the error. hence not significant. for the P and I stars. while it is about 4 times the error for the P and E stars.," To see whether there is a difference in colour among the different populations we summed up the P, I, and E components in all clusters and obtained the following weighted : We see that $\Delta(V-K)$ is of the same order of the error, hence not significant, for the P and I stars, while it is about 4 times the error for the P and E stars."
 We conclude that the extreme populations is significantly bluer than the primordial one., We conclude that the extreme populations is significantly bluer than the primordial one.
 This A(V—K) corresponds to a AY of about 0.05 or 0.08. according to the models cited. for case e) and 5). respectively.," This $\Delta(V-K)$ corresponds to a $\Delta$ Y of about 0.05 or 0.08, according to the models cited, for case $a$ ) and $b$ ), respectively."
 The differences in colour and in metallicity (see Sec., The differences in colour and in metallicity (see Sec.
 2.3) are summarised in Table 2.. while the corresponding AY values are shown in Table 3.. for simplicity.," 2.3) are summarised in Table \ref{t:delta}, , while the corresponding $\Delta$ Y values are shown in Table \ref{t:deltaY}, for simplicity."
 A complication may arise from the presence of CO bands in the K filter., A complication may arise from the presence of CO bands in the $K$ filter.
 They are stronger in P stars than in O-depleted I and E ones. and may depress the flux in the K filter by about1-2%.. making the P stars to appear bluer. hence decreasing the difference in aamong populations.," They are stronger in P stars than in O-depleted I and E ones, and may depress the flux in the $K$ filter by about, making the P stars to appear bluer, hence decreasing the difference in among populations."
 From simulations based on the CO index by Cohenetal.(1978) we have estimated that the oof P stars should beabout 0.013 nag redder than they appear., From simulations based on the CO index by \cite{cpf78} we have estimated that the of P stars should beabout 0.013 mag redder than they appear.
 That means that thecorrected, That means that thecorrected
the index p of the power-law distribution of shock-accelerated electrons is not universal.,the index $p$ of the power-law distribution of shock-accelerated electrons is not universal.
 In four of the afterglows analyzed here. the shallow fall-off of either the radio or the optical light- after the jet break requires p~1.5.," In four of the afterglows analyzed here, the shallow fall-off of either the radio or the optical light-curve after the jet break requires $p \sim 1.5$."
 Mésszárros. Rees Wijers (1998) have shown that. for a fixed p. variations 1n the jet energy per solid angle could lead to range of light-curve decay.," Mésszárros, Rees Wijers (1998) have shown that, for a fixed $p$, variations in the jet energy per solid angle could lead to range of light-curve decay."
" Because the observer receives radiation from the entire jet surface after the jet-break time 7;. the internal structure of the jet has little effect on the light-curve decay index after f;. thus we believe that the values of jp, determined by modelling the post jet-break afterglow decay are not sensitive to the angular structure of the outflow."," Because the observer receives radiation from the entire jet surface after the jet-break time $t_j$, the internal structure of the jet has little effect on the light-curve decay index after $t_j$, thus we believe that the values of $p$ determined by modelling the post jet-break afterglow decay are not sensitive to the angular structure of the outflow."
 We note that. for a fractional energy electron ε close to equipartition. the hard electron distributions (p.<2) identified in the 991208. 991216. 000301c. and 010222 afterglows. lead toav.-break passing through the optical band at/after few days. yielding the steepening seen in the optical emission of these afterglows.," We note that, for a fractional energy electron $\epsilon$ close to equipartition, the hard electron distributions $p < 2$ ) identified in the 991208, 991216, 000301c, and 010222 afterglows, lead to a $\nu_*$ -break passing through the optical band at/after few days, yielding the steepening seen in the optical emission of these afterglows."
" Our modelling of the broadband emission of ten afterglows reveals several properties of GRB jets. which represent constraints on the models for GRB progenitors (Woosley 1993. Paezyfisski 1998. Vietri Stella 1998. MacFadyen Woosley 1999, Mésszárros. Rees Wijers 1999, MacFadyen. Woosley Heger 2001): 1) the jet energy has a relatively narrow distribution. the values determined here being within a factor of 5. around 5s10 erg. 2) the jet initial Lorentz factor is between ~LOO and ) narrower jets are less massive and more relativistic than WICer Jets 1) the baryonic mass encountered by the jet (as it breaks out) is less than 10.! of the material that the GRB progenitor had initially within the jet aperture. 5) the surrounding medium does not have. in general. the +? profile expected for the unperturbed wind of à massive GRB progenitor."," Our modelling of the broadband emission of ten afterglows reveals several properties of GRB jets, which represent constraints on the models for GRB progenitors (Woosley 1993, Paczyńsski 1998, Vietri Stella 1998, MacFadyen Woosley 1999, Mésszárros, Rees Wijers 1999, MacFadyen, Woosley Heger 2001): $1)$ the jet energy has a relatively narrow distribution, the values determined here being within a factor of 5, around $\sim 5 \times 10^{50}$ erg, $2)$ the jet initial Lorentz factor is between $\sim 100$ and $3)$ narrower jets are less massive and more relativistic than wider jets $4)$ the baryonic mass encountered by the jet (as it breaks out) is less than $10^{-4}$ of the material that the GRB progenitor had initially within the jet aperture, $5)$ the surrounding medium does not have, in general, the $r^{-2}$ profile expected for the unperturbed wind of a massive GRB progenitor."
 In most cases we find that the density of the external medium is between 0.1ciu. and 100cum., In most cases we find that the density of the external medium is between $0.1\cm3$ and $100\cm3$.
 The conclusions and the jet parameters presented here were obtained by modelling the afterglow data within a specific framework and under certain assumptions. the most notable being the uniformity of the jet and the constancy of the energy release parameters (©... 2p).," The conclusions and the jet parameters presented here were obtained by modelling the afterglow data within a specific framework and under certain assumptions, the most notable being the uniformity of the jet and the constancy of the energy release parameters $\epsel$, $\epsmag$ )."
 For simplicity. the observer was located on the jet symmetry axis.," For simplicity, the observer was located on the jet symmetry axis."
 Until the time when first observations are done (few hours to | day). the narrow jets considered here undergo significant lateral spreading. so that the afterglow light-curves seen by an observer located off the jet axis (but still within the initial jet opening. to allow the GRB to be detected and localized) differ little from those seen by an on-axis observer.," Until the time when first observations are done (few hours to 1 day), the narrow jets considered here undergo significant lateral spreading, so that the afterglow light-curves seen by an observer located off the jet axis (but still within the initial jet opening, to allow the GRB to be detected and localized) differ little from those seen by an on-axis observer."
 More complex jet models for GRB afterglows. such as that of a structured jet proposed by Rossi (2001). or a hydrodynamical treatment of the jet lateral spreading (Granot 2001). may yield different jet parameters and constraints on GRB progenitors than presented here.," More complex jet models for GRB afterglows, such as that of a structured jet proposed by Rossi (2001), or a hydrodynamical treatment of the jet lateral spreading (Granot 2001), may yield different jet parameters and constraints on GRB progenitors than presented here."
 We note that the existence of a quasi-universal jet energy has also been established in a less model-dependent way by Piran (2001). based on the narrow width of the afterglow X-ray luminosity at 1/2 day. when a good fraction of the entire jet is visible to the observer. thus this property should also be present in more sophisticated jet models.," We note that the existence of a quasi-universal jet energy has also been established in a less model-dependent way by Piran (2001), based on the narrow width of the afterglow $X$ -ray luminosity at 1/2 day, when a good fraction of the entire jet is visible to the observer, thus this property should also be present in more sophisticated jet models."
beam distance from VLA B0253+192.,beam distance from VLA B0253+192.
 To separate a possible contribution of JO25616+192703. we subtracted the emission of VLA B0253-192 from the NVSS image adopting a circular beam size.," To separate a possible contribution of J025616+192703, we subtracted the emission of VLA B0253+192 from the NVSS image adopting a circular beam size."
 This way we find an upper limit of 0.3 mJy for the flux of JO25616+192703 at MMHz., This way we find an upper limit of 0.3 mJy for the flux of J025616+192703 at MHz.
 The spectral index derived with that value would close to or larger than zero. thus synchrotron radiation can be excluded.," The spectral index derived with that value would close to or larger than zero, thus synchrotron radiation can be excluded."
 The spectrum 1s thus clearly different from that of BB0253+192. making à physical association between both sources unlikely.," The spectrum is thus clearly different from that of B0253+192, making a physical association between both sources unlikely."
 Both sources are also visible in the Spitzer 244m. map (see Fig. 4)., Both sources are also visible in the Spitzer $\mu$ m map (see Fig. \ref{l1457-mips}) ).
 Fluxes and positions are given in Table 3.., Fluxes and positions are given in Table \ref{l1457-sources}.
 The Spitzer maps at uim and 170im are unfortunately disturbed and cannot be used to determine the infrared spectra of the sources., The Spitzer maps at $\mu$ m and $\mu$ m are unfortunately disturbed and cannot be used to determine the infrared spectra of the sources.
 In the IRAS im map and our bolometer map the emission Is extensive. so we can only determine upper limits for the source fluxes at those wavelengths.," In the IRAS $\mu$ m map and our bolometer map the emission is extensive, so we can only determine upper limits for the source fluxes at those wavelengths."
 As can be seen in Fig., As can be seen in Fig.
 3. the source JO25616+192703 is projected precisely on a local emission maximum in our bolometer map of L11457-C. A chance alignment of a background source is very low., \ref{vla} the source J025616+192703 is projected precisely on a local emission maximum in our bolometer map of 1457-C. A chance alignment of a background source is very low.
 It is therefore possible that this source is indeed associated with the dense core. possibly a protostellar condensation. as a Class 0 object. deeply embedded in the core.," It is therefore possible that this source is indeed associated with the dense core, possibly a protostellar condensation, as a Class 0 object, deeply embedded in the core."
 The few points of the spectrum of JO25616+192703 that we know so far are consistent with established Class 0 sources., The few points of the spectrum of J025616+192703 that we know so far are consistent with established Class 0 sources.
 For example. 00419141522 (André et al. 1999))," For example, 04191+1522 (André et al. \cite{andre:etal99}) )"
 would have a similar infrared intensity at 244 (Dunham et al. 2006)).," would have a similar infrared intensity at $\mu$ m (Dunham et al. \cite{dunham:etal06}) ),"
 if put at the same distance., if put at the same distance.
 This source is also seen at radio wavelength with a similar flux (André et al. 1999))., This source is also seen at radio wavelength with a similar flux (André et al. \cite{andre:etal99}) ).
 Recently. Kauffmann et al. (2005))," Recently, Kauffmann et al. \cite{kauffmann:etal05}) )"
 and Bourke et al. (2006)), and Bourke et al. \cite{bourke:etal06}) )
 have detected protostellar objects with very low luminosities., have detected protostellar objects with very low luminosities.
 These sources cannot be distinguished from the extensive dust emission at mm-wavelength. but are more easily detected at shorter wavelengths. similar to J0256164-192703.," These sources cannot be distinguished from the extensive dust emission at mm-wavelength, but are more easily detected at shorter wavelengths, similar to J025616+192703."
 Without further points in the spectrum it is. however. not possible to draw more conclusions about the evolutionary state about our source.," Without further points in the spectrum it is, however, not possible to draw more conclusions about the evolutionary state about our source."
 We have presented a bolometer map at 1.2mm. which shows at least three dense dust condensations with peak H» column densities of >1077 cem? and solar masses.," We have presented a bolometer map at 1.2mm, which shows at least three dense dust condensations with peak $_2$ column densities of $\ge10^{22}$ $^{-2}$ and solar masses."
 These are closer to virial equilibrium than is the source as a whole., These are closer to virial equilibrium than is the source as a whole.
 Towards one of the condensations. we find two point sources seen in the radio regime at cem and in the infrared at 24m. The radio spectral index of one of the sources indicates nonthermal emission. thus the source is most likely a background object.," Towards one of the condensations, we find two point sources seen in the radio regime at cm and in the infrared at $\mu$ m. The radio spectral index of one of the sources indicates nonthermal emission, thus the source is most likely a background object."
 For the other source the spectrum is only sparsely known., For the other source the spectrum is only sparsely known.
 The exact location of the source at a local maximum in L.11457-C suggests a physical assocation with the cloud., The exact location of the source at a local maximum in 1457-C suggests a physical assocation with the cloud.
 We speculate that this source could be a protostellar condensation that is still deeply embedded in the core., We speculate that this source could be a protostellar condensation that is still deeply embedded in the core.
 Its low luminosity and the detection only at radio and infrared wavelengths could be caused by an object with a temperature similar to that of the surrounding material., Its low luminosity and the detection only at radio and infrared wavelengths could be caused by an object with a temperature similar to that of the surrounding material.
 To support this hypothesis more observations of the spectrum of this object between radio and infrared wavelength or a search for an associated outflow are necessary., To support this hypothesis more observations of the spectrum of this object between radio and infrared wavelength or a search for an associated outflow are necessary.
 Possibly owing to the multiple velocity components in the CO data of L1457 (Zimmermann 1993)) no such outflow has been detected so far., Possibly owing to the multiple velocity components in the CO data of L1457 (Zimmermann \cite{zimmermann93}) ) no such outflow has been detected so far.
 If the protostellar nature of the source is confirmed. the study of this source will significantly increase our knowledge of low-mass star formation outside of the Galactic plane.," If the protostellar nature of the source is confirmed, the study of this source will significantly increase our knowledge of low-mass star formation outside of the Galactic plane."
weak lensing detection of IHToekstra et al. (,weak lensing detection of Hoekstra et al. (
2001).,2001).
" We will also present the results when the sample of groups is split iuto two samples of ""rich galaxw groups aud “poor ealaxy groups. divided by velocity dispersion."," We will also present the results when the sample of groups is split into two samples of “rich galaxy groups” and “poor galaxy groups”, divided by velocity dispersion."
 A second paper will follow with the results of ealaxy-galaxy lensing in these fields aud a masta likelihood analysis of the shear., A second paper will follow with the results of galaxy-galaxy lensing in these fields and a maximum likelihood analysis of the shear.
 Our galaxy group catalogs were geuerated using a friends-of-fricnds aleorithin with the CNOC2 redshift survey data (Yee et al., Our galaxy group catalogs were generated using a friends-of-friends algorithm with the CNOC2 redshift survey data (Yee et al.
 2000: Carlbere et al., 2000; Carlberg et al.
 2001)., 2001).
 The CNOC?2 area contains | fields well-spaced iu right ascension and was iutended to better understaud the properties of field galaxies., The CNOC2 area contains 4 fields well-spaced in right ascension and was intended to better understand the properties of field galaxies.
 The CNOC2 ealaxy suuple contains 6200 ealaxies with redshifts to z of 0.7., The CNOC2 galaxy sample contains 6200 galaxies with redshifts to z of 0.7.
 From this galaxv catalog a sample of 192 ealaxy groups was identified., From this galaxy catalog a sample of 192 galaxy groups was identified.
 The average nuniber of galaxies identified iu cach group is  Land the groups have a median redshift of 0.33., The average number of galaxies identified in each group is $\sim$ 4 and the groups have a median redshift of 0.33.
 The eroups have a median dyviauuically determined velocity dispersion of 190 lau |., The groups have a median dynamically determined velocity dispersion of 190 km $^{-1}$.
 For this project we observed the 1 central patches of the CNOC?2 fields. where most of the galaxy eroups are located.," For this project we observed the 4 central patches of the CNOC2 fields, where most of the galaxy groups are located."
 The observations were carried out mostly at the Canada-France-Hawail Telescope with 2 additional uiehts at the Kitt Peak National Observatory Mavall 1-ii Telescope., The observations were carried out mostly at the Canada-France-Hawaii Telescope with 2 additional nights at the Kitt Peak National Observatory Mayall 4-m Telescope.
 The fields were observed. in V... aud I..," The fields were observed in $R_c$ , and $I_c$."
" Deep exposures (71 hours) were taken iu the P, aud I. bands. which were used for the lensing micasurements."," Deep exposures $\sim$ 4 hours) were taken in the $R_c$ and $I_c$ bands, which were used for the lensing measurements."
 The characteristics of the data obtained are outlined iu Table 1., The characteristics of the data obtained are outlined in Table 1.
 Gravitational lensing is usually lited by svstematics and if is nmnuportaut to cusure no spurious shear is introduced in the stacking procedure., Gravitational lensing is usually limited by systematics and it is important to ensure no spurious shear is introduced in the stacking procedure.
 This cau be achieved by carefully monitoring the astrometry over cach input image that enters the stack., This can be achieved by carefully monitoring the astrometry over each input image that enters the stack.
 Wide-field cameras in use today have larger distortions than earlier. sanaller CCD cameras.," Wide-field cameras in use today have larger distortions than earlier, smaller CCD cameras."
 This distortion ust be properly mapped and corrected iu order to eusure no artificial source of shear is iuported curing the stacking process., This distortion must be properly mapped and corrected in order to ensure no artificial source of shear is imported during the stacking process.
 Note. however. that eroup lensing is less affected by systematics than cosmic shear studies (the weal. leusiue sjeual from the large scale structure in the Universe).," Note, however, that group lensing is less affected by systematics than cosmic shear studies (the weak lensing signal from the large scale structure in the Universe)."
 This is due to the random orientation of the ealaxy-eroup pairs across the field. as opposed to looking or a preferred orientation as cosmiüc shear studies do.," This is due to the random orientation of the galaxy-group pairs across the field, as opposed to looking for a preferred orientation as cosmic shear studies do."
 Cosmic shear measurements use the patterus iu the laree-scale distortion feld of background sources to map out he matter distribution in the Universe., Cosmic shear measurements use the patterns in the large-scale distortion field of background sources to map out the matter distribution in the Universe.
 This signal is inv and more susceptible to svstematics than galaxy-ealaxy or eroup-galaxy lensing where the shear sigual is averaged in radial bius around each leus., This signal is tiny and more susceptible to systematics than galaxy-galaxy or group-galaxy lensing where the shear signal is averaged in radial bins around each lens.
 In this analysis he image reduction aud stacking was carried out sine he IRAF inosaic package mscred (Valdes. E.C.. 1997).," In this analysis the image reduction and stacking was carried out using the IRAF mosaic package mscred (Valdes, F.G., 1997)."
 Object catalogs were extracted from our stacked nuages using the iucat software. an iuplementation of the I&kaiser. Squires and Broadhurst (1995. hereafter ISB) method.," Object catalogs were extracted from our stacked images using the imcat software, an implementation of the Kaiser, Squires and Broadhurst (1995, hereafter KSB) method."
 This software is optimized for 1icasuriug the shapes of faint sources., This software is optimized for measuring the shapes of faint sources.
 The object detection algoritlin works by smoothing the images usine different sized filters aud then detecting the “peaks” which are then added to the source catalog., The object detection algorithm works by smoothing the images using different sized filters and then detecting the “peaks” which are then added to the source catalog.
" For each detected object weighted quadrupole moments were measured aud the resulting polarizations were calculated: The polarization nieasuremieuts need to be corrected for the effects of ποσο, camera distortion and PSF anisotropy."," For each detected object weighted quadrupole moments were measured and the resulting polarizations were calculated: The polarization measurements need to be corrected for the effects of seeing, camera distortion and PSF anisotropy."
 These corrections to correct for these concerns lave been discussed in NSB aud Luppiuo Kaiser (1997) with some improvements made ly Hoekstra et al (, These corrections to correct for these concerns have been discussed in KSB and Luppino Kaiser (1997) with some improvements made by Hoekstra et al. (
1998 aud 2000).,1998 and 2000).
 The techniques work well for ground-based data where the PSF is stable aud not very anisotropic. aud where the fields contain iiy stars which are used in the correction aleorithius.," The techniques work well for ground-based data where the PSF is stable and not very anisotropic, and where the fields contain many stars which are used in the correction algorithms."
 The source catalogs are trimmed so that all stars are removed., The source catalogs are trimmed so that all stars are removed.
 The stars can easily be located by comparing magnitude and half light radius., The stars can easily be located by comparing magnitude and half light radius.
 We kept oulv those objects for which the half helt radi were greater than 1.2 times the stellar PSF. thus eusuriug the contamination frou stars In our source catalog is very small," We kept only those objects for which the half light radii were greater than 1.2 times the stellar PSF, thus ensuring the contamination from stars in our source catalog is very small."
 The Πιο magnitude of our inages is approximately 25 in 14..., The limiting magnitude of our images is approximately 25 in $_c$ .
 Iu the weak leusiug analysis we used a source catalog of approximately 150 000 objects (~ LO per sq arcninute) and a galaxy group catalogs containing the 116 CNOC?2 ealaxy eroup centers that were within the area we observed., In the weak lensing analysis we used a source catalog of approximately 150 000 objects $\sim$ 40 per sq arcminute) and a galaxy group catalog containing the 116 CNOC2 galaxy group centers that were within the area we observed.
 The faint members of the galaxy groups are included iu the source catalog. but as shown by IHockstra et al. (," The faint members of the galaxy groups are included in the source catalog, but as shown by Hoekstra et al. ("
2001). this does not coutanunate the fiual result.,"2001), this does not contaminate the final result."
 This is incicated by the fact that the πας density of faint galaxies does not increase significautlv towards the group center. thus fünt group ieniboers are do nof influence the fal shear measurement.," This is indicated by the fact that the number density of faint galaxies does not increase significantly towards the group center, thus faint group members are do not influence the final shear measurement."
 The source deusitv of backeround objects is not sutficicutly high to extract a signal from iudividual ealaxv groups. except for the most massive eroups. and so the ealaxy groups must be stacked ancl the weak leusing signal measured around the stacked eroups.," The source density of background objects is not sufficiently high to extract a signal from individual galaxy groups, except for the most massive groups, and so the galaxy groups must be stacked and the weak lensing signal measured around the stacked groups."
 The source galaxies around the stacked ealaxy group were divided iuto radial bius aud the average distortion was calculated in cach bins, The source galaxies around the stacked galaxy group were divided into radial bins and the average distortion was calculated in each bin.
 The component of the average distortion tangential to the eroup center is the weak. leusine signal and is displaved in Figure la., The component of the average distortion tangential to the group center is the weak lensing signal and is displayed in Figure 1a.
 The tanecutial shear is plotted inm physical bius (nuits of !Mpe) nee. the redshift of cach galaxy group is precisely known from the CNOC?2 redshift survey., The tangential shear is plotted in physical bins (units of $^{-1}$ Mpc) since the redshift of each galaxy group is precisely known from the CNOC2 redshift survey.
 Usine equation (1). the best fit isothermal sphere to the average tanecutial shear profile vielded an Γκι radius of (σσ ," Using equation (1), the best fit isothermal sphere to the average tangential shear profile yielded an Einstein radius of $\pm$ 13."
"We can 10,alternatively fit the tangential shear data witli a Navarro. Freuk and White (NEW) dark matter profile (Navarro. Freuk White. 1996)."," We can alternatively fit the tangential shear data with a Navarro, Frenk and White (NFW) dark matter profile (Navarro, Frenk White, 1996)."
 This density profile. which las been observedto fit mass distributionswell over a wide range of scales. is given by," This density profile, which has been observedto fit mass distributionswell over a wide range of scales, is given by"
compact radio continuum emission with flux densities of >4o~34jJv. +. which in (urn implies an upper limit to the intrinsic brightness temperature (corresponding to a rest frequency of ~8 GllIz) of 5.6x10? IX for any compact radio source in BRI 13350417.,"compact radio continuum emission with flux densities of $\geq 4\sigma \simeq
34~\mu$ Jy $^{-1}$, which in turn implies an upper limit to the intrinsic brightness temperature (corresponding to a rest frequency of $\sim$ 8 GHz) of $5.6 \times 10^5$ K for any compact radio source in BRI 1335–0417."
 Our coherence tests during these observations using two VLBI calibrators show that the lack of a strong point source in BRI 13350417 al the full resolution of the array cannot be due to the phase relerencing procedure., Our coherence tests during these observations using two VLBI calibrators show that the lack of a strong point source in BRI 1335–0417 at the full resolution of the array cannot be due to the phase referencing procedure.
 The 4o flux limit reported above is almost an order οἱ magnitude lower than the flux measured by the VLA (220443Jv) at 1.4 GlIIz 1999)., The $\sigma$ flux limit reported above is almost an order of magnitude lower than the flux measured by the VLA $220 \pm 43~\mu$ Jy) at 1.4 GHz \citep{CMY99}.
. This immediately implies that more than ~80% of the radio contimuun emission in BRI 13350417 is extended and not confined to the central ACN., This immediately implies that more than $\sim$ of the radio continuum emission in BRI 1335–0417 is extended and not confined to the central AGN.
 In the following. we will assess if our LIISA observations can recover the flux seen by the VLA.," In the following, we will assess if our HSA observations can recover the flux seen by the VLA."
 To do so. we applied a two-dimensional Gaussian taper falling to at 1 MÀ to the visibility data in both the w and i-directions.," To do so, we applied a two-dimensional Gaussian taper falling to at 1 $\lambda$ to the visibility data in both the - and -directions."
 This gives a beam size of 189x113 mas (1.25x0.75 kpe. P. A.= —48°). and the resulting image is shown in Figure 2.," This gives a beam size of $189 \times 113$ mas $1.25 \times 0.75$ kpc, P. A.= $-48^{\circ}$ ), and the resulting image is shown in Figure 2."
 The rms noise level in this naturally weighted image is 28 j/Jv +., The rms noise level in this naturally weighted image is $28~\mu$ Jy $^{-1}$.
 The peak flux density of the detected continuum source is 126£28 Jy +. and the total flux density is 208446pJy. which agrees well with the flux density measured with the VLA at 1.4 GlIz (Carillietal.1999).," The peak flux density of the detected continuum source is $126 \pm 28~\mu$ Jy $^{-1}$, and the total flux density is $208 \pm 46~\mu$ Jy, which agrees well with the flux density measured with the VLA at 1.4 GHz \citep{CMY99}."
. The size of the source at full width half maximum (EWIIM) is 255x138 mas (1.70.9 kpe). and the derived intrinsic brightness temperature is (3.6€0.9)xLO? Is. Figure 3 shows the continuum emission in DRI 1335.0417 at the highest angular resolutions for which there is at least 4.50 detection (0=19 pJv 1)).," The size of the source at full width half maximum (FWHM) is $255 \times 138$ mas $1.7 \times 0.9$ kpc), and the derived intrinsic brightness temperature is $(3.6 \pm 0.9) \times 10^4$ K. Figure 3 shows the continuum emission in BRI 1335–0417 at the highest angular resolutions for which there is at least $\sigma$ detection $\sigma=19~\mu$ Jy $^{-1}$ )."
 This corresponds to a resolution of 53x27 mas (0.35x0.18 kpc) in position angle —6., This corresponds to a resolution of $53 \times 27$ mas $0.35 \times 0.18$ kpc) in position angle $-6^{\circ}$.
 This image was obtained bv applving a two-dimensional Gaussian taper falling to at 5 MA in both the w and idirectious of the visibility data., This image was obtained by applying a two-dimensional Gaussian taper falling to at 5 $\lambda$ in both the - and -directions of the visibility data.
 The peak flux density of the continuum source detected ab this resolution is 8Td19 yey +. and the total [lux density is 131430 yJv.," The peak flux density of the continuum source detected at this resolution is $87 \pm 19~\mu$ Jy $^{-1}$, and the total flux density is $131 \pm 30~\mu$ Jy."
 The size ol the source is 64x35 mas (0.42x0.23 kpe) al FWHAL and its intrinsic brightness temperature is (2.240.5)x10°Ix. The results obtained at this resolution imply that about two thirds of the total radio emission emerges [rom the central ~0.3 kpe of BRI 1335.0417. and. likewise. one third arises [rom more extended scales (71.3 kpc).," The size of the source is $64 \times 35$ mas $0.42 \times 0.23$ kpc) at FWHM, and its intrinsic brightness temperature is $(2.2 \pm 0.5) \times 10^5$K. The results obtained at this resolution imply that about two thirds of the total radio emission emerges from the central $\sim$ 0.3 kpc of BRI 1335–0417, and, likewise, one third arises from more extended scales $\sim$ 1.3 kpc)."
 We have detected. 1.4 GllIz emission Irom BRI 133500417 (corresponding to a rest [rame lrequeney of ~8 Gllz) using the ISA., We have detected 1.4 GHz emission from BRI 1335–00417 (corresponding to a rest frame frequency of $\sim$ 8 GHz) using the HSA.
 At a moderate resolution (189x113 mas: Figure 2) the intrinsic brightiess temperature value of (he detected continuum structure is ~3.5x10! K. The measured {lux density at this resolution is consistent with the VLA flux density of (his source (Carillietal. 1999).., At a moderate resolution $189 \times 113$ mas; Figure 2) the intrinsic brightness temperature value of the detected continuum structure is $\sim 3.5\times 10^4$ K. The measured flux density at this resolution is consistent with the VLA flux density of this source \citep{CMY99}. .
 This implies (hat the radio continuum emission, This implies that the radio continuum emission
lost bx a Roche-lobe filling low-inass stellar companion.,lost by a Roche-lobe filling low-mass stellar companion.
 The accreting matter forms a dise whose instabilities trigecr outbursts (see Lasota 2001 for a review of the iustabilitv model)., The accreting matter forms a disc whose instabilities trigger outbursts (see Lasota \cite{l01} for a review of the instability model).
 Naravan. MeClintock Y1 (1996: see also Lasota. Naravan Yi 1996. aud Naravan. Barret MeCliutock 1997)) proposed that quicscent SAT discs are truncated and that the inner accretion flow forms au ADAF.," Narayan, McClintock Yi \cite{nmy96}; see also Lasota, Narayan Yi \cite{lny96} and Narayan, Barret McClintock \cite{nbm97}) ) proposed that quiescent SXT discs are truncated and that the inner accretion flow forms an ADAF."
 This hypothesis has beeu recently vindicated frou the theoretical point of view by Dubus. Ibuneury Lasota (P001)) are is supported by observations (see Done 2002 for a review).," This hypothesis has been recently vindicated from the theoretical point of view by Dubus, Hameury Lasota \cite{dhl01}) ) and is supported by observations (see Done \cite{cd02} for a review)."
" Naravan. Garcia AlcClintock (1997)) compared quiescent Iuninosities of SNTs supposed to contain black holes with hose of neutron-star SNTs aud realized that. in accordance with the prediction of the ADAF model. systems coutaimie black-hole ""candidates? are cdinunuer."," Narayan, Garcia McClintock \cite{ngm97}) ) compared quiescent luminosities of SXTs supposed to contain black holes with those of neutron-star SXTs and realized that, in accordance with the prediction of the ADAF model, systems containing black-hole “candidates"" are dimmer."
 They came to the couclusion that they found evidence for the presence of event horizons., They came to the conclusion that they found evidence for the presence of event horizons.
 This conclusion has been challenged. by. Chen et al. (1998)), This conclusion has been challenged by Chen et al. \cite{chetal98}) )
 who asserted that the relative dimmess of black- candidate svstenis was due solely to Naravan et al. (1997)), who asserted that the relative dimness of black-hole candidate systems was due solely to Narayan et al. \cite{ngm97}) )
 comparison method., comparison method.
 Things were clarified by Lasota ILhuneurw (1998)) who suggested. comparing systems with similar orbital period ou the assumption such systems would have similar accretion rates — the ADAF model asserting only that accreting black holes should be cinuinuer than neutron stars for the accretion rate., Things were clarified by Lasota Hameury \cite{lh98}) ) who suggested comparing systems with similar orbital period on the assumption such systems would have similar accretion rates – the ADAF model asserting only that accreting black holes should be dimmer than neutron stars for the accretion rate.
 The new method showed. however. the sale effect (Lasota Tameury 1998:; Menon et al. 1999)).," The new method showed, however, the same effect (Lasota Hameury \cite{lh98}; Menou et al. \cite{metal99}) ),"
 recently confirmed by Carcia et al. (2001)):, recently confirmed by Garcia et al. \cite{getal01}) ):
 black holes (candidates) are dinuuer than systems kuown to coutaiu neutron stars. or at least stars with surface.," black holes (candidates) are dimmer than systems known to contain neutron stars, or at least stars with surface."
 This is a very strong arguiieut in favour ofthe presence of event horizons. in fact this is the most conservative conclusion.," This is a very strong argument in favour of the presence of event horizons, in fact this is the most conservative conclusion."
 However. it is not a proof.," However, it is not a proof."
 The arguueuts against the claim that the relative cinimess of black-hole caudidates is the proof of existence of event horizons are of two. not unrelated. types.," The arguments against the claim that the relative dimness of black-hole candidates is the proof of existence of event horizons are of two, not unrelated, types."
 First. it has been argued that the accretion flow in quiescent SNTs are not represented by ADAF-s.," First, it has been argued that the accretion flow in quiescent SXTs are not represented by ADAFs."
 Naravan Yi (1995a)) aud Dlaudford Beecluan (1999]) argued. (see however Paczyviisski 1995 and Abramowicz. Lasota Teauenushchey 2000 for criticisin of the argumeut) that ADAFs are subject to mass loss and therefore the dinness of quiescent SATs could result from the low accretion rate onto the compact object - most of the matter being lost with the wind.," Narayan Yi \cite{ny95a}) ) and Blandford Begelman \cite{bb99}) ) argued (see however Paczyńsski \cite{bp98} and Abramowicz, Lasota Igumenshchev \cite{ali00} for criticism of the argument) that ADAFs are subject to mass loss and therefore the dimness of quiescent SXTs could result from the low accretion rate onto the compact object - most of the matter being lost with the wind."
 However. as shown by Aleuou et al. (19999).," However, as shown by Menou et al. \cite{metal99}) ),"
 such wind models do not offer an explanation of the luminosity difference between neutron-star systems and those xesuned to contain black holes., such wind models do not offer an explanation of the luminosity difference between neutron-star systems and those presumed to contain black holes.
" Iu fact. these authors also pointed out that the quiesceut ""nnuimositv of neutron-star binaries is not consistent with he asstuption of a ~105€ radiative efficiency."," In fact, these authors also pointed out that the quiescent luminosity of neutron-star binaries is not consistent with the assumption of a $\sim 10\%$ radiative efficiency."
 Since the attempt to apply to these systems the windy-ADAF model of Quatacrt Naravan (1999)) failed. they proposed hat the action of a magnetic propeller could be auswoer.," Since the attempt to apply to these systems the windy-ADAF model of Quataert Narayan \cite{qn99}) ) failed, they proposed that the action of a magnetic propeller could be answer."
 However. a compelling signature of this effect has vet to )o found.," However, a compelling signature of this effect has yet to be found."
 Despite of this. Abramowicz Ieumneushcheyv. (2001) sugeested that the observed differences between quiesceut ininosities of accreting black- holes and neutron stars is well explained by the occurrence iu such systems of a CDAF (Convection Dominated Accretion Flow: SCC Naravan. Iemmenushchey Abyamowicz 2000) instead of an ADAF.," Despite of this, Abramowicz Igumenshchev (2001) suggested that the observed differences between quiescent luminosities of accreting black holes and neutron stars is well explained by the occurrence in such systems of a CDAF (Convection Dominated Accretion Flow; see Narayan, Igumenshchev Abramowicz 2000) instead of an ADAF."
 Thev found that for low viscosities accretion flows around compact bodies form ADAFs ouly in their innermost regions but are couvectively dominated at radii RomWeRe where Ra=2CAfe? is the Selvarzschil radius).," They found that for low viscosities accretion flows around compact bodies form ADAFs only in their innermost regions but are convectively dominated at radii $R\gta
10^2R_{\rm S}$ (where $R_{\rm S}=2GM/c^2$ is the Schwarzschild radius)."
 Du such flows emission comes mostly from the convective region: the radiative efficieucy. is independent of accretion rate and equals spy=10.7., In such flows emission comes mostly from the convective region; the radiative efficiency is independent of accretion rate and equals $\varepsilon_{\rm BH}= 10^{-3}$.
 Assuming tha the efficiency of accretion onto a neutron star is £x«c0.1 one obtains the observed ratio between black-hole auc ueutrou-tar hnuuinosities., Assuming that the efficiency of accretion onto a neutron star is $\varepsilon_{\rm NS}\approx 0.1$ one obtains the observed ratio between black-hole and neutron-star luminosities.
 Unufortunatelv this came be the correct explanation of the hunuinositv difference (Lasota 2002)) because. as ueutioned above. neutron stars In quiesceut trausieut svstenis do not seem to accrete with a 0.1 efficiency.," Unfortunately this cannot be the correct explanation of the luminosity difference (Lasota \cite{l02}) ) because, as mentioned above, neutron stars in quiescent transient systems do not seem to accrete with a 0.1 efficiency."
 Another class of argument asserts that N-ravs du quiescent SNTs ire not euitted by the accretion flow., Another class of argument asserts that X-rays in quiescent SXTs are not emitted by the accretion flow.
 Brown. Bildsten. Rutledge (19983) suegeestedOO that. in neutron-star svstenis. uost (or all) of the quiesceut N-vay πιο] is not due to accretion but results from cooliug of the neutrou-star crust jcated by clear reactions.," Brown, Bildsten, Rutledge \cite{bbr98}) ) suggested that, in neutron-star systems, most (or all) of the quiescent X-ray luminosity is not due to accretion but results from cooling of the neutron-star crust heated by nuclear reactions."
 This crust-cooliug. model does not sccm to be in perfect agreement with observations showing two spectral components and a variable flux (sce Rutledge et al., This crust-cooling model does not seem to be in perfect agreement with observations showing two spectral components and a variable flux (see Rutledge et al.
 2002 aud references therein)., \cite{retal02} and references therein).
 Tf the crustalcooling. model were righ it would imply different X-rav endssion niechaudsnas or the two classes of quiesceut SNTs., If the crustal-cooling model were right it would imply different X-ray emission mechanisms for the two classes of quiescent SXTs.
 Towever. DIuuimositv variations observed also iu quiescent black-hole svstenis (see e.g. Garcia et al. 2001))," However, luminosity variations observed also in quiescent black-hole systems (see e.g. Garcia et al. \cite{getal01}) )"
 would rather sugecst a conunuon origin., would rather suggest a common origin.
 Attempts to ascribe quiescent N-rav huuinositv iu black-hole svsteuis to active stellar companions (Bildsten Rutledge 2000)) are not based ou a sound theoretical foundation (Lasota 2001)) and have been refuted by observations (Carcia ct al. 2001))., Attempts to ascribe quiescent X-ray luminosity in black-hole systems to active stellar companions (Bildsten Rutledge \cite{br00}) ) are not based on a sound theoretical foundation (Lasota \cite{l01}) ) and have been refuted by observations (Garcia et al. \cite{getal01}) ).
 Alenou (20011) presented an argunieut based on the settling-flow model of \Ledvecev Narayan (2001)) iu which the accretion flow arrives with very low angular momenta a the surface of à rapidly rotating compact object., Menou \cite{m02}) ) presented an argument based on the settling-flow model of Medvedev Narayan \cite{medna01}) ) in which the accretion flow arrives with very low angular momentum at the surface of a rapidly rotating compact object.
 The X-rav luninosity is then due to rotation- loss by the accreting body., The X-ray luminosity is then due to rotation-energy loss by the accreting body.
 This requires viscous contact between this body and the accreting matter., This requires viscous contact between this body and the accreting matter.
 Menou (2001)) pointed out that if black-hole candidates hac. coutrary to neutron stars. radii smaller than the," Menou \cite{m02}) ) pointed out that if black-hole candidates had, contrary to neutron stars, radii smaller than the"
the disk the solid phase is the dominant component. and because of the importance of grain chemistrv in determining the gas phase composition and molecular deuteration.,"the disk the solid phase is the dominant component, and because of the importance of grain chemistry in determining the gas phase composition and molecular deuteration."
 We look ab how the molecular abundances vary with heieht above the midplane. and at the radial column density distribution.," We look at how the molecular abundances vary with height above the midplane, and at the radial column density distribution."
 Finally. we compare our results wilh the available observations and use this to determine whether we can sav anvthing about the nature of the desorption processes acting in the disk.," Finally, we compare our results with the available observations and use this to determine whether we can say anything about the nature of the desorption processes acting in the disk."
 We beein bv looking at the vertical abundance distributions at #2 = 250 AU and a time of 1 Mvrs., We begin by looking at the vertical abundance distributions at $R$ = 250 AU and a time of 1 Myrs.
 In this section we concentrate on \lodel D: dillerences in the chemistry that arise because of desorption processes included in the other models are discussed in the next seclions., In this section we concentrate on Model B; differences in the chemistry that arise because of desorption processes included in the other models are discussed in the next sections.
 The fractional abundances as à function of height. z. above (he midplane are displaved in Figure 11..," The fractional abundances as a function of height, $z$, above the midplane are displayed in Figure \ref{fig:frac_250_a}."
 The three laver structure (see Fieure 10)) lound by previous authors (2???) is Clearly seen. with most molecules having low abundances in the midplane (due to freezeout ab z « 50 AU) and in the surlace lavers (due to photodissociation at z > 100 AU). and abundance peaks in the molecular laver. which for this model is between 50 and 95 AU above the midplane.," The three layer structure (see Figure \ref{fig:three}) ) found by previous authors \citep{aikawa02,wl00, ah99}
 is clearly seen, with most molecules having low abundances in the midplane (due to freezeout at $z$ $<$ 50 AU) and in the surface layers (due to photodissociation at $z$ $>$ 100 AU), and abundance peaks in the molecular layer, which for this model is between 50 and 95 AU above the midplane."
 The midplane is not completely devoid of molecules since desorption by CRIT can maintain low levels of CO and No in this region., The midplane is not completely devoid of molecules since desorption by CRH can maintain low levels of CO and $_2$ in this region.
" The reaction of these two species with IL, and its deuterated isotopomers results in the formation of Noll . NoD . and ."," The reaction of these two species with $_3^+$ and its deuterated isotopomers results in the formation of $_2$ $^+$, $_2$ $^+$, $^+$ and $^+$."
" Since the temperature is cold. there is a high level of deuteration of IL,. producing hieh D/II ratios in the daughter molecules."," Since the temperature is cold, there is a high level of deuteration of $_3^+$, producing high D/H ratios in the daughter molecules."
" The presence of CO and No in the gas reduces the deuteration of IL, compared to Model A (which does not have nonthermal desorption) from D, /IL, ~ 70 to a value of ~ 25.", The presence of CO and $_2$ in the gas reduces the deuteration of $_3^+$ compared to Model A (which does not have non–thermal desorption) from $_3^+$ $_3^+$ $\sim$ 70 to a value of $\sim$ 25.
" Deuterated IL, forms from the reaction of ILD with (he previous molecule in the chain ΠΟΙΟ | reacts with ILD to form HD,.", Deuterated $_3^+$ forms from the reaction of HD with the previous molecule in the chain $_2$ $^+$ reacts with HD to form $_2^+$.
" If CO and No are present in the gas. IL, and its deuterated isolopomers are more likely to react with these molecules than with IID. reducing the formation rate of the deuterated molecules (see ο [or details)."," If CO and $_2$ are present in the gas, $_3^+$ and its deuterated isotopomers are more likely to react with these molecules than with HD, reducing the formation rate of the deuterated molecules (see \citet{cd05} for details)."
 The molecular laver begins al the point where thermal desorpüon can remove some molecules. mainly CO. Ne and CIL; from the grains.," The molecular layer begins at the point where thermal desorption can remove some molecules, mainly CO, $_2$ and $_4$ from the grains."
 The upper boundary is set by photodissociation., The upper boundary is set by photodissociation.
 In (he molecular laver most of (he CO cvcles between its gaseous and solid phases. but a small proportion is broken up by reaction with releasing — and O. some of which eo on to," In the molecular layer most of the CO cycles between its gaseous and solid phases, but a small proportion is broken up by reaction with $^+$ releasing $^+$ and O, some of which go on to"
The 5gravitational instability scenario of large-scale5 structure formation relates peculiar velocities of galaxies with their peculiar gravitational accelerations.,The gravitational instability scenario of large-scale structure formation relates peculiar velocities of galaxies with their peculiar gravitational accelerations.
" In linear theory. this relation has a particularly simple form ?:: llere. //,—100km/s/Mpe is the IIubble constant. Q4, is the current value of the eeiccal density parameter of non-relativistic matter. f/(0,,)=(dlnD/dlnz)|.2o (with D being the erowth factor) and oy is the mean matter density of the background."," In linear theory, this relation has a particularly simple form \cite{Pe80}: Here, $H_0=100\,h\,\kms / \mathrm{Mpc}$ is the Hubble constant, $\Omega_\mrm$ is the current value of the cal density parameter of non-relativistic matter, $f(\Omega_\mrm)\equiv\left(\de\ln D\slash\de \ln z\right)\vert_{z=0}$ (with $D$ being the growth factor) and $\rho_b$ is the mean matter density of the background."
" Within ACDM . models with a cosmological constant. the growth parameter f is very well fitted by /(O,,)2OP? (2) and is virtually independent of A. citeLLPBR."," Within $\Lambda$ CDM models with a cosmological constant, the growth parameter $f$ is very well fitted by $f(\Omega_\mrm)\simeq\Omega_\mrm^{0.55}$ \citep{Linder} and is virtually independent of $\Lambda$, \\cite{LLPR}."
.The acceleration vector at a position ‘7 is given by (he integral: where Oy(7)—[py(T)pi]/py is the density contrast of non-relativistic matter at the point T., The acceleration vector at a position $\bmr$ ' is given by the integral: where $\delta_\mrm(\bmr)=\left[\rho_\mrm(\bmr)-\rho_b\right]\slash \rho_b$ is the density contrast of non-relativistic matter at the point $\bmr$.
 llowever. as what we observe are galaxies. we have to assume some relation between their density field and that of matter.," However, as what we observe are galaxies, we have to assume some relation between their density field and that of matter."
 This is usually done via the paradigm: ὃς=504.," This is usually done via the paradigm: $\delta_\mrg=b\, \delta_\mrm$."
 This biasing scheme. valid in linear theory that we use in our whole analysis. neglects the stochasticity. as well as possible scale- ancl galaxy-tvpe dependence in the relation between (he two density fields.," This biasing scheme, valid in linear theory that we use in our whole analysis, neglects the stochasticity, as well as possible scale- and galaxy-type dependence in the relation between the two density fields."
 For more details on possible non-linear biasing. see for example (he review bv ?..," For more details on possible non-linear biasing, see for example the review by \cite{LaSu}."
" Including the biasing relation into ((2)) and using the [act that for a spherical survey [roedr!=0. we get the following expression for the peculiar acceleration: The biasing parameter b is usually combined with the [actor /(£2,) into the parameter 1Ξξ f((Qy)/b. Comparing ((1)) and (3)). we get the proportionality valid in linear theory:"," Including the biasing relation into \ref{eq:g.theor}) ) and using the fact that for a spherical survey $\int\frac{\bmr '-\bmr}{|\bmr '-\bmr |^{3}}\,\de^3 \bmr '=0$, we get the following expression for the peculiar acceleration: The biasing parameter $b$ is usually combined with the factor $f(\Omega_\mrm)$ into the parameter $\beta\equiv f(\Omega_\mrm)\slash b$ Comparing \ref{eq:v.and.g}) ) and \ref{eq:g.galaxies}) ), we get the proportionality valid in linear theory:"
We do not wish to be more quantitative about the enuüssion altitudes at this stage for two reasons.,We do not wish to be more quantitative about the emission altitudes at this stage for two reasons.
 Firstly. ou account of the previously mentioned Iuuitatious of the formula we need a refinement in it.," Firstly, on account of the previously mentioned limitations of the formula we need a refinement in it."
 This should treat the sweepback more precisely using the Deutsch solution for the maguetic field of au oblique rotator and we intend to report on it soon., This should treat the sweepback more precisely using the Deutsch solution for the magnetic field of an oblique rotator and we intend to report on it soon.
 Secondly. the offset translates oulv into a difference between emission altitudes of the two compoucuts.," Secondly, the offset translates only into a difference between emission altitudes of the two components."
 It can determine individual altitudes oulv if at least one of the altitudes (core or cone) is known definitely., It can determine individual altitudes only if at least one of the altitudes (core or cone) is known definitely.
 Tu the foregoing we have ignored the possible contributions to phase offsets by the emuüssiou nechanisuni(s)., In the foregoing we have ignored the possible contributions to phase offsets by the emission mechanism(s).
 Can enmuüssion mechauisuis contribute siguificautly to the phase offset of a pulse coniponeut?, Can emission mechanisms contribute significantly to the phase offset of a pulse component?
 Note firstly that aberration aud the Üufs are first order effects., Note firstly that aberration and the 'mfs' are first order effects.
 As secu in Figs., As seen in Figs.
 1 and 2. these offsets are individually quite large over most of the range of e as well as a.," 1 and 2, these offsets are individually quite large over most of the range of $v$ as well as $\alpha$."
 Simularly from Fies., Similarly from Figs.
" 3 aud £. the same is true even when we deal with their difference. 1οι, Ag."," 3 and 4, the same is true even when we deal with their difference, i.e., $\Delta
 \varphi$."
 Even at sinall (but not extremely simall) values of c. aud. also a. the uct offset is Significant.," Even at small (but not extremely small) values of $v$, and, also $\alpha$, the net offset is significant."
 For the cussion niechauisui to affect this. its contribution should also be equally sjeuificaut.," For the emission mechanism to affect this, its contribution should also be equally significant."
 If this is true then our considerations would need modifications., If this is true then our considerations would need modifications.
 Ou the other haud. irrespective of these dynamical contributions which are also not known. the effects considered by us mist be present.," On the other hand, irrespective of these dynamical contributions which are also not known, the effects considered by us must be present."
 Frou. the foregoing there is ample indication that they suffice for describing the observed core/cone longitude offsets coming into plav due to the cussion mechanisi selecting different altitudes of emission. perhaps stronely varving with period. inclination angle etc.," From the foregoing there is ample indication that they suffice for describing the observed core/cone longitude offsets coming into play due to the emission mechanism selecting different altitudes of emission, perhaps strongly varying with period, inclination angle etc."
" In addition. they are providing constraiuts on the enission mechanisius. ο,οι, the unexpectedly sinall fille factors for the core componcuts for those pulsars in which core leads."," In addition, they are providing constraints on the emission mechanisms, e.g., the unexpectedly small filling factors for the core components for those pulsars in which core leads."
 Now. we briefly comment on multifrequeuncy observations.," Now, we briefly comment on multifrequency observations."
" As the frequency of observation decreases. A,,. reCluains negative but increases in maenitude for some pulsars (c.g.. PSR 0329|51. sec. cg.. Malov aud Suletmanova 1998) while it decreases to such an extent that it changes frou negative to positive for PSR 18521105 (see. e.g... Weisberg et al."," As the frequency of observation decreases, $\Delta_{cc}$ remains negative but increases in magnitude for some pulsars (e.g., PSR 0329+54, see, e.g., Malov and Suleimanova 1998) while it decreases to such an extent that it changes from negative to positive for PSR 1821+05 (see, e.g., Weisberg et al."
 1999)., 1999).
 Since our considerations brine iu filline factor variations. both frequency aud geometry of emission change with the altitude.," Since our considerations bring in filling factor variations, both frequency and geometry of emission change with the altitude."
 Thus. our picture potentially can explain the otherwise perplexing observations and its application will lead to very iusightful coustraiuts on the eniüsson nmechanisu.," Thus, our picture potentially can explain the otherwise perplexing observations and its application will lead to very insightful constraints on the emission mechanism."
 Ig addition.possibilitg.," In addition,."
. This secus to have been observed by Rankin (2001. private conummulcation).," This seems to have been observed by Rankin (2001, private communication)."
are of undamenutal inportauce iu determiune star formation in shallow potential wells typical ol pre-reionization dwarfs.,are of fundamental importance in determining star formation in shallow potential wells typical of pre-reionization dwarfs.
 Figure | (taken from Figure 7 in ROSOS) shows tlat below a [ew 10* M... halos with exactly tje sane dark 1jass can be either dark or luminous (clepencding on the environment).," Figure \ref{fstar} (taken from Figure 7 in RGS08) shows that below a few $10^7$ $_\odot$, halos with exactly the same dark mass can be either dark or luminous (depending on the environment)."
 Thus. feedbacx from non-iotizing and ionizing UV. radiation. mechanical feedback aud chemical enrichment. cai produce two ialos with the same dark mass aid ve‘y different. star ormation elliciencies.," Thus, feedback from non-ionizing and ionizing UV radiation, mechanical feedback and chemical enrichment can produce two halos with the same dark mass and very different star formation efficiencies."
 This :ippears to be the main efect respousible for the observed spread. iu jetallicity. for a given luitosity. in the 1ew dwarfs.," This appears to be the main effect responsible for the observed spread in metallicity for a given luminosity, in the new dwarfs."
" In Figure 5.. we plot the uetallicity as a unction of the mean star foruation elficiency lor the |alo. f/,."," In Figure \ref{Zfs}, we plot the metallicity as a function of the mean star formation efficiency for the halo, $f_*$."
" As expected. simulaed dwarfs with igher values of the star 1jetallicity are the ones with he larger value of f,."," As expected, simulated dwarfs with higher values of the star metallicity are the ones with the larger value of $f_*$."
 However. this effect aloie cannot account for all the observed. scatter of the metallicity as illustrated by the color codi© of simulated chwarls in Figure 3..," However, this effect alone cannot account for all the observed scatter of the metallicity as illustrated by the color coding of simulated dwarfs in Figure \ref{ZL}."
" It appears that the metallicity is οἱ simply proportiot alto f, (otherwise the boundaries between symbols of differeut color would be orizoutal).", It appears that the metallicity is not simply proportional to $f_*$ (otherwise the boundaries between symbols of different color would be horizontal).
" Iustead. or a given value of /,. tte metallicity is larger for fainter cbwarfs."," Instead, for a given value of $f_*$, the metallicity is larger for fainter dwarfs."
 It is uot too rprisiug tha Z is not simply proportional to [ει, It is not too surprising that $Z$ is not simply proportional to $f_*$.
" Even when using a very simple cheical evolution odel. neglecing gas inflows aud assuine’ iststantaneous metal recycling. the mea metallicity of e stars is p'oportional to AL./ALgas ‘ather han f,AL/Map. where Mj, is |e initial value the gas inass avallable for star formation."," Even when using a very simple chemical evolution model, neglecting gas inflows and assuming instantaneous metal recycling, the mean metallicity of the stars is proportional to $M_*/M_{gas}$ rather than $f_*=M_*/M_{bar}$, where $M_{gas}$ is the initial value of the gas mass available for star formation."
 Thus. ZxfiAlbayρα).," Thus, $Z \propto f_*
(M_{bar}/M_{gas})$."
 M feedback elTects reduce e value of iLoosALpay below unity in the SLjiallest. anc lowest |uuluosity primorlal cbwarls. the uetallicity of the stars will )e larger [o ra fixed value of f. as observed in Figure 5 and Figure 3..," If feedback effects reduce the value of $M_{gas}/M_{bar}$ below unity in the smallest and lowest luminosity primordial dwarfs, the metallicity of the stars will be larger for a fixed value of $f_*$, as observed in Figure \ref{Zfs} and Figure \ref{ZL}."
" Tje. reduction of. Mou/Mp,. below ulity ca1 be prexluced. by three effects: the jiucrease of the Jeaus mass of the [GML ove ‘the virial mass 0 “the halo due to rejeatiug (see Figure 6 iu RGS035). heating of the gas via ioniziug radiation from stars within the |alo. aud by iuitiple episodes of star formation with a first burst tha lowers Alga, substantially. out does not p""OC|ice sullicientlv large values of f, aud Z whe1 compared to subsequent bIss."," The reduction of $M_{gas}/M_{bar}$ below unity can be produced by three effects: the increase of the Jeans mass of the IGM over the virial mass of the halo due to reheating (see Figure 6 in RGS08), heating of the gas via ionizing radiation from stars within the halo, and by multiple episodes of star formation with a first burst that lowers $M_{gas}$ substantially, but does not produce sufficiently large values of $f_*$ and $Z$ when compared to subsequent bursts."
 Figure 6 shows [Fe/H] versus t1e surface brightness in the V-band. X ," Figure \ref{ZS} shows [Fe/H] versus the surface brightness in the V-band, $\Sigma_V$ ."
TIe symbols are the sale as in Fie., The symbols are the same as in Fig.
 1 and the soid line s1i0ws the SDSS sensitialy limit., \ref{Kor} and the solid line shows the SDSS sensitivity limit.
 No trend is olserved between metallicity aud vX4. for the siinulated dwarls., No trend is observed between metallicity and $\Sigma_V$ for the simulated dwarfs.
 Observed dwars show less scatter or τμ. relation thau for the Iuiuiuosiy-inetalicity relation in Figwe 3.., Observed dwarfs show less scatter for $\Sigma_V$ -metallicity relation than for the luminosity-metallicity relation in Figure \ref{ZL}.
 There is one dwa with metallicity below [Fe/H]2—2.5: Seeue-1 that has [Fe/H]=—2.8., There is one dwarf with metallicity below $=-2.5$: Segue-1 that has $=-2.8$.
 Sitce he spectral syutesis method used iu ? and ? may not be subjec to the overestimation of metalicities seen with measurements Slug the CA triplet. the lack of «warls with /H]«—)al could be a sign of a chatee in tle IMF at very low [Fe/H].," Since the spectral synthesis method used in \cite{Kirby:08} and \cite{Gehaetal08} may not be subject to the overestimation of metallicities seen with measurements using the CA triplet, the lack of dwarfs with $< -3.0$ could be a sign of a change in the IMF at very low [Fe/H]."
 Finally. we have exzuninec whether there is a depeudeuce of the metallicities on the «instance of the ealaxy from the Milky Way or Audromeda.," Finally, we have examined whether there is a dependence of the metallicities on the distance of the galaxy from the Milky Way or Andromeda."
 For the new Milky Way chwarls there is slieht trend oL higher metallicities at smaller Galactoceutric distauces: however. the upward {θέ is «oniated by Ursa Major IE aud Coma Ber..," For the new Milky Way dwarfs there is slight trend of higher metallicities at smaller Galactocentric distances; however, the upward trend is dominated by Ursa Major II and Coma Ber.,"
 both of ofwhich show evidence for tidal disruption., both of ofwhich show evidence for tidal disruption.
 Fo: the M31 dwarfs. no trend was observed.," For the M31 dwarfs, no trend was observed."
Piro et al,Piro et al.
/s paper on GRB 011121 at the end of Feb. and ας E. Jiang. L. J. ou. Z. Li and IH. TE. Ma for kind. help.,"'s paper on GRB 011121 at the end of Feb, and G. F. Jiang, L. J. Gou, Z. Li and H. T. Ma for kind help."
 We also appreciate the referees for their helpful comments and the third referee. for her/his great help., We also appreciate the referees for their helpful comments and the third referee for her/his great help.
 This work is supported. by the National Natural Science. Foundation (erants 10225314 and 10233010) of China. and the National 973 Project on Fundamental Researches of China (NABRSE 6119090754).," This work is supported by the National Natural Science Foundation (grants 10225314 and 10233010) of China, and the National 973 Project on Fundamental Researches of China (NKBRSF G19990754)."
In this paper we studied the process of the momentun transfer between cold. clouds ejected from the disc. and the corona of disc galaxies Like the Milkv Way. which has a significant impact on the kinematics of both the coronal and the cold extraplanar gas.,"In this paper we studied the process of the momentum transfer between cold clouds ejected from the disc and the corona of disc galaxies like the Milky Way, which has a significant impact on the kinematics of both the coronal and the cold extra–planar gas."
 In. particular. we presented: two-dimensional hydrodynamical simulations of cloud-corona interaction.," In particular, we presented two-dimensional hydrodynamical simulations of cloud-corona interaction."
 In the more realistic simulations the eas is allowed to cool radiatively. but. for comparison. we also considered the corresponding adiabatic cases.," In the more realistic simulations the gas is allowed to cool radiatively, but, for comparison, we also considered the corresponding adiabatic cases."
 The simulations are similar. but significantly improved. with respect to those carried. out by 2: here we explored. a range of relative speeds between the cloud and the hot medium (50200kms +). we followed the evolution of the system for longer time (60 Myr). we treated the metallicity as a dynamical variable (assigning initially to the cloud a metallicity ten times higher than that of the ambient. eas) and we accounted for the variation with temperature of the μιαν molecular weight yr.," The simulations are similar, but significantly improved, with respect to those carried out by \citetalias{Marinacci10}: here we explored a range of relative speeds between the cloud and the hot medium $50-200\kms$ ), we followed the evolution of the system for longer time $60\Myr$ ), we treated the metallicity as a dynamical variable (assigning initially to the cloud a metallicity ten times higher than that of the ambient gas) and we accounted for the variation with temperature of the fluid's molecular weight $\mu$."
 From the simulations we draw the following conclusions., From the simulations we draw the following conclusions.
 The process studied here has an important consequence on the kinematies of the corona., The process studied here has an important consequence on the kinematics of the corona.
 The corona cannot be static. otherwise the momentum transfer from. the cold clouds. would. be very ellective. but it cannot be rotating at a speed. close to the rotation speed. of the disc. because at low relative disc-corona (and thus cloud-corona) velocity the cooling hampers any momentum transfer.," The corona cannot be static, otherwise the momentum transfer from the cold clouds would be very effective, but it cannot be rotating at a speed close to the rotation speed of the disc, because at low relative disc-corona (and thus cloud-corona) velocity the cooling hampers any momentum transfer."
 In a galaxy like the Alilky Way we expect the corona to lag. in the inner regions. by I00knms with respect to the cold. disc at oo2kpe.," In a galaxy like the Milky Way we expect the corona to lag, in the inner regions, by $\sim 100
\kms$ with respect to the cold disc at $z\simeq2\kpc$."
 Such a spinning corona must be characterised by a significantIv llattened density cistribution: based on simple isothermal models. we estimate values of the axis ratios of its isodensity surfaces in the range 0.5—0.7.," Such a spinning corona must be characterised by a significantly flattened density distribution: based on simple isothermal models, we estimate values of the axis ratios of its isodensity surfaces in the range $0.5-0.7$."
 Due to interaction with the corona. fountain clouds eain mass and decelerate: based on our calculations. this cllect can account for zThmsΚροτ of the vertical eracient in rotational speed of the extraplanar gas in disc ealaxies like the Milky Way.," Due to interaction with the corona, fountain clouds gain mass and decelerate: based on our calculations, this effect can account for $\approx-7 \kms\kpc^{-1}$ of the vertical gradient in rotational speed of the extra–planar gas in disc galaxies like the Milky Way."
" Such a value can be sullicient to reconcile purely ballistic fountain models with the observed kinematics of the eextraplanar gas. giving further support to the ""fountain plus accretion” model of ?.."," Such a value can be sufficient to reconcile purely ballistic fountain models with the observed kinematics of the extra–planar gas, giving further support to the “fountain plus accretion” model of \citetalias{FraternaliB08}."
 1n conclusion. we think that the presented. results contribute quantitatively to refine an emerging consistent scenario in which galactic fountains. cold extra.planar gas and galactic aare strictly interlaced.," In conclusion, we think that the presented results contribute quantitatively to refine an emerging consistent scenario in which galactic fountains, cold extra–planar gas and galactic are strictly interlaced."
 “Phis scenario makes specific predictions that can be tested in the future with soft. rav observations detecting the elusive galactic aand with improved nuneasures of the cold extraplanar eas., This scenario makes specific predictions that can be tested in the future with soft X-ray observations detecting the elusive galactic and with improved measures of the cold extra–planar gas.
" From the theoretical point of view. the next step will be the extension to three dimensions of the two-dimensional calculations here presented. and ideally one would. like to realise ""elobal simulations. in which the entire galaxy is modelled. with supernova-driven fountain clouds orbiting through the hot coronal medium under the clleet of the ealactic gravitational Ποια."," From the theoretical point of view, the next step will be the extension to three dimensions of the two-dimensional calculations here presented, and ideally one would like to realise “global” simulations, in which the entire galaxy is modelled, with supernova-driven fountain clouds orbiting through the hot coronal medium under the effect of the galactic gravitational field."
In (his section. we use the MICAIC technique to constrain the model parameters.,"In this section, we use the MCMC technique to constrain the model parameters."
 The MCAIC method is well suitable for high dimensional parameter space investigation., The MCMC method is well suitable for high dimensional parameter space investigation.
 The Metropolis-Lastines algorithm is used when sampling the model parameters., The Metropolis-Hastings algorithm is used when sampling the model parameters.
 The probability density distributions of the model parameters can also be simply approximated by the number densitv of the sample points., The probability density distributions of the model parameters can also be simply approximated by the number density of the sample points.
 A brief introduction to the basic. procedure of the MCAIC sampling can be found in Fanetal.(2010b).., A brief introduction to the basic procedure of the MCMC sampling can be found in \cite{2010A&A...517L...4F}.
 For more details about the MCMC method. please reler to Neal(1993):Gamerman(1997):Mackay.(2003)..," For more details about the MCMC method, please refer to \cite{Neal1993,Gamerman1997,
2003itil.book.....M}."
 We also discuss implications of model parameters from the best fits to multi-wavelength data of SNR RX J1713.7-3946 lor three scenarios of the 5-rav emission., We also discuss implications of model parameters from the best fits to multi-wavelength data of SNR RX J1713.7-3946 for three scenarios of the $\gamma$ -ray emission.
 In all these scenarios. the radio to X-ray emissions are generated through svnchrotron of relativistic electrons.," In all these scenarios, the radio to X-ray emissions are generated through synchrotron of relativistic electrons."
 The hieh energy 5-ravs are produced with different mechanisms., The high energy $\gamma$ -rays are produced with different mechanisms.
 The basic physical parameters ol SNR. RN J1713.7-3946 are adopted as: Age Ti.21600 vr. Distance dzz1 kpe. and Radius 2210 pe (Wangetal.1997).. and we assume a uniform emission sphere with a radius of R in deriving related quantities.," The basic physical parameters of SNR RX J1713.7-3946 are adopted as: Age $T_{\rm life}\approx 1600$ yr, Distance $d\approx 1$ kpc, and Radius $R\approx 10$ pc \citep{1997A&A...318L..59W}, and we assume a uniform emission sphere with a radius of $R$ in deriving related quantities."
 Although the errors of the Fermi data are large. we still include (hese data in the spectral fits (Abdoetal.2011)..," Although the errors of the Fermi data are large, we still include these data in the spectral fits \citep{fermi:rxj1713}."
 The procedure described in (his paper can be applied to future observations with improved data to evaluate different emission models., The procedure described in this paper can be applied to future observations with improved data to evaluate different emission models.
 In the leptonic scenario the ταν emission is produced through inverse Compton (IC) scaltering of energetic electrons olf the background radiation field. including the interstellar inlrared. optical radiation. and the cosmic microwave background (CMD).," In the leptonic scenario the $\gamma$ -ray emission is produced through inverse Compton (IC) scattering of energetic electrons off the background radiation field, including the interstellar infrared, optical radiation, and the cosmic microwave background (CMB)."
" The energy spectrum of accelerated electrons is prescribed as ΕΕ)xEoexp[οn|. where E. a,. E; ave the electron energy. power law spectral index. high-energy cutoff energy. respectively. and 9, describes the sharpness of this cutoff."," The energy spectrum of accelerated electrons is prescribed as $F_e(E)\propto E^{-\alpha_e}\exp
\left[-(E/E_c^e)^{\delta_e}\right]$, where $E$, $\alpha_e$ , $E_c^e$ are the electron energy, power law spectral index, high-energy cutoff energy, respectively, and $\delta_e$ describes the sharpness of this cutoff."
 The normalization is given through the total energy ol electrons above 1 GeV. Wi.," The normalization is given through the total energy of electrons above 1 GeV, $W_e$."
 The svnchrotron radiation also depends on (he magnetic field strength 2. The interstellar radiation field (ISRE) other than the CMD max be important for the calculation of IC οταν spectrum., The synchrotron radiation also depends on the magnetic field strength $B$ The interstellar radiation field (ISRF) other than the CMB may be important for the calculation of IC $\gamma$ -ray spectrum.
 The inclusion of ISRE has been proposed to improve the fit tothe HESS data (Porteretal.2006).., The inclusion of ISRF has been proposed to improve the fit to the $HESS$ data \citep{2006ApJ...648L..29P}.
 IIowever. given the new data of X-ray by Suzaku aud TeV 5-rav by HEESS. i was shown that only if the intensity of ISRE is artificially boosted by more (han one order of magnitude. the goocdness-ol-fit can be improved significantly 2009a).," However, given the new data of X-ray by $Suzaku$ and TeV $\gamma$ -ray by $HESS$, it was shown that only if the intensity of ISRF is artificially boosted by more than one order of magnitude, the goodness-of-fit can be improved significantly\citep{2008ApJ...685..988T,2009MNRAS.392..240M}."
.In this work the ISRF is adopted as that given by, .In this work the ISRF is adopted as that given by
These muubers clearly show that (P; in BLRGs is significautly lower (by ~5o) than that iu Sevtert 1s.,These numbers clearly show that $\langle \Gamma\rangle$ in BLRGs is significantly lower (by $\sim 5\sigma$ ) than that in Seyfert 1s.
 The probability that the mean of the above joiut sample of BLRGs equals 1.95 G.c.. the mean of Sevtert Is of Nandra Pounds 1991) is 2.3«10 (οίαλος frou the Student £ distribution|.," The probability that the mean of the above joint sample of BLRGs equals 1.95 (i.e., the mean of Seyfert 1s of Nandra Pounds 1994) is $2.3\times 10^{-5}$ (obtained from the Student $t$ distribution)."
 Also. we note that 15 best-fit values of P iu the above 3 samples of BLERGs are less than (T? of Nandra Pounds (199D.," Also, we note that 15 best-fit values of $\Gamma$ in the above 3 samples of BLRGs are less than $\langle \Gamma \rangle$ of Nandra Pounds (1994)."
 If Ts of BLRCs were sunpled from the same distribution as that of Seyfert Ls. the probability of such au eveut would beouly 2P34107.," If $\Gamma$ s of BLRGs were sampled from the same distribution as that of Seyfert 1s, the probability of such an event would be only $2^{-15}\simeq 3\times 10^{-5}$."
" Furtheriiore, we can tes the hypothesis that the joiut sample of BLRGs has the mean of 1.95 the standard deviation of 0.15 (as in the saniple of Naudra Pounds 1991). for which we obtain the probability of L7«10.7 (based ou the normal distribution)."," Furthermore, we can test the hypothesis that the joint sample of BLRGs has the mean of 1.95 the standard deviation of 0.15 (as in the sample of Nandra Pounds 1994), for which we obtain the probability of $4.7\times 10^{-8}$ (based on the normal distribution)."
 We note that E00 stated that their distribution of D did not «differ from that of Sevfert 1s fromGinya., We note that E00 stated that their distribution of $\Gamma$ did not differ from that of Seyfert 1s from.
. However. he probability that their values of Ε (1.78. 1.90. 1.72. 1.92) belong to the distribution with (D?=15 and op=0.15 is oul54.," However, the probability that their values of $\Gamma$ (1.78, 1.90, 1.72, 1.92) belong to the distribution with $\langle\Gamma \rangle=1.95$ and $\sigma_\Gamma= 0.15$ is only."
 Furthermore. we should note the difference of ATz0.1 between ‘and other X-ray iustruineuts iu fits to he Crab as well as other sources (e.g. Cüerlisshi et 11999: Done. Madejski Zvveki 2000). which. if taken into account. would make the difference in (TD) between BLRCs aud Seyfert 185 even more pronounced. reducing the above probability to LF«107.," Furthermore, we should note the difference of $\Delta\Gamma \simeq 0.1$ between and other X-ray instruments in fits to the Crab as well as other sources (e.g., Gierlińsski et 1999; Done, Madejski Żyycki 2000), which, if taken into account, would make the difference in $\langle \Gamma \rangle$ between BLRGs and Seyfert 1s even more pronounced, reducing the above probability to $1.7\times 10^{-3}$."
 Ou the other hand. the distribution of P iu BLRGs is fron that in Sevfert Ls.," On the other hand, the distribution of $\Gamma$ in BLRGs is from that in Seyfert 1s."
 This can be noted by considering the iutriusic dispersion. op. of the values of TP iu the samples (given bv their standard cdeviatious).," This can be noted by considering the intrinsic dispersion, $\sigma_\Gamma$, of the values of $\Gamma$ in the samples (given by their standard deviations)."
" For the BLRCs samples fromGinga.. audAXTE.. op=0.18. 0.05. aud 0,10. respectively. while op of their sum equals 0.11."," For the BLRGs samples from, and, $\sigma_\Gamma=0.18$, 0.08, and 0.10, respectively, while $\sigma_\Gamma$ of their sum equals 0.14."
 The dispersion iu the Sevfert-1 samples from. and CNandia Pounds 199 I:Matt 2000) is gp=0.15 and 0.20. respectively.," The dispersion in the Seyfert-1 samples from and (Nandra Pounds 1994; Matt 2000) is $\sigma_\Gamma=0.15$ and 0.20, respectively."
 Given the values of (D; listed above. the distributions for BLRGs aud Sevtert Is significantly overlap Gvhich. we stress again. docs not uuply that the 2 distributions are inclistinguislhalle).," Given the values of $\langle
\Gamma \rangle$ listed above, the distributions for BLRGs and Seyfert 1s significantly overlap (which, we stress again, does not imply that the 2 distributions are indistinguishable)."
 Furthermore. BLRGs obey an overall correlation between these D iud Q seen in Sevferts (Zdziarski. Lubinsslki Suuth 1999. hereafter Z99).," Furthermore, BLRGs obey an overall correlation between these $\Gamma$ and $\Omega$ seen in Seyferts (Zdziarski, Lubińsski Smith 1999, hereafter Z99)."
 This is illustrated in Figure 5. (where we can also note the AT~0.1 offset between the yesults and those of and X))., This is illustrated in Figure \ref{f:reflection} (where we can also note the $\Delta\Gamma\sim 0.1$ offset between the results and those of and ).
 We sce that BLRCs (includiug 3C 120) suuplv occupy a lower left part of the broader D-O parineter space occupied by radio-quiet. Sevterts., We see that BLRGs (including 3C 120) simply occupy a lower left part of the broader $\Gamma$ $\Omega$ parameter space occupied by radio-quiet Seyferts.
 Also. the average streneth of the broad component of the Fe Ίνα liue iu racdio-quiet Sevterts with P close to that of 3€ 120 observed by (Lubiüsski Zdziarski 2001) corresponds to the streneth of Compton reflection cousistcu within error bars with that observed in 3€ 120.," Also, the average strength of the broad component of the Fe $\alpha$ line in radio-quiet Seyferts with $\Gamma$ close to that of 3C 120 observed by (Lubińsski Zdziarski 2001) corresponds to the strength of Compton reflection consistent within error bars with that observed in 3C 120."
 This Προς a general similarity between the Ttwsical couditious iu the X-ray sources of Sevterts >xl. BLRGs Gucluding 3€ 120). but with some aukuowu paraueter determine (D. 09) in BLRGs offset with respect to that in radio-«quiet Sevferts.," This implies a general similarity between the physical conditions in the X-ray sources of Seyferts and BLRGs (including 3C 120), but with some unknown parameter determining $\Gamma$, $\Omega$ ) in BLRGs offset with respect to that in radio-quiet Seyferts."
 We discuss this issue in the 86.3. below., We discuss this issue in the \ref{s:geometry} below.
 The overall similarity between the X-ray properties of BLRGs aud radio-quiet Sevterts is further confirmed by the variability pattern of 3€ 120 refftvarbb). where the N-rav spectrum becomes softer with the increasing N-rav flux.," The overall similarity between the X-ray properties of BLRGs and radio-quiet Seyferts is further confirmed by the variability pattern of 3C 120 \\ref{f:var}b b), where the X-ray spectrum becomes softer with the increasing X-ray flux."
 This pattern is the same as that often seen in racdio-quiet Sevtert ls. owe. in NGC 55I8 (Alaedziarz et 11998) and NGC 1151 (Yaqoob Warwick 1991).," This pattern is the same as that often seen in radio-quiet Seyfert 1s, e.g., in NGC 5548 (Magdziarz et 1998) and NGC 4151 (Yaqoob Warwick 1991)."
 We also uote that black-hole binaries im the hard state show a similar behavior (e... Zdziuski. Wen Paciesas 2001).," We also note that black-hole binaries in the hard state show a similar behavior (e.g., Zdziarski, Wen Paciesas 2001)."
 Clearly. the dominant radiative process respousible Or N-rav contiuua of many classes of compact objects is Compton upscatteriug of soft photous w enersetie electrons.," Clearly, the dominant radiative process responsible for X-ray continua of many classes of compact objects is Compton upscattering of soft photons by energetic electrons."
 Less clear issues are the orm of the electron distribution aud the source of he seed soft photons., Less clear issues are the form of the electron distribution and the source of the seed soft photons.
 Iu the case of racdio-quiet Sevtert ls. the form of the soft 5-—raw spectra roni OSSE rules out distributiousdomuüuated by 1j01-thormial electrous (Coudel et 11996).," In the case of radio-quiet Seyfert 1s, the form of the soft -ray spectra from OSSE rules out distributionsdominated by non-thermal electrons (Gondek et 1996)."
 The strongest evidence for an olectron distribution close to Maxwellian with AL~50 100 keV and το Lin Sevferts comes from the average OSSE spectrum) of NCC 1151 (Johuson et 11997), The strongest evidence for an electron distribution close to Maxwellian with $kT\sim 50$ –100 keV and $\tau\sim 1$ in Seyferts comes from the average OSSE spectrum of NGC 4151 (Johnson et 1997)
"Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc. The EVN is a joint facility of European, Chinese, South African and other radio astronomy institutes funded by their national research councils.","Observatory is a facility of the National Science Foundation operated under cooperative agreement by Associated Universities, Inc. The EVN is a joint facility of European, Chinese, South African and other radio astronomy institutes funded by their national research councils."
 The WSRT is operated by ASTRON with support from the Netherlands Foundation for Scientific Research., The WSRT is operated by ASTRON with support from the Netherlands Foundation for Scientific Research.
vorlices in the are under (he same conditions as those considered in current post-glitch moclels.,vortices in the are under the same conditions as those considered in current post-glitch models.
" Although the model is quite schematic. it contains a plausible mechanism [or storing ancl releasing vorlicily, as actually confirmed by parallel dynamical simulations (llaskellοἱal.2011)."," Although the model is quite schematic, it contains a plausible mechanism for storing and releasing vorticity, as actually confirmed by parallel dynamical simulations \citep{hps11}."
" In particular. the model allows to calculate.directly the angular momentum £,.(7) of the vortex sheet atv."," In particular, the model allows to calculate the angular momentum $L_v(x)$ of the vortex sheet at $x$ ."
" In Figure 4 we show the reduction of angular momentum. μία). /L,C0). when uniformly. distributed. vorlicity contained within -r is accumulated in a sheet at ar: at the peak. .r,,. the reduction is of order LO7."," In Figure \ref{fig4} we show the reduction of angular momentum, $\ell_v(x)=L_v(x)/L_v(0)$ , when uniformly distributed vorticity contained within $x$ is accumulated in a sheet at $x$; at the peak, $x_m$, the reduction is of order $10^{-3}$."
" For comparison. we also show the significantly different results for a uniform-densityv. evlindrical or spherical stir: we see how spherical sviumetry and realistic density profile are crucial to obtain the correct order ol magnitude of 6,r)."," For comparison, we also show the significantly different results for a uniform-density, cylindrical or spherical star; we see how spherical symmetry and realistic density profile are crucial to obtain the correct order of magnitude of $\ell_v(x)$ ."
" The angular momentum stored during M, and released at the glitch. Z4. can be caleulated [rom the number of vortices removed [rom the interior and aecumulated at 7,,. namely ονΑρα)=2xR20?uuurs."," The angular momentum stored during $\Delta t_{\rm gl}$ and released at the glitch, $\Delta L_{\rm gl}$, can be calculated from the number of vortices removed from the interior and accumulated at $x_m$, namely $\Delta N_v(x_m)=2\pi R_s^2x_m^2\omega_{\rm max}/\kappa$."
 We find (cf., We find (cf.
 Paper I) with an effective moment of inertia where The glitch rise-time is very short. τω«40 s (Dodsonetal.2002): we introduce a new parameter. Σω which globally describes the of vorticitycoupled to the normal crust on timescales of order 744 (the steady-state coupled Traction. corresponding to long timescales and to pre-glitch conditions. is 34.= 1).," Paper I) with an effective moment of inertia where The glitch rise-time is very short, $\tau_{\rm gl}<40$ s \citep{dod02}; we introduce a new parameter, $Y_{\rm gl}$, which globally describes the of vorticity to the normal crust on timescales of order $\tau_{\rm gl}$ (the steady-state coupled fraction, corresponding to long timescales and to pre-glitch conditions, is $Y_{\infty}=1$ )."
 The value of Y; depends on the detailed short-time dvnamics of the vorticitv: in order to get an estimate of the observables. only (his quantity is needed.," The value of $Y_{\rm gl}$ depends on the detailed short-time dynamics of the vorticity; in order to get an estimate of the observables, only this quantity is needed."
 From angular momentum conservationand variationof the crust equation of motion we findthe elitch jump parameters, From angular momentum conservationand variationof the crust equation of motion we findthe glitch jump parameters
temperature.,temperature.
 Tf the spectral index 9 due to Calactic eniüssiou is coustaut in some area. then the following relations hold:," If the spectral index $\beta$ due to Galactic emission is constant in some area, then the following relations hold:"
cau significantly alter the radial abundauce profiles calculated. by the mocel.,can significantly alter the radial abundance profiles calculated by the model.
 However. undersuch circumstances the main couclusious ol this study remain the same.," However, under such circumstances the main conclusions of this study remain the same."
 The location of the 15” deusity-enhauced shell has been fixed in order to best reproduce the observational data. which may be considered a contrivance olthe mocel.," The location of the $15''$ density-enhanced shell has been fixed in order to best reproduce the observational data, which may be considered a contrivance of the model."
 However. in Figure 8 of Maurou&Huggius (2000).. this cau be icentilied as the radius at whieh the first. distinct. dust shells occur.," However, in Figure 8 of \citet{mau00}, this can be identified as the radius at which the first distinct dust shells occur."
 To permit comparison between this mode and previous models published in the literature (which do not include deusityv-enlanced gas and dust shells). we have also rut the mode without any cdensitv-enbauced shells d -0) aud also with. the cdensity-enhancec dust shells of (Brown&Millar2003) to produce the data shown inL.," To permit comparison between this model and previous models published in the literature (which do not include density-enhanced gas and dust shells), we have also run the model without any density-enhanced shells $\beta=0$ ) and also with the density-enhanced dust shells of \citep{bro03} to produce the data shown in."
. To [acilitate comparison] of the inodel ‘esults with observed mapsal of molecularD nicrowave emission⋅⋅ line⋅ flux. the rotational excitation of some of the molecules of interest ins been calculated as a function⋅ of⋅ raclits⋅ sine a modified version of the munuline ∢∙∩∐↕↥↽≻⋯≺↵↕⋅∢∙⋯⇂↩↸≺⇂≺↲↜∖∢∙↕⋅∣≻≺↲≼≻∖⇁∙∣⋃↜∖⊓⋜↕∐∩∐↕≺↵↕∣⋅al.199 1).," To facilitate comparison of the model results with observed maps of molecular microwave emission line flux, the rotational excitation of some of the molecules of interest has been calculated as a function of radius using a modified version of the mmline computer code \citep[described by][]{jus94}."
. The central stellar radius was taken o be 3.1x10 em and the temperature 2650 ⊾⊁↸⋡⊓↵↓⋅∖⇁≺↵≺⇂⋅↕⋅∩⊔↥∖↓≺↵⇩⊳∖∐∢⋅∏⊆, The central stellar radius was taken to be $\times10^{13}$ cm and the temperature 2650 K \citep[derived from][]{men01}.
∩∖⇁≺↵↕⋜↕↥⋅⊇∪∪↽⊔⋅−⋅ The dust opacity was taken from Figure 6 of Memshehikovetal.(2001)., The dust opacity was taken from Figure 6 of \citet{men01}.
. Due to the ack of published. collision: excitation rates aud vibrational trausition strengths.B only a ‘ough estimate of the rotational excitation is possible for most molecules.," Due to the lack of published collisional excitation rates and vibrational transition strengths, only a rough estimate of the rotational excitation is possible for most molecules."
 For CoH aud CoH the rates from Flower(1999) are used. and for CjH. CjH. ο and CG5H he HC3N rates from Creen&Chapman(1978) are used.," For $_2$ H and $_2$ $^-$ the $^+$ rates from \citet{flo99} are used, and for $_4$ H, $_4$ $^-$, $_6$ H and $_6$ $^-$ the $_3$ N rates from \citet{gre78} are used."
 For CoH aud CoH the rates lave been extrapolated up to J=οἱ., For $_6$ H and $_6$ $^-$ the rates have been extrapolated up to $J=31$.
 The collisional rates used are for closed electronic-shell specieslap] so the spectroscopic.E structure. ol the correspoucling (closed. electronic-shell) anions has been used for the calculations of the excitation of the open-shell CoH. CH aud CGH radicals.," The collisional rates used are for closed electronic-shell species so the spectroscopic structure of the corresponding (closed electronic-shell) anions has been used for the calculations of the excitation of the open-shell $_2$ H, $_4$ H and $_6$ H radicals."
 This approximation is reasonable because the structure of these hydrocarbon anious aud neutrals are very similar., This approximation is reasonable because the structure of these hydrocarbon anions and neutrals are very similar.
" Only the ""Hy states have. been cousidered for the ueutval hydrocarbons: the population of the ?1L;; states are not expected to sieutficautly affect. the relative populations of the states olf interest here.", Only the $^2\Pi_{1/2}$ states have been considered for the neutral hydrocarbons; the population of the $^2\Pi_{3/2}$ states are not expected to significantly affect the relative populations of the states of interest here.
 Rotational Einstein A coefficients were calculated using dipole-moments from Woon—(1995) and Blauksbyetal.(2001) for the neutrals and the anions. respectively.," Rotational Einstein $A$ coefficients were calculated using dipole-moments from \citet{woo95} and \citet{bla01} for the neutrals and the anions, respectively."
 I. pumpiug of rotational levels has been caleulated through consideration of the radiative excitation of a stugle vibrational state ~10 jun above the ground state (ef.1981)., IR pumping of rotational levels has been calculated through consideration of the radiative excitation of a single vibrational state $\sim10$ $\mu$ m above the ground state \citep[\emph{cf}.
 CoH has been calculated. to have a strong vibrational transition (4=0.6 1) at 12.5 pan (Tarroni&Carter20010)... which we assume to also oceur ln⋅↿ CoH .," $_2$ H has been calculated to have a strong vibrational transition $A=0.6$ $^{-1}$ ) at 12.5 $\mu$ m \citep{tar04}, which we assume to also occur in $_2$ $^-$."
 The vibrational⋅ spectra ofopa CyH. CoH aud. thelr. associated. anious⋅ are less well kuown.," The vibrational spectra of $_4$ H, $_6$ H and their associated anions are less well known."
 For⊲ these speciesME a trausition⋅⋅ has been assumed to occur at 4912.5x pm with. 4=.1s 1ον, For these species a transition has been assumed to occur at 12.5 $\mu$ m with $A=1$ $^{-1}$.
 IR pumping:ον has a ⋅⋅⋅ ⊳∖↥∑≟∐∐∐∙⋜↕∐↕≺↵∐≺↵∢∙↕∩∐↕∐≺↵⋯∩↥≺↲∢∙⋃⋜∐⋅≺↵⊸∖∢∙↥⋜↕⋃∩∐⋅. ⋅⋅ but changes in⋅ the vibrational⋅ transition⋅⋅ waveleneths aud Eiusteiu A coellicieuts by up . ↕∩⋜↕∐∩↓⋅≺⇂≺↵↓⋅∩↥∐↕⋜↕∑≟∐∐⋯⇂≺↵≺⇂∩↕∩↕⊳∖↓∑≟∐∐∢∙⋜⋃∐↥⊽∖⇁. ⋅⋅⋅ affect. the results of the present. study.," IR pumping has a significant effect on the molecular excitation, but changes in the vibrational transition wavelengths and Einstein $A$ coefficients by up to an order of magnitude do not significantly affect the results of the present study."
 Το assess the impact of errors in the collisional excitation rates on the calculated molecular emission profiles. the rates were varied by au order of⋅ maguitude⋅ either⋅ way.," To assess the impact of errors in the collisional excitation rates on the calculated molecular emission profiles, the rates were varied by an order of magnitude either way."
 The overall features of the emission profiles remained the salle., The overall features of the emission profiles remained the same.
 The upper pauel of shows the, The upper panel of shows the
P and ie represent the equilibrium pressure and enthalpy density. respectively. and ey. and ς represent the speed of sound. shear and bulk viscosities.,"$P$ and $w$ represent the equilibrium pressure and enthalpy density, respectively, and $v_s$, $\eta$ and $\zeta$ represent the speed of sound, shear and bulk viscosities."
 Of course (he energv-momentunm tensor obevs (he conservation equation. will the Greek indices ji7 running over all four coordinates. — 0.," Of course the energy-momentum tensor obeys the conservation equation, with the Greek indices $\mu,\nu$ running over all four coordinates, = 0."
 These linearized hydrodsynamic equations acit (wo normal modes. corresponding to whether the momentum density fluctuations are transverse or longitudinal to the fluid flow.," These linearized hydrodynamic equations admit two normal modes, corresponding to whether the momentum density fluctuations are transverse or longitudinal to the fluid flow."
 Transverse f[Inctuations lead to the which is the one considered in this paper., Transverse fluctuations lead to the which is the one considered in this paper.
 There is also the (from longitudinal moment [Iuctuations). and the (ii the presence ol a conserved current): (hese modes are not considered in (his work.," There is also the (from longitudinal momentum fluctuations), and the (in the presence of a conserved current); these modes are not considered in this work."
 For a more complete introduction to relativistic hydrodynamics see |11].., For a more complete introduction to relativistic hydrodynamics see \cite{Kovtun2}.
 IIycdrodsnamic [Iuctuations in (he gauge theory plasma are dual to perturbations of a specified eravitational background in the dual theory., Hydrodynamic fluctuations in the gauge theory plasma are dual to perturbations of a specified gravitational background in the dual theory.
 Throughout Chis paper we use an effective 5-dimensional metric., Throughout this paper we use an effective 5-dimensional metric.
 In. principle. one can view this metric as being derived by cimensionally reducing some theory in a higher dimension.," In principle, one can view this metric as being derived by dimensionally reducing some theory in a higher dimension."
 Our results are easily generalizable to different numbers of dimensions in the same way as the results of [9]..., Our results are easily generalizable to different numbers of dimensions in the same way as the results of \cite{stretched}.
 Our convention is to use /.a.uy.2 to denote the usual four space-time coordinates: the coordinate r denotes the extra dimension.," Our convention is to use $t,x,y,z$ to denote the usual four space-time coordinates; the coordinate $r$ denotes the extra dimension."
 In (he case of the shear mode. one only needs to consider metric perturbations ui—UgoFy.," In the case of the shear mode, one only needs to consider metric perturbations $g_{\mu \nu} \rightarrow g_{\mu \nu} + h_{\mu \nu}$."
 The components of the perturbation των can be decomposed into three irreducible sets: the field equations for each set decouple Lom the others., The components of the perturbation $h_{\mu \nu}$ can be decomposed into three irreducible sets; the field equations for each set decouple from the others.
 The sets are found by classifvine the perturbations under. O(2) rotations [4]: see also [8].., The sets are found by classifying the perturbations under $O(2)$ rotations \cite{hydroI}; see also \cite{Mas}.
" The set of gravitational perturbations dual to the hvdrodynanmic shear mode have ρω. h,.. and h,,yr all nonzero with all other components of h,,, vanishing."," The set of gravitational perturbations dual to the hydrodynamic shear mode have $h_{y0}$, $h_{yz}$, and $h_{yr}$ all nonzero with all other components of $h_{\mu \nu}$ vanishing."
 We assume a standard gauge choice Πω which leaves us with only two nonvanishing components of μυ., We assume a standard gauge choice $h_{\mu r} = 0$ which leaves us with only two nonvanishing components of $h_{\mu \nu}$.
 The hydroclvuamic transport coefficients are found by solving the linearized Einstein equations for the perturbations with appropriate boundary conditions., The hydrodynamic transport coefficients are found by solving the linearized Einstein equations for the perturbations with appropriate boundary conditions.
 The resulting dispersion relation for (hese gravitational perturbations can be compared with the expected dispersion relation from the boundary (gauge) hvdrodynamies., The resulting dispersion relation for these gravitational perturbations can be compared with the expected dispersion relation from the boundary (gauge) hydrodynamics.
 In first order hydrodynamics. this method has been used to calculate the speed of sound. and the shear ancl bulk viscosities in a wide variety of gravity. duals.," In first order hydrodynamics, this method has been used to calculate the speed of sound and the shear and bulk viscosities in a wide variety of gravity duals."
 In [9].. the authors also derive a general formula for the shear viscosity which depends only on components of the dual metric: (his formula is applicable for a large class of gravitational backgrounds.," In \cite{stretched}, the authors also derive a general formula for the shear viscosity which depends only on components of the dual metric; this formula is applicable for a large class of gravitational backgrounds."
a result of resonant interactions with the disk.,a result of resonant interactions with the disk.
 The eap-clearing timescale can be estimated by assundue that it would occur on the viscous timescale to spread across the scale height in the protoplanetary disk (?).., The gap-clearing timescale can be estimated by assuming that it would occur on the viscous timescale to spread across the scale height in the protoplanetary disk \citep{sasaki10}.
 The actual timescale for a growing plauct to open a gap is likely longer than this estimate. but it naust be substantially shorter than the planets accretion timescale in order to limit the final mass of the planct.," The actual timescale for a growing planet to open a gap is likely longer than this estimate, but it must be substantially shorter than the planet's accretion timescale in order to limit the final mass of the planet."
 A reasonable estimate for the gap opening timescale would be to assume a median value of Των222.5«10? vx., A reasonable estimate for the gap opening timescale would be to assume a median value of $\tau_{\rm gap} \approx 2.5\times 10^3$ yr.
 Woeassuimoe that the infall rate decays exponentially over some timescale. rg.," We assume that the infall rate decays exponentially over some timescale, $\tau_{\rm off}$."
 This assumed. form for the iufall decay rate is also made by other authors in similar mvoestisations (?).., This assumed form for the infall decay rate is also made by other authors in similar investigations \citep{canup10}.
 We further asstune that the iufall decay time is of the same order as the gap opening timescale. Tog=Teap2.5ν10? Vr.," We further assume that the infall decay time is of the same order as the gap opening timescale, $\tau_{\rm off} = \tau_{\rm gap} = 2.5\times 10^3$ yr."
 Although it seems likely that eas counties to accrete through. the gap im a cimeumstellar disk opened by a erowing protoplauet. itf is uncertain if this eas is able to accrete either outo the protoplauet or even onto a cireuuplauetary disk.," Although it seems likely that gas continues to accrete through the gap in a circumstellar disk opened by a growing protoplanet, it is uncertain if this gas is able to accrete either onto the protoplanet or even onto a circumplanetary disk."
 We have chosen the exponcutial decay of πιαπιο material shown iu Equ., We have chosen the exponential decay of infalling material shown in Eqn.
 Ll because it is simple and easy to understaud., \ref{eqn:infall_decay} because it is simple and easy to understand.
 Also. it allows a direct. comparisonτσ with ?.. which use the same prescription for the expoucutial decay aud short πια. clecay nuescales.," Also, it allows a direct comparison with \citet{sasaki10}, which use the same prescription for the exponential decay and short infall decay timescales."
 Furthermore. whether it be via gap opening or the elobal depletion of the solar nebula. it is certain that the iufall from the solar nebula should halt at some point iu time.," Furthermore, whether it be via gap opening or the global depletion of the solar nebula, it is certain that the infall from the solar nebula should halt at some point in time."
 The only uncertainty is the timescale over which the infall wanes., The only uncertainty is the timescale over which the infall wanes.
 Receut observations of 56 weallined and classical T Tauri stars was used to determine the FUV cussion of vouus 110 Myr stars and whether the observed flux is cousistent with that required by models of photoevaporative dispersal of eirciuustellar disks (?).., Recent observations of $56$ weak-lined and classical T Tauri stars was used to determine the FUV emission of young $1-10$ Myr stars and whether the observed flux is consistent with that required by models of photoevaporative dispersal of circumstellar disks \citep{ingleby11}.
 These authors concluded that radiation fields sufiicicutly strong for the removal of gas are present during the disk dispersal phase., These authors concluded that radiation fields sufficiently strong for the removal of gas are present during the disk dispersal phase.
 Although not considered iu this study. x-ray flux las been shown to cuhanuce FUV-driven photoevaporation (?)..," Although not considered in this study, x-ray flux has been shown to enhance FUV-driven photoevaporation \citep{gorti09}."
 2 also investigated the x-ray cussion of l—10 Myr stars and found that the x-ray fux remains high. :xl constant. throughout he duration of the dispersal phase.," \citet{ingleby11} also investigated the x-ray emission of $1-10$ Myr stars and found that the x-ray flux remains high, and constant, throughout the duration of the dispersal phase."
 Clearly. 10 voung Sun would have exposed the ciretunplanetary nebulae of Jupiter aud Saturn to FUV radiation.," Clearly, the young Sun would have exposed the circumplanetary nebulae of Jupiter and Saturn to FUV radiation."
 Althoueh it is nuclear at this time row iuch FUV from the Sun i$ able to reach these circuniplanetary nebulae. one must remember tha he Sun most likely formed in a cluster of 10010 stars which could have contibuted to the FUV field in which these nebulae were embedded (27)..," Although it is unclear at this time how much FUV from the Sun is able to reach these circumplanetary nebulae, one must remember that the Sun most likely formed in a cluster of $10^3-10^4$ stars which could have contributed to the FUV field in which these nebulae were embedded \citep{fatuzzo08,adams10}."
 Whatever its source. the FUV radiation woul rave heated the periphery of the circumplanctary disks.," Whatever its source, the FUV radiation would have heated the periphery of the circumplanetary disks."
 Gas heated to sufficient temperatures wouk hen have become uubouud from these disks., Gas heated to sufficient temperatures would then have become unbound from these disks.
" The eravitational radius. ry. is defined as the radius at which the sound speed of the heated eas equals he escape speed from the system. where fis the Boltzmann coustaut aud T is the temperature of the super-heated αποριαο, or what we will refer to as the cuvelope temperature. Juss "," The gravitational radius, $r_{\rm g}$, is defined as the radius at which the sound speed of the heated gas equals the escape speed from the system, where $k$ is the Boltzmann constant and $T$ is the temperature of the super-heated atmosphere, or what we will refer to as the envelope temperature, $T_{\rm env}$."
Gas bevoud the eravitational radius will escape from the system., Gas beyond the gravitational radius will escape from the system.
" The gravitational radius is the canonical radius bevoud which eas heated to a temperature Ti, will escape from the disk.", The gravitational radius is the canonical radius beyond which gas heated to a temperature $T_{\rm env}$ will escape from the disk.
 Tn actuality. eas can escape from the disk at radii substantially smaller than ry.," In actuality, gas can escape from the disk at radii substantially smaller than $r_{\rm g}$."
" The heating of the disks surface by FUV radiation aud the resultant outflow are complicated processes, but if is useful to emiplov a simplified model with an isothermal. heated atmosphere."," The heating of the disk's surface by FUV radiation and the resultant outflow are complicated processes, but it is useful to employ a simplified model with an isothermal, heated atmosphere."
 Consider a disk irradiated and heated by external FUV radiation from either the carly Sun or the fellow members of its birth cluster., Consider a disk irradiated and heated by external FUV radiation from either the early Sun or the fellow members of its birth cluster.
 Depeucding ou the streneth of the PFUV flux. the heated eas will reach temperatures in the range 100EK«T«3000Iv (?).," Depending on the strength of the FUV flux, the heated gas will reach temperatures in the range $100\ {\rm K} < T< 3000\ {\rm K}$ \citep{adams04}."
 As the eas heats. it expands eeneratius a neutral outflow.," As the gas heats, it expands generating a neutral outflow."
 The cexpaucing outflow beeius subsonically but becomes supersonic by the time it reaches the gravitational radius., The expanding outflow begins subsonically but becomes supersonic by the time it reaches the gravitational radius.
 This outflow is eonerally isotropic. but the majority of nass loss is doiuinated by nass loss from the," This outflow is generally isotropic, but the majority of mass loss is dominated by mass loss from the"
"1,---,M, where M is the number of observed multiplets).",", where $M$ is the number of observed multiplets)."
" Here, 6v© corresponds to the average separation within a given multiplet."," Here, $\delta \nu_{i}^{\rm O}$ corresponds to the average separation within a given multiplet."
" Because the problem is intrinsically ill-posed, the inversion of the above equation necessarily must be regularized (Jeffrey 1988, Kawaler et al."," Because the problem is intrinsically ill-posed, the inversion of the above equation necessarily must be regularized (Jeffrey 1988, Kawaler et al."
 1999)., 1999).
" To do so we minimize: where the second term of the right hand side is, precisely, the regularization term."," To do so we minimize: where the second term of the right hand side is, precisely, the regularization term."
 The form of the regularization term determines some additional constraints to the solution., The form of the regularization term determines some additional constraints to the solution.
" For instance, if Q(r) cannot have a steep spatial gradient, then £= d/dr."," For instance, if $\Omega(r)$ cannot have a steep spatial gradient, then ${\cal
L}\equiv d/dr$ ."
" On the other hand, if O(r) must be smooth, then L= d?/dr?.The theoretical splittings óv7 are computed numerically: where Ω/=Q(rj) and K;;=Ki(r;), being the index j (j= 1,---,N) associated to the radial mesh point in the stellar model on which the kernel K;(r) and the rotation rate Q(r) are evaluated, and wj=rj41—rj."," On the other hand, if $\Omega(r)$ must be smooth, then ${\cal L}\equiv d^2/dr^2$ .The theoretical splittings $\delta
\nu_{i}^{\rm T}$ are computed numerically: where $\Omega_j= \Omega(r_j)$ and $K_{ij}= K_i(r_j)$, being the index $j$ $j=1, \cdots, N$ ) associated to the radial mesh point in the stellar model on which the kernel $K_i(r)$ and the rotation rate $\Omega(r)$ are evaluated, and $w_j= r_{j+1}-r_j$."
" In the least squares method, the derivative of the sum of the squared residuals (the function $) with respect to 0; is taken and equated to zero."," In the least squares method, the derivative of the sum of the squared residuals (the function $S$ ) with respect to $\Omega_j$ is taken and equated to zero."
" After some algebra, the following matrix equation is derived: where K is a (NxM) matrix with elements (K);;=wy;Kij /o1, Dis a (Nx1) vector with components (£2);=€, and Y is a (Mx1) vector with components (Y);=ὃν,/σι."," After some algebra, the following matrix equation is derived: where $\mathbf{K}$ is a $(N \times M)$ matrix with elements $(\mathbf{K})_{ij}= w_j K_{ij}/\sigma_i$ , $\mathbf{\Omega}$ is a $(N
\times 1)$ vector with components $(\mathbf{\Omega})_{j}= \Omega_j$, and $\mathbf{\Upsilon}$ is a $(M \times 1)$ vector with components $(\mathbf{\Upsilon})_{i}= \delta \nu_i^{\rm O}/\sigma_i$."
" Finally H is the (NxN) regularization matrix, which adopts a tridiagonal form (that is, (H);;=0 for |i—j|> 1) or a pentadiagonal structure (that is, (H);;=0 for |i—j|> 2) depending on whether £=d/dr or £&d?/dr?."," Finally, $\mathbf{H}$ is the $(N \times N)$ regularization matrix, which adopts a tridiagonal form (that is, $(\mathbf{H})_{ij}= 0$ for $|i-j|>1$ ) or a pentadiagonal structure (that is, $(\mathbf{H})_{ij}= 0$ for $|i-j|>2$ ) depending on whether ${\cal L}\equiv d/dr$ or ${\cal L}
\equiv d^2/dr^2$."
 Eq. (10)), Eq. \ref{matricial}) )
 constitutes a (ΝxN) system of linear equations that must be solved for the unknown rotation velocities €; that minimize S., constitutes a $(N \times N)$ system of linear equations that must be solved for the unknown rotation velocities $\Omega_j$ that minimize $S$ .
" We have tested the reliability of our RLS scheme by employing this technique on synthetic"" (free of uncertainties) frequency splittings generated with the asteroseismological model of tthrough the forward approach.", We have tested the reliability of our RLS scheme by employing this technique on “synthetic” (free of uncertainties) frequency splittings generated with the asteroseismological model of through the forward approach.
" Specifically, we considered rotational splittings corresponding to consecutive /(=1 g-modes with k=1,---,40."," Specifically, we considered rotational splittings corresponding to consecutive $\ell= 1$ $g$ -modes with $k= 1,\cdots, 40$."
 We employed the LU decomposition and also the Gauss-Jordan methods (Press et al., We employed the LU decomposition and also the Gauss-Jordan methods (Press et al.
 1992) for solve the system., 1992) for solve the system.
 Both methods give almost identical results., Both methods give almost identical results.
" In all of the cases we have examined, the inversions are able to recover the input rotation profile that we used to compute the synthetic splittings, provided that an adequate range of values of the parameter A is adopted."," In all of the cases we have examined, the inversions are able to recover the input rotation profile that we used to compute the synthetic splittings, provided that an adequate range of values of the parameter $\lambda$ is adopted."
 We have applied the RLS method to infer the internal rotation profile of01224-200., We have applied the RLS method to infer the internal rotation profile of.
. We have employed the M-Τ averaged £=1 splittings., We have employed the $M= 7$ averaged $\ell= 1$ splittings.
 The regularization matrix corresponds to the smoothing of the second derivative of Q(r)., The regularization matrix corresponds to the smoothing of the second derivative of $\Omega(r)$ .
 In Fig., In Fig.
 6 we show the inverted rotation profiles for ffor several values of A., \ref{inv-pg0122} we show the inverted rotation profiles for for several values of $\lambda$.
" For very small values of A, the inverted profiles exhibit strong variations that lack physical meaning."," For very small values of $\lambda$, the inverted profiles exhibit strong variations that lack physical meaning."
" However, as the value of A isincreased, the inverted solution gradually stabilizes."," However, as the value of $\lambda$ isincreased, the inverted solution gradually stabilizes."
" The resulting rotation profile (corresponding to A>10?) consists of an almost linearly decreasing rotation rate with €),~10.75Hz and Ως~4.58 uHz, in excellent agreement with the results of the forward approach."," The resulting rotation profile (corresponding to $\lambda \gtrsim
10^{-2}$ ) consists of an almost linearly decreasing rotation rate with $\Omega_{\rm c} \sim 10.75\, \mu$ Hz and $\Omega_{\rm s} \sim 4.58\,
\mu$ Hz, in excellent agreement with the results of the forward approach."
" The monotonic linear functional form characterizing the inverted rotation profiles should not be surprising, since we are forcing Q(r) to have a small value of its second derivative at the outset."," The monotonic linear functional form characterizing the inverted rotation profiles should not be surprising, since we are forcing $\Omega(r)$ to have a small value of its second derivative at the outset."
" An analysis of the uncertainties similar to that performed for the forward approach leads to the conclusion that even with the inclusion of uncertainties in the observed splittings, the rotation of iis faster at the central regions than at the surface."," An analysis of the uncertainties similar to that performed for the forward approach leads to the conclusion that even with the inclusion of uncertainties in the observed splittings, the rotation of is faster at the central regions than at the surface."
" Specifically, we found Q.=10.75+ 2.4uHz and Ως=4.58+1.7 wHz."," Specifically, we found $\Omega_{\rm c}= 10.75\pm 2.4\, \mu$ Hz and $\Omega_{\rm s}= 4.58\pm 1.7\, \mu$ Hz."
" We have also made rotational inversions onto a fixed functional basis — a method called “function fitting”, see Kawaleret al. ("," We have also made rotational inversions onto a fixed functional basis — a method called “function fitting”, see Kawaleret al. ("
1999) for details.,1999) for details.
" In this inversion technique, an explicit assumption about the functional form of Q(r) is made."," In this inversion technique, an explicit assumption about the functional form of $\Omega(r)$ is made."
" For instance, if Q(r) is assumed to be a polynomial in r of degree K—1, then Q(r)=σιαλα."," For instance, if $\Omega(r)$ is assumed to be a polynomial in $r$ of degree $K-1$, then $\Omega(r)= \sum_{k=1}^{K} a_k r^{k-1}$."
" Following the least squares method, in which we minimize the sum of the squared residuals, we obtain the matrix equation: where A is a (MxK) matrix with elements ais a (dKx1) vector with components (a),=ax, and Y isa (Mx1) vectorwith components (Y);=óv?/o;. Eq. (11))"," Following the least squares method, in which we minimize the sum of the squared residuals, we obtain the matrix equation: where $\mathbf{A}$ is a $(M \times K)$ matrix with elements $\mathbf{a}$ is a $(K \times 1)$ vector with components $(\mathbf{a})_{k}= a_k$, and $\mathbf{\Upsilon}$ is a $(M \times 1)$ vectorwith components $(\mathbf{\Upsilon})_{i}= \delta \nu_i^{\rm
 O}/\sigma_i$ Eq. \ref{matricial2}) )"
 is a (KxK) system of linear equations that must be solved for the unknown set of values of αι. that minimize S., is a $(K \times K)$ system of linear equations that must be solved for the unknown set of values of $a_k$ that minimize $S$ .
" Specifically,we have explored linear (K— 2) functional formsfor Q(r), defined by the two parameters a1=X and a»=(Qs— Ὡς)."," Specifically,we have explored linear $K= 2$ ) functional formsfor $\Omega(r)$ , defined by the two parameters $a_1= \Omega_{\rm c}$ and $a_2= (\Omega_{\rm s}- \Omega_{\rm c})$ ."
" The optimal values we found for these parameters are Q,=10.74+2.9 wHz and Ως=4.57+1.8 wHz, in excellent agreement with the RLS fits and also withthe forward approach."," The optimal values we found for these parameters are $\Omega_{\rm c}= 10.74\pm 2.9\, \mu$ Hz and $\Omega_{\rm s}= 4.57\pm 1.8\, \mu$ Hz, in excellent agreement with the RLS fits and also withthe forward approach."
Now that the offset aid rotation for each chip/filter combination had been determined to place each of the chips iuto a «‘OUMNOL nmnaste ‘frame. the overall scale was tle last of tle linear parameters o be solved for.,"Now that the offset and rotation for each chip/filter combination had been determined to place each of the chips into a common master frame, the overall scale was the last of the linear parameters to be solved for."
 To do this. we coustructed a master frame based on ouly the FOOGW exposures ine oulv trausformatiois allowed. for offset. and. rotation. oit no scale changes.," To do this, we constructed a master frame based on only the F606W exposures using only transformations allowed for offset and rotation, but no scale changes."
 This way. the rane would represent {he average scale ol the F606W expostLes.," This way, the frame would represent the average scale of the F606W exposures."
 We then transformec the positions «X tlie stars tn each of the exposures for each filter (203 in otal. including the PID-!L209f cat aj}to tlis reference frame and took note of the scale factor of he linear translormatiou.," We then transformed the positions of the stars in each of the exposures for each filter (203 in total, including the PID-12094 data) into this reference frame and took note of the scale factor of the linear transformation."
 We plot his scale [actor for each exposure on the top pauels of Figure 9.., We plot this scale factor for each exposure on the top panels of Figure \ref{fig:scale_by_filt}.
 The images are ordere by filter aud tιο filers are separated by a vertical dotted line., The images are ordered by filter and the filters are separated by a vertical dotted line.
 It is clear hat there is some (rei with filter. but there is cousiderable ittra-filler scatter as well.," It is clear that there is some trend with filter, but there is considerable intra-filter scatter as well."
 This scatter could be due o velocity aberratiol Du realing., This scatter could be due to velocity aberration or breathing.
 We divided each o “the above scale measurements by the keyword. taken [roi the image header.," We divided each of the above scale measurements by the keyword, taken from the image header."
 It. reports the expecte special-relativistic variation in the plate scale. which is related to the dot-prodwt between the <werage telescope velocity vector during the exposure and the clirecticmu to the target (see Cox Ciillilaxl 2002).," It reports the expected special-relativistic variation in the plate scale, which is related to the dot-product between the average telescope velocity vector during the exposure and the direction to the target (see Cox Gilliland 2002)."
 After making this deterministic correction. we found ttat the global exposures for each Lilter were now cousistent to witun 0.01 pixel (as shown ou the bottom pauels of Fig. 9)).," After making this deterministic correction, we found that the global exposures for each filter were now consistent to within 0.01 pixel (as shown on the bottom panels of Fig. \ref{fig:scale_by_filt}) )."
 We solved or tlie average scaling for each filter aud included this correctjou in the meta distortion solution., We solved for the average scaling for each filter and included this correction in the meta distortion solution.
 The aveage scaling we found or each filter is given iu the eighth column of Table )3.., The average scaling we found for each filter is given in the eighth column of Table \ref{tab3}.
 They are d1 good qualitative agreement with the findings in Figure 2 0- WNoghurina-Platais et ((2010)., They are in good qualitative agreement with the findings in Figure 2 of Kozhurina-Platais et (2010).
 We note hat the linear terius o. UVIS appear to be considerabM7 nore stable than those of ACS/WFC., We note that the linear terms of UVIS appear to be considerably more stable than those of ACS/WFC.
 Fi[n]eure 6 of Anderson (20t06) shows that the ACS FOV clanges in radius by £0.03 pixel due to breatune. while the residuas here are about 0.01 pixel.," Figure 6 of Anderson (2006) shows that the ACS FOV changes in radius by $\pm$ 0.03 pixel due to breathing, while the residuals here are about 0.01 pixel."
 We want enijohasize that breathing can introduce errors larger tiu 0.008 pixel into the, We want emphasize that breathing can introduce errors larger than 0.008 pixel into the
((52-100 mpe).,(52-100 mpc).
 A similar structure is also seen in the optical image. but it is much fainter and its shape is not identical to the ring in the 3.29 image.," A similar structure is also seen in the optical image, but it is much fainter and its shape is not identical to the ring in the 3.29 image."
 The diffuse nebulosity at 3.29 iis located primarily between the filaments and the central star and is only faintly detected farther from the star than the filaments., The diffuse nebulosity at 3.29 is located primarily between the filaments and the central star and is only faintly detected farther from the star than the filaments.
 Our efforts to correct the 3.29 image for ~3.3 ccontinuum emission. by obtaining an L’ image and combining it with à 2 ccontinuum image to interpolate a ~3.3 ccontinuum image. were unsuccessful due to the high sky background at L’.," Our efforts to correct the 3.29 image for $\sim$ 3.3 continuum emission, by obtaining an $L'$ image and combining it with a 2 continuum image to interpolate a $\sim$ 3.3 continuum image, were unsuccessful due to the high sky background at $L'$."
 Previous studies. however. that have quantified the relative strengths of the 3.29 IIEF emission to the underlying continuum emission. establish that the ~3.3 ccontinuum contributes less than to narrowband 3.29 pphotometry at all observed positions in citepSWD83.SWA96..," Previous studies, however, that have quantified the relative strengths of the 3.29 IEF emission to the underlying continuum emission, establish that the $\sim$ 3.3 continuum contributes less than to narrowband 3.29 photometry at all observed positions in \\citep{SWD83,SWA96}."
 We therefore believe that our 3.29 image is dominated by 3.29 HEF emission., We therefore believe that our 3.29 image is dominated by 3.29 IEF emission.
 Figure 4+ illustrates positions. overlaid on our 3.294 image. where previous photometry and spectrophotometry have quantified the contribution of the underlying ~3.3 ccontinuum surface brightness to the total observed surface brightness at 3.29 ((hereafter £).," Figure \ref{f4} illustrates positions, overlaid on our 3.29 image, where previous photometry and spectrophotometry have quantified the contribution of the underlying $\sim$ 3.3 continuum surface brightness to the total observed surface brightness at 3.29 (hereafter $f$ )."
 At position C in Figure 4.. the 3.0-3.7 sspectrophotometry of Sellgrenetal.(1983) implies f =14%.," At position C in Figure \ref{f4}, the 3.0-3.7 spectrophotometry of \citet{SWD83} implies $f$ =."
. For all other positions plotted on Figure 4.. we derive f from surface photometry at K (broadband 2.2 jm). narrowband 3.29 µπι.. and { of Sellgrenetal.(1996).," For all other positions plotted on Figure \ref{f4}, , we derive $f$ from surface photometry at $K$ (broadband 2.2 ), narrowband 3.29 , and $L'$ of \citet{SWA96}."
. We interpolate the continuum at 3.3 aat each position from K and L’ data (Sellgrenetal.1996).. assuming that the continuum emission between 2.2 aand 3.5 ccan be approximated by a power-law (Sellgrenetal.1983.1985.1996).," We interpolate the continuum at 3.3 at each position from $K$ and $L'$ data \citep{SWA96}, assuming that the continuum emission between 2.2 and 3.8 can be approximated by a power-law \citep{SWD83,SAB85,SWA96}."
. We find that f =18%..13%..12%.. 10%... and 17%.. at positions A. B. I. 2. and 3. respectively.," We find that $f$ =, and , at positions A, B, 1, 2, and 3, respectively."
 Note that most positions illustrated in Figure 4+ fall on diffuse 3.29 HEF emission. with f = 12-18%.. except for Position 2. which falls on Filament [ and has f =10%.," Note that most positions illustrated in Figure \ref{f4} fall on diffuse 3.29 IEF emission, with $f$ = , except for Position 2, which falls on Filament I and has $f$ =."
. These published observations demonstrate that the 3.294 HEF emission accounts for >82% of the emission in our 3.29 image.," These published observations demonstrate that the 3.29 IEF emission accounts for $\ga
82\%$ of the emission in our 3.29 image."
 Narrow-band images at 2.09 aand 2.14 ((shown as a combination in Fig. 3)).," Narrow-band images at 2.09 and 2.14 (shown as a combination in Fig. \ref{f3}) ),"
" and at 2.18 alsoshowninFig.3]]LFG96 are at wavelengths free of eemission lines (Martinietal.1999)., and thus represent the 2 ccontinuum emission."," and at 2.18 \\citep[ also shown
in Fig.~\ref{f3}] are at wavelengths free of emission lines \citep{MSD99}, and thus represent the 2 continuum emission."
 Apart from differences in seeing and a few isolated instrumental artifacts. the two images of the 2 ccontinuum emission shown in Figure 3. display a very similar morphology.," Apart from differences in seeing and a few isolated instrumental artifacts, the two images of the 2 continuum emission shown in Figure \ref{f3} display a very similar morphology."
 While neither 2 ccontinuum image in Figure 3 has been corrected for instrumental scattered light. they were obtained on different telescopes. with different cameras and different orientations of the diffraction spikes of oon the image.," While neither 2 continuum image in Figure \ref{f3} has been corrected for instrumental scattered light, they were obtained on different telescopes, with different cameras and different orientations of the diffraction spikes of on the image."
 Hence. their similarity implies that the main morphological structures ineach image are intrinsic to aandare not artifacts of the telescope. the infrared camera. orthe instrumental scattered light.," Hence, their similarity implies that the main morphological structures ineach image are intrinsic to andare not artifacts of the telescope, the infrared camera, orthe instrumental scattered light."
 The 2, The 2
 Lin the 0.5-2.0 keV band. or a bolometric luminosity of —-3x 10H ergs +.,"$^{-1}$ in the 0.5-2.0 keV band, or a bolometric luminosity of $\sim$ $\times$ $^{41}$ ergs $^{-1}$."
" The luminosity- relation for radio-quiet groups from Croston,Hardeastle,Birkinshaw(2005) predicts a luminosity of ~4x 10! eres ! for groups with a temperature of ~1.8 keV, so that the 333 group appears to be roughly two orders of magnitude less luminous than would be expected for its temperature, unless both and are failing to detect a very large, smooth extended environment."," The luminosity-temperature relation for radio-quiet groups from \citet{cro05a} predicts a luminosity of $\sim$ $\times$ $^{43}$ ergs $^{-1}$ for groups with a temperature of $\sim$ 1.8 keV, so that the 33 group appears to be roughly two orders of magnitude less luminous than would be expected for its temperature, unless both and are failing to detect a very large, smooth extended environment."
 A temperature of —0.5 keV is predicted for our measured luminosity., A temperature of $\sim$ 0.5 keV is predicted for our measured luminosity.
" This suggests either that 333 (or previous generations of outbursts from its AGN) has had a dramatic effect on the surrounding group, or that powertul radio galaxies can form in extremely poor environments."," This suggests either that 33 (or previous generations of outbursts from its AGN) has had a dramatic effect on the surrounding group, or that powerful radio galaxies can form in extremely poor environments."
" We have fitted three models to the radio through X-ray flux density distributions of the seven regions shown in Figures 7. through 10:: a broken power-law synchrotron model with a high energy cut-off, a synchrotron self-Compton model (SSC) in which the X-ray emission is the result of inverse-Compton scattering of the synchrotron radio photons off the relativistic electrons, and a model of inverse Compton scattering of CMB photons (IC/CMB) off the radio synchrotron emitting relativistic electrons."," We have fitted three models to the radio through X-ray flux density distributions of the seven regions shown in Figures \ref{s12sed} through \ref{n3sed}: a broken power-law synchrotron model with a high energy cut-off, a synchrotron self-Compton model (SSC) in which the X-ray emission is the result of inverse-Compton scattering of the synchrotron radio photons off the relativistic electrons, and a model of inverse Compton scattering of CMB photons (IC/CMB) off the radio synchrotron emitting relativistic electrons."
" Other sources of seed photons in the inverse-Compton scenario including optical photons from the stellar population of the host galaxy or beamed optical/IR photons from a hidden nuclear jet are implausible because of the distance of the hot spots from the nucleus (see Hardcastle,Birkinshaw,&Worrall(1998) for details).", Other sources of seed photons in the inverse-Compton scenario including optical photons from the stellar population of the host galaxy or beamed optical/IR photons from a hidden nuclear jet are implausible because of the distance of the hot spots from the nucleus (see \citet{mjh98} for details).
" For example, the energy density of IR photons emitted from the nucleus (based on the nuclear flux measured in theSpitzer images) is roughly two orders of magnitude below the energy density of the CMB at the distance of the hot spots."," For example, the energy density of IR photons emitted from the nucleus (based on the nuclear flux measured in the images) is roughly two orders of magnitude below the energy density of the CMB at the distance of the hot spots."
" The relevant model parameters for the synchrotron, IC/CMB, and SSC curves shown in Figures 7 through 10 are summarized in Table 6.."," The relevant model parameters for the synchrotron, IC/CMB, and SSC curves shown in Figures \ref{s12sed} through \ref{n3sed} are summarized in Table \ref{fitptab}."
" We assume a power-law electron energy distribution for the synchrotron fits with +,,,;,=100 and 5,,,.210*.", We assume a power-law electron energy distribution for the synchrotron fits with $\gamma_{min}$ =100 and $\gamma_{max}$ $^{8}$.
 We also allow for a break in the electron energy distribution where the index of the power law steepens., We also allow for a break in the electron energy distribution where the index of the power law steepens.
 We choose to plot synchrotron models for S3 and S4 with relatively large breaks for consistency with SI and S2., We choose to plot synchrotron models for S3 and S4 with relatively large breaks for consistency with S1 and S2.
" For the latter two regions, the large change in power-law index is required by the optical and IR data."," For the latter two regions, the large change in power-law index is required by the optical and IR data."
" The flux density distributions of S3 and S4 could be modeled with a smaller change in index which would be roughly consistent with the optical and IR limits, although the lack of detections in these bands prevent a definitive conclusion."," The flux density distributions of S3 and S4 could be modeled with a smaller change in index which would be roughly consistent with the optical and IR limits, although the lack of detections in these bands prevent a definitive conclusion."
 We emphasize that in regions S3 and S4. we must be averaging over considerable substructure in the synchrotron scenario because," We emphasize that in regions S3 and S4, we must be averaging over considerable substructure in the synchrotron scenario because"
metallicity depends on the adopted reddeuiug.,metallicity depends on the adopted reddening.
 In Table 5 we present the results for our calibration (column L} in the range of published reddeniugs (colum 3) iu comparison to otherwise determined iron abuudauces (column 5. SCC ILhuxis 1996.. Fricl 1995.. aud the Authouy-Twaroe ct al.," In Table 5 we present the results for our calibration (column 4) in the range of published reddenings (column 3) in comparison to otherwise determined iron abundances (column 5, see Harris \cite{harr96a}, Friel \cite{frie}, and the Anthony-Twarog et al."
 papers for references)., papers for references).
 Coluun 2 gives the umber of stars duvolved in the metallicity determination., Column 2 gives the number of stars involved in the metallicity determination.
 Our results agree well with the published irom abuudances., Our results agree well with the published iron abundances.
 Also the isoanetallicitv lines of the xed ejauts in the (bgu) diaerain agree very well with the slopes of the jw. calibration for all clusters.," Also the iso-metallicity lines of the red giants in the $(b-y),m_1$ diagram agree very well with the slopes of the new calibration for all clusters."
 NGC 2395 has a high intrinsic dispersion simular to M22 and w Cen., NGC 2395 has a high intrinsic dispersion similar to M22 and $\omega$ Cen.
 The high dispersion for NGC 6397 is due to the fact that all giauts have (bdg) colors bluer than 0.55 mae. a very metal insensitive regime for Strónuueren metallicities.," The high dispersion for NGC 6397 is due to the fact that all giants have $(b-y)$ colors bluer than 0.55 mag, a very metal insensitive regime for Strömmgren metallicities."
 Red eiauts in the elobular clusters w Centauri. M55. aud AI22 together with field eiauts from Authouv-Twarog Twarog (1998)) have been used to revise the Carebel Richtler (19923) metallicity calibration of the Stromuueren (bg)any diaerain.," Red giants in the globular clusters $\omega$ Centauri, M55, and M22 together with field giants from Anthony-Twarog Twarog \cite{anth98}) ) have been used to revise the Grebel Richtler \cite{greb92}) ) metallicity calibration of the Strömmgren $(b-y),m_1$ diagram."
 For all eiauts iun w Cen and M22. accurate and homogeneous iron abundauces from high resolution spectroscopy are availalde inthe literature.," For all giants in $\omega$ Cen and M22, accurate and homogeneous iron abundances from high resolution spectroscopy are available in the literature."
 M55 jas a well determined average mon abundance value., M55 has a well determined average iron abundance value.
 In otal. 62 CN-weak giants have been used.," In total, 62 CN-weak giants have been used."
 CN-vich stars rave been excluded. since their ns value mnüdndes a too Hel iron abundance in the (5gU).mn diagrai.," CN-rich stars have been excluded, since their $m_1$ value mimics a too high iron abundance in the $(b-y),m_1$ diagram."
 Iu order ο cover a wide metallicity range. 2.0< /TI]«0.0 dex. our new calibration is connected to a previous calibration w Crebel Richtler (1992)) around solar metallicitics.," In order to cover a wide metallicity range, $-2.0 <$ $< 0.0$ dex, our new calibration is connected to a previous calibration by Grebel Richtler \cite{greb92}) ) around solar metallicities."
 Iu the color range 0.5<(yy)«1.l mag. for which our calibration is valid. the loci of equal iron abundances lic on straight lines.," In the color range $0.5 < (b-y)
< 1.1$ mag, for which our calibration is valid, the loci of equal iron abundances lie on straight lines."
 We emphasize that it was possible to find a nifonu metallicity calibration in the indicated parameter range that secs to have no obvious depeudeucies ou Inuuinosity witlin the errors., We emphasize that it was possible to find a uniform metallicity calibration in the indicated parameter range that seems to have no obvious dependencies on luminosity within the errors.
 In particular. the new calibration seenis to be valid for globular cluster as well as for field giants.," In particular, the new calibration seems to be valid for globular cluster as well as for field giants."
 No variation of metallicity seusitivitv with metallicity iu the (b—g).my diagram has been found.," No variation of metallicity sensitivity with metallicity in the $(b-y),m_1$ diagram has been found."
 Ou average. the precision of a ποσατν deteziuunation with our new calibration for a single ejut is in the order of 0.11 dex.," On average, the precision of a metallicity determination with our new calibration for a single giant is in the order of 0.11 dex."
 Average ατασος of giants within a cluster andl relative abuudauces between clusters can be deteriuued with much higher recision. depeudiugc» on the nmuuber of red [m]eiauts per cluster sample.," Average abundances of giants within a cluster and relative abundances between clusters can be determined with much higher precision, depending on the number of red giants per cluster sample."
 The application of the new calibration to indepeudent samples of red giants with published Stronunerch photometry agrees very well with otherwise determined abundances., The application of the new calibration to independent samples of red giants with published Strömmgren photometry agrees very well with otherwise determined abundances.
 For the red giauts iu ιο Cen and M22. the influence of CN-strong stars on the metallicity calibration has been studied.," For the red giants in $\omega$ Cen and M22, the influence of CN-strong stars on the metallicity calibration has been studied."
 For Stromunercu metallicities higher thau 1.0 dex. CNaveak stars cannot be distinguished iu the (b.g).ing diagiun froin stars with lower iron abundances but higher CN 1und strengths.," For Strömmgren metallicities higher than $-$ 1.0 dex, CN-weak stars cannot be distinguished in the $(b-y),m_1$ diagram from stars with lower iron abundances but higher CN band strengths."
 Iowever. the difference between the Stromunercu inetallicitv of CN-rich stars and their spectroscopically deteriunued mon abundance is correlated to the CN baud streneths.," However, the difference between the Strömmgren metallicity of CN-rich stars and their spectroscopically determined iron abundance is correlated to the CN band strengths."
 This night be used to determine the iron abundance of a red eint if its Strónumueren colors aud one of the CN iudices are kuown. or alternatively to detect CN-vich stars when their iron abundances are known.," This might be used to determine the iron abundance of a red giant if its Strömmgren colors and one of the CN indices are known, or alternatively to detect CN-rich stars when their iron abundances are known."
 This. research was supported through νο yecto FONDECYT 3980032, This research was supported through `Pro- yecto FONDECYT 3980032'.
"""I thank Tom Richtler and Boris Dirsch for interesting discussions. and the referee for very helpful comments that improved the paper."," I thank Tom Richtler and Boris Dirsch for interesting discussions, and the referee for very helpful comments that improved the paper."
spherical coordinates (16) assuming axisymmetry around the rotation axis of the star-disk system.,"spherical coordinates $(R,\theta)$ assuming axisymmetry around the rotation axis of the star-disk system."
" Note that, in this configuration, the heat pulse is not localized in a relatively small portion of the disk (as in our 3D simulation), but is distributed in a ring."," Note that, in this configuration, the heat pulse is not localized in a relatively small portion of the disk (as in our 3D simulation), but is distributed in a ring."
" Nevertheless, we found that the evolution of the flare and the stream described by the 2.5D simulation is quite similar to that of our 3D simulation."," Nevertheless, we found that the evolution of the flare and the stream described by the 2.5D simulation is quite similar to that of our 3D simulation."
" The 2.5D simulation allowed us to extend our analysis of the stream dynamics, following its evolution until t=100 hours (i.e. approximately four days)."," The 2.5D simulation allowed us to extend our analysis of the stream dynamics, following its evolution until $t= 100$ hours (i.e. approximately four days)."
 Figure 8 shows the profiles of the relevant velocities derived from the 2.5D simulation at t=36 hours (upper panel) and t=68 hours (lower panel)., Figure \ref{stream_prof_2D} shows the profiles of the relevant velocities derived from the 2.5D simulation at $t=36$ hours (upper panel) and $t=68$ hours (lower panel).
 The velocity profiles at t=36 hours resemble those derived from the 3D simulation (lower panel in Fig. 7))., The velocity profiles at $t=36$ hours resemble those derived from the 3D simulation (lower panel in Fig. \ref{fig4_s}) ).
" In particular, the flow slightly brakes approaching the stellar surface due to a pressure gradient force slowing down the free-fall of matter, as found in the 3D simulation."," In particular, the flow slightly brakes approaching the stellar surface due to a pressure gradient force slowing down the free-fall of matter, as found in the 3D simulation."
" In addition, the 2.5D simulation shows that, at this stage, the stream has not reached yet a quasi-stationary condition."," In addition, the 2.5D simulation shows that, at this stage, the stream has not reached yet a quasi-stationary condition."
" Later, the stream stabilizes (after t~60 hours); the gravitational force becomes dominant along the stream and the matter is accelerated until it impacts the stellar surface, reaching there a maximum velocity of us&300 km s! corresponding to z0.9ug."," Later, the stream stabilizes (after $t\approx 60$ hours); the gravitational force becomes dominant along the stream and the matter is accelerated until it impacts the stellar surface, reaching there a maximum velocity of $u\rs{str}\approx
300$ km $^{-1}$ corresponding to $\approx 0.9\, u\rs{ff}$."
" The above results, therefore, show that the physical characteristics of the accretion stream triggered by the flare closely recall those, largely discussed in the literature, of streams driven by the accumulation of mass at the disk truncation radius under the effect of the viscosity and pushed out of the equatorial plane because of the growing pressure gradient there (????))."," The above results, therefore, show that the physical characteristics of the accretion stream triggered by the flare closely recall those, largely discussed in the literature, of streams driven by the accumulation of mass at the disk truncation radius under the effect of the viscosity and pushed out of the equatorial plane because of the growing pressure gradient there \citealt{2002ApJ...578..420R, Romanova2003ApJ, 2008A&A...478..155B,
2009A&A...508.1117Z}) )."
 Our simulation shows that the stream is relatively cold (its temperature remains below 1 MK)., Our simulation shows that the stream is relatively cold (its temperature remains below 1 MK).
" After the disk material enters the stream, its density slightly decreases and, close to the stellar surface, increases again as a result of the gas compression by the dipolar magnetic field."," After the disk material enters the stream, its density slightly decreases and, close to the stellar surface, increases again as a result of the gas compression by the dipolar magnetic field."
" The stream velocity ustrgradually increases toward the star: it becomes supersonic already at a distance of zz2x10"" cm from the disk, and us approaches the free- speed ug close to the stellar surface."," The stream velocity $u\rs{str}$gradually increases toward the star: it becomes supersonic already at a distance of $\approx 2\times
10^{11}$ cm from the disk, and $u\rs{str}$ approaches the free-fall speed $u\rs{ff}$ close to the stellar surface."
 Figure 9 shows the density distribution on the stellar surface at t=36 hours; the dense spot on the surface is the region of impact of the stream., Figure \ref{acc_spot} shows the density distribution on the stellar surface at $t=36$ hours; the dense spot on the surface is the region of impact of the stream.
" The spot covers a small percentage of the stellar surface and the stream is inhomogeneous with its mass density varying across the stream and being the largest in the inner region, according to ?.."," The spot covers a small percentage of the stellar surface and the stream is inhomogeneous with its mass density varying across the stream and being the largest in the inner region, according to \cite{Romanova2004ApJ}. ."
" Finally, we derived the mass accretion rate M due to the"," Finally, we derived the mass accretion rate $\dot{M}$ due to the"
allows a nearly uniform. coverage of the epoch difference of 92 years.,allows a nearly uniform coverage of the epoch difference of 92 years.
 The limiting magnitude of the plates is of the order of V—15.5 mag. while the corresponding one of the CCD frames ranges from V—16 to 19 mag.," The limiting magnitude of the plates is of the order of $V = 15.5\;$ mag, while the corresponding one of the CCD frames ranges from $V = 16$ to 19 mag."
 The refractor plates of Bonn were mainly scanned at the PDS 2020GM of University of Münnster., The refractor plates of Bonn were mainly scanned at the PDS 2020GM of University of Münnster.
 The plates R1874 and R1875 with lower limiting magnitude have been measured at the ASCORECORD of Hoher List observatory., The plates R1874 and R1875 with lower limiting magnitude have been measured at the ASCORECORD of Hoher List observatory.
 On these plates only 80 stars for each plate were measured., On these plates only 80 stars for each plate were measured.
 Stars were extracted and rectangular coordinates x and y were determined from the PDS measurements using standard procedures (e.g. Tucholke 1994)., Stars were extracted and rectangular coordinates $x$ and $y$ were determined from the PDS measurements using standard procedures (e.g. Tucholke 1994).
 The first epoch plates of the Bonn refractor contained scratches and reseau lines. which led to problems for a significant number of the stars.," The first epoch plates of the Bonn refractor contained scratches and reseau lines, which led to problems for a significant number of the stars."
 Therefore for some of these stars no rectangular coordinates could be obtained., Therefore for some of these stars no rectangular coordinates could be obtained.
 The plates from Shanghai were scanned at the PDS 1010MS of the Purple Mountain Observatory. Chinese Academy of sciences (see also Wang et 1999).," The plates from Shanghai were scanned at the PDS 1010MS of the Purple Mountain Observatory, Chinese Academy of sciences (see also Wang et 1999)."
 The determination of the rectangular coordinates x. v of the stars on the CCD frames was performed for the observations from Hoher List by standard CCD reduction techniques (DAOPHOT. ΗΑΕ) routines.," The determination of the rectangular coordinates $x$, $y$ of the stars on the CCD frames was performed for the observations from Hoher List by standard CCD reduction techniques (DAOPHOT, IRAF) routines."
 Magnitudes. x and v positions were determined via PSF fit.," Magnitudes, $x$ and $y$ positions were determined via PSF fit."
 The observations from Calar Alto were reduced earlier (Geffert et 1994) by the IMEX routine of the IRAF program package., The observations from Calar Alto were reduced earlier (Geffert et 1994) by the IMEX routine of the IRAF program package.
 The astrometric reduction was performed by a central overlap algorithm., The astrometric reduction was performed by a central overlap algorithm.
 Due to the small field of the CCD frames we had to use only the plates in the first step of the reduction., Due to the small field of the CCD frames we had to use only the plates in the first step of the reduction.
 A catalogue of positions and proper motions of 450 stars was established in this first step., A catalogue of positions and proper motions of 450 stars was established in this first step.
 In the following steps of the reduction the CCD frames were included., In the following steps of the reduction the CCD frames were included.
 While for the plates only quadratic. polynomials of the rectangular. coordinates had to be taken into account. third order polynomials were necessary for the reduction of the CCD frames.," While for the plates only quadratic polynomials of the rectangular coordinates had to be taken into account, third order polynomials were necessary for the reduction of the CCD frames."
 The third order polynomials for the reduction of the CCD frames had to be used due to the distortion of the optics of the focal reducer of the WWFPP camera (Geffert et 1994)., The third order polynomials for the reduction of the CCD frames had to be used due to the distortion of the optics of the focal reducer of the WWFPP camera (Geffert et 1994).
 From the different position and time pairs we determined for each star for a certain epoch the mean position and the proper motion using least squares technique., From the different position and time pairs we determined for each star for a certain epoch the mean position and the proper motion using least squares technique.
 All stars with proper motion errors larger than 4 mas/yr were omitted., All stars with proper motion errors larger than 4 mas/yr were omitted.
 The final catalogue contains 532 positions and proper motions of stars in the region of M10., The final catalogue contains 532 positions and proper motions of stars in the region of M10.
 The median of the internal errors was about +1 mas/yr., The median of the internal errors was about $\plm 1\;$ mas/yr.
 From a plot of the proper motions versus magnitude no magnitude equation was found in our data., From a plot of the proper motions versus magnitude no magnitude equation was found in our data.
 We have performed two independent reductions of the M10 data with reference stars from Hippareos (ESA 1997) and the ACT catalogue (Urban et 1998)., We have performed two independent reductions of the M10 data with reference stars from Hipparcos (ESA 1997) and the ACT catalogue (Urban et 1998).
 Although the Hipparcos stars proper motions are more accurate. they seemed to be of only limited usefulness for our work. since the majority of the plates contained only four Hipparcos stars. while about 17 stars could be used from the ACT catalogue.," Although the Hipparcos stars proper motions are more accurate, they seemed to be of only limited usefulness for our work, since the majority of the plates contained only four Hipparcos stars, while about 17 stars could be used from the ACT catalogue."
 Therefore we consider both solutions as equivalent., Therefore we consider both solutions as equivalent.
 Table 3. (the complete table is available in. electronic. form) gives the catalogue of our positions and proper motions for the complete field with respect to the ACT catalogue., Table \ref{t_cdsdat} (the complete table is available in electronic form) gives the catalogue of our positions and proper motions for the complete field with respect to the ACT catalogue.
 The limiting magnitude of this catalogue is about V=15.5 mag. which corresponds to the limiting magnitude of the first epoch plates.," The limiting magnitude of this catalogue is about $V = 15.5\;$ mag, which corresponds to the limiting magnitude of the first epoch plates."
 The size of the field is approximately 75.\75aremin? centered on MIO., The size of the field is approximately $75\;\times\;75\;{\rm arcmin}^2$ centered on M10.
 We chose the ACT solution for the catalogue in Table 3.., We chose the ACT solution for the catalogue in Table \ref{t_cdsdat}. .
 The proper motions may be transfered to the Hipparcos system by adding Αμαcosó=1.5 mas/yr and Aja=|0.1 mas/yr to the proper motions from Table 3.., The proper motions may be transfered to the Hipparcos system by adding $\Delta\mua = -1.5\;$ mas/yr and $\Delta\mud = +0.1\;$ mas/yr to the proper motions from Table \ref{t_cdsdat}.
 For the determination of the membership we will also use the solution based on the ACT catalogue. while for the determination of the absolute proper motion of MIO we will take the mean of both solutions.," For the determination of the membership we will also use the solution based on the ACT catalogue, while for the determination of the absolute proper motion of M10 we will take the mean of both solutions."
 CCDs have been used in astrometry since several years. e.g. for the determination of parallaxes. double stars and for meridian circle observations (see references in Geffert 1998).," CCDs have been used in astrometry since several years, e.g. for the determination of parallaxes, double stars and for meridian circle observations (see references in Geffert 1998)."
 However most of these observations are based on CCD observations with fields of (10)? and smaller., However most of these observations are based on CCD observations with fields of $(10\arcmin)^2$ and smaller.
 Since our study uses CCD observations of fields with a size at least ς lit seems necessary to evaluate the accuracy. which may be obtained with such telescope detector combinations.," Since our study uses CCD observations of fields with a size at least $\;\times\;$ it seems necessary to evaluate the accuracy, which may be obtained with such telescope detector combinations."
 While the CCD frames of Calar Alto were already tested in an earlier study (Geffert et 1994). we will concentrate here on the Hoher List observations.," While the CCD frames of Calar Alto were already tested in an earlier study (Geffert et 1994), we will concentrate here on the Hoher List observations."
 In a first step we have compared positions of stars from pairs of CCD frames., In a first step we have compared positions of stars from pairs of CCD frames.
 The positions of one frame were transformed by an affinetransformation to à second frame and, The positions of one frame were transformed by an affinetransformation to a second frame and
the dust should be located at a coustaut radius so that when it is exposed to the QSO continu it equilibrates iV a constaut temperature.,the dust should be located at a constant radius so that when it is exposed to the QSO continuum it equilibrates at a constant temperature.
 This means that the dust is iu iv shell eeoimetzy: if there is dust at a rauge of radii. then it will radiate at lower temperatures than the sublimation temperature which will make the iufrared. SED wider aud flatter than a blackbody. typically rising iu flux towards longer wavelengths (Pier Rrolik 1992).," This means that the dust is in a shell geometry; if there is dust at a range of radii, then it will radiate at lower temperatures than the sublimation temperature which will make the infrared SED wider and flatter than a blackbody, typically rising in flux towards longer wavelengths (Pier Krolik 1992)."
 To quantify the coustraimt on ecomery. the ucar-infrared bunip was fit with the spectral nioel of Barvainis (1987).," To quantify the constraint on geometry, the near-infrared bump was fit with the spectral model of Barvainis (1987)."
 Barvainis’ model assumes that the dust is iu a shell subtending BiO steradiaus of the ceutral source with au iuner radius (04) set by the πανπιοαν of the source (CL) aud the (lust sublimation temperature (£144)., Barvainis' model assumes that the dust is in a shell subtending $\Omega$ steradians of the central source with an inner radius $r_1$ ) set by the luminosity of the source $L$ ) and the dust sublimation temperature $T_{\rm max}$ ).
 The shell is assumed o be optically thin in the near-IR (Az Tynan). with a UV optical depth of 3.," The shell is assumed to be optically thin in the near-IR $\lambda \ga 1\micron$ ), with a UV optical depth of 3."
" Barvainis’ uniform dust model was adopted with ο=(05s. Cluissivity proportional o pL, and we additionally assiune that the dust eraius ave a constant nuniber density. that £L=2«1019 cre/s. iuid that Zia LSOOT. fiue the imuer radius at 1.1 oc. consistent with the mieroleusime size coustraiut."," Barvainis' uniform dust model was adopted with $a=0.05 \micron$, emissivity proportional to $\nu^{1.7}$, and we additionally assume that the dust grains have a constant number density, that $L=2\times 10^{46}$ erg/s, and that $T_{\rm max}=1800$ K, fixing the inner radius at 1.1 pc, consistent with the microlensing size constraint."
 The wo free parameters left are ο/r4 (the ratio of the outer o inner radius of the shell) and Q., The two free parameters left are $r_2/r_1$ (the ratio of the outer to inner radius of the shell) and $\Omega$.
 Au adequate fit was ound for ro/rj=3 aud Q=πε., An adequate fit was found for $r_2/r_1 = 3$ and $\Omega = 7\pi/4$.
 In this model. the (lust outside ro is shielded. aud sodoes not contribute o the reprocessed light.," In this model, the dust outside $r_2$ is shielded, and sodoes not contribute to the reprocessed light."
 The eram number deusitv at he inuer edge is 5<10? 72 so the mnuplied dust uass is about [ML for a dust densitv of 1 ¢/an?.," The grain number density at the inner edge is $5\times 10^9$ $^{-3}$, so the implied dust mass is about $4 {\rm M_\odot}$ for a dust density of 1 $^3$."
 The Toomme Q-piriuueer is quite small. ~107 (for a mass 6of 2< LOPAL.). which indicates that this region should © eravitationally unstable.," The Toomre Q-parameter is quite small, $\sim 10^{-9}$ (for a mass of $2\times 10^9 M_\odot$ ), which indicates that this region should be gravitationally unstable."
 A clumpy disk model should lave a simular spectrum since the smaller filliug factor cau © conipeusated by an increase in dust muuber deusity., A clumpy disk model should have a similar spectrum since the smaller filling factor can be compensated by an increase in dust number density.
 The conclusion is hat the narrow infrared peak requires i Luarrow range of radii (factor of 3) so that reprocessed i1unuinositv is absorbed by a region with a narrow ranec of 6lust teniperature., The conclusion is that the narrow infrared peak requires a narrow range of radii (factor of 3) so that reprocessed luminosity is absorbed by a region with a narrow range of dust temperature.
 This narrow range is not uuprecedeuted Hw Barvainis (1992) found a narrow rauge of radii for the 6lust emission in Faivall 9 frou reverberation ryxi 0.3pc iuid ry~ 1.3pc., This narrow range is not unprecedented as Barvainis (1992) found a narrow range of radii for the dust emission in Fairall 9 from reverberation $r_1 \la 0.3$ pc and $r_2 \sim 1.3$ pc.
 This. of course. ix a rather simplistic model.," This, of course, is a rather simplistic model."
 N-ravs may penetrate deeper. heating dust at larger radii.," X-rays may penetrate deeper, heating dust at larger radii."
 Transicut heating bv N-ravs causes temperature flickering which broadens the spectrum: this is only iiportaut for dust grains with radi a<ομαι (Voit 1991)., Transient heating by X-rays causes temperature flickering which broadens the spectrum; this is only important for dust grains with radii $a \la 0.01\micron$ (Voit 1991).
 Iufrared radiation transfer has been ignored since the hell is optically thin by construction: reality may nof ο so sinple;, Infrared radiation transfer has been ignored since the shell is optically thin by construction; reality may not be so simple.
 A warped disk with a warp anele which Scales logarithinically racius would produce a uch wider infrared bump than observed., A warped disk with a warp angle which scales logarithmically radius would produce a much wider infrared bump than observed.
 Normal galaxies can be well described by a 3500 Iv dackbody in the near infrared spectral region (Scliuitt al., Normal galaxies can be well described by a 3500 K blackbody in the near infrared spectral region (Schmitt et al.
 1997). which peaks at a wavelcueth of about Linn (rest). while the Eimstein Cross spectrum peaks closer o 24nu (rest).," 1997), which peaks at a wavelength of about $1 \micron$ (rest), while the Einstein Cross spectrum peaks closer to $2 \micron$ (rest)."
 Thus. to fit the infrared bump requires Srong extinction as iu a starburst salaxv: an extinction of Ay~T ds needed to fit the infrared bump: most of 1e light is absorbed bv dust.," Thus, to fit the infrared bump requires strong extinction as in a starburst galaxy: an extinction of $A_V\sim 7$ is needed to fit the infrared bump; most of the light is absorbed by dust."
" The absorbed bpuuinositv is Las=6x1010 cress.es, Which mast be re-cradiated 1 jo far-infrared (FIR)."," The absorbed luminosity is $L_{\rm abs}=6\times 10^{46}$ erg/s, which must be re-radiated in the far-infrared (FIR)."
 The radio-EIR. relation. (Condo1 1992) Innit« the FIR Iuuiuosity to 6«10% ere/s. a factoY of 10 times smaller han the absorbed Iuninositv.," The radio-FIR relation (Condon 1992) limits the FIR luminosity to $6\times 10^{45}$ erg/s, a factor of 10 times smaller than the absorbed luminosity."
 A test of the host galaxy hwpothesis would be to look for a1 extended source im he IR images., A test of the host galaxy hypothesis would be to look for an extended source in the IR images.
 Meurer et al. (, Meurer et al. (
1997) have found an ορΊσα luit on the bolometric surface i+ . ⋝↥⋅↕∶↴∙⊾∐↑∐↸∖↴∖∷∖↴∪↕↴∖↴↑⋜∐⋅↴∏∐⋅↴∖↴↑∶↴∙⊾⋜↧↕⋜⋯↕↸∖↴∖↴∪↕⊇∖↓∩⊔∫⇀∙↨↘↴⋉⊳−∙↖↖⇁↕∐↸⊳↕ + . ↽⋅≽ ∙ would indicate that his starburst has a size of > 3.5 kpe. or 1 (includiug the factor of ~2 stretch due to lensing).,"1997) have found an empirical limit on the bolometric surface brightness of starburst galaxies of $2\times 10^{11}{\rm L}_\odot/
{\rm kpc}^2$ , which would indicate that this starburst has a size of $\ga$ 3.5 kpc, or $\sim1\arcsec$ (including the factor of $\sim 2$ stretch due to lensing)."
 This is huger than t1ο infrared upper huit of 0725. ruling out a starburst as the origin of the infrared bun.," This is larger than the infrared upper lmit of $0\farcs25$, ruling out a starburst as the origin of the infrared bump."
 One can roughly estimate the host galaxy bpuuinositv from the QSO luuinosity as follows., One can roughly estimate the host galaxy luminosity from the QSO luminosity as follows.
 Rauch Dlaudford (1991) have argued that sub-Eddiugton accretion disk models are too large to be consistent with the optical Size constraint from mücroleusiug., Rauch Blandford (1991) have argued that sub-Eddington accretion disk models are too large to be consistent with the optical size constraint from microlensing.
 We confiiued this bv fitting the optical/UV SED with various thin accretion disk models from IHubenuy et al. (, We confirmed this by fitting the optical/UV SED with various thin accretion disk models from Hubeny et al. (
2000): all disks are a factor of a few too big to be consistent with the mucrolensing constraint.,2000); all disks are a factor of a few too big to be consistent with the microlensing constraint.
 A sinaller ciission region müsht be achieved w decreasing the mass of the black hole: however. ouce he accretion rate reaches Eddington. the luuinosity will  roughly fixed at the Eddineton rate.," A smaller emission region might be achieved by decreasing the mass of the black hole; however, once the accretion rate reaches Eddington, the luminosity will be roughly fixed at the Eddington rate."
 Thus. if one ASSTILLCS hat the observed (isotropic) hunünositv equals he Eddiuston rate. then the black hole mass is about 241074 ," Thus, if one assumes that the observed (isotropic) luminosity equals the Eddington rate, then the black hole mass is about $2\times 10^9 {\rm M_\odot}$."
Using the Magorran ct al. (, Using the Magorrian et al. (
1998) relation ETWeCLL dack hole mass and bulge luminosity the bulge uninositv is estimated to be about 6«I0?L ... with au uncertaiutv of about a factor of 10.,"1998) relation between black hole mass and bulge luminosity, the bulge luminosity is estimated to be about $6\times 10^9 {\rm L_\odot}$ , with an uncertainty of about a factor of 10."
 This is more than 150 iues smaller than thenear-IR liuinesitv. vet another," This is more than 150 times smaller than thenear-IR luminosity, yet another"
into two main threads.,into two main threads.
 Phe first is the radial velocity. study and its results., The first is the radial velocity study and its results.
 Phe extraction of the velocities is described in Section 3. after which we discuss the summed. spectra and new cphemeris (Sections 4 and 5))., The extraction of the velocities is described in Section \ref{rad_vel} after which we discuss the summed spectra and new ephemeris (Sections \ref{disentangle} and \ref{ephemeris}) ).
 In Section 6 we discuss the absence of apparent cecentricity in our data. and conclude the svstem has changed in some wav.," In Section \ref{circularity} we discuss the absence of apparent eccentricity in our data, and conclude the system has changed in some way."
 We use our racial velocity semi-amplitude to derive new svstenm parameters in Section 7.., We use our radial velocity semi-amplitude to derive new system parameters in Section \ref{masses}.
 The second thread. isolating the lighteurve of the secondary star. is explored in Section 8S..," The second thread, isolating the lightcurve of the secondary star, is explored in Section \ref{flux_deficits}."
 Photometry and spectroscopy of LP Peg were obtained over four nights in 1995 October from the 2.5-m INT ancl 1.0-m JKT on La Palma., Photometry and spectroscopy of IP Peg were obtained over four nights in 1995 October from the 2.5-m INT and 1.0-m JKT on La Palma.
 Only tsvo nights of piotonmietryv were taken simultaneous with t1ο spectroscoov (see Section 2.4))., Only two nights of photometry were taken – simultaneous with the spectroscopy (see Section \ref{spec_phot}) ).
 Table 1. contains a log of the observations.," Table \ref{observations}
 contains a log of the observations."
 Data from the AAVSO show that in the latter half of 1995 LP Pee underwent two outbursts. one in mid-Septemboer and one in mid-December.," Data from the AAVSO show that in the latter half of 1995 IP Peg underwent two outbursts, one in mid-September and one in mid-December."
 The data presented in this paper were taken in the first week of October. approximately two weeks after the end of the mid-September outburst. and are thus during equiescence.," The data presented in this paper were taken in the first week of October, approximately two weeks after the end of the mid-September outburst, and are thus during quiescence."
 Data were taken with the ΝΤ on 1995 October 4th 5th using an EEVYT-chip with a scale of 0.31 arcsec per pixel., Data were taken with the JKT on 1995 October 4th 5th using an -chip with a scale of 0.31 arcsec per pixel.
 A narrow band (rwii-46A)) filter with maximunr ransmission at a wavelength of A7322.A was used on both nights., A narrow band ) filter with maximum transmission at a wavelength of $\lambda$ was used on both nights.
 The images were processed. by subtracting olf a single as image in the case of the Oct. 4th data. anc a median of three biases for the Oct Sth cata.," The images were processed by subtracting off a single bias image in the case of the Oct. 4th data, and a median of three biases for the Oct 5th data."
 A single sky Lat aken at the beginning of cach night was used to Llatliclel he images., A single sky flat taken at the beginning of each night was used to flatfield the images.
 The photometry was then extracted using an optimal extraction method. (Navlor 1998)., The photometry was then extracted using an optimal extraction method (Naylor 1998).
 The photometry of LP Peg was divided by that of a neighbouring star (the second star placed on the slit during the spectroscopic runsee below) to remove the ellects of sky transparcney variations and. thus. the photometry is relative.," The photometry of IP Peg was divided by that of a neighbouring star (the second star placed on the slit during the spectroscopic run—see below) to remove the effects of sky transparency variations and, thus, the photometry is relative."
 Phe phase-Folded light curves for both nights are shown in Figs., The phase-folded light curves for both nights are shown in Figs.
 1. and 2 with the bright-spot hump at ó0.2 and the primary eclipse at ó=0.0.," \ref{4th} and \ref{5th}
 with the bright-spot hump at $\phi\sim -0.2$ and the primary eclipse at $\phi=0.0$."
 Data were taken on 1995 October 4th. 5th. 9th and. 10th with the INT. using a rEKJ-chip with a scale of 0.33 aresec »er pixel. but binned by a factor two in the spatial direction.," Data were taken on 1995 October 4th, 5th, 9th and 10th with the INT, using a -chip with a scale of 0.33 arcsec per pixel, but binned by a factor two in the spatial direction."
 AX two-aresee slit was used set at PA 58.7 degrees to obtain he spectrum ofa neighbouring star simultaneously (the star used to divide the IP Peg photometry bysee above)., A two-arcsec slit was used set at PA 58.7 degrees to obtain the spectrum of a neighbouring star simultaneously (the star used to divide the IP Peg photometry by—see above).
 Εις allowed: us to correct for slit losses (see Section 2.4))., This allowed us to correct for slit losses (see Section \ref{spec_phot}) ).
 The 15211. grating was used on the 4th and Sth. giving ~1.22 oper pixel: the 11120011 grating was used on the Oth and 1011. giving —0.84 pper pixel.," The R831R grating was used on the 4th and 5th, giving $\sim$ 1.22 per pixel; the R1200R grating was used on the 9th and 10th, giving $\sim$ 0.84 per pixel."
 An are spectrum was taken each time the telescope's position on the skv was changed (7 every hour)., An arc spectrum was taken each time the telescope's position on the sky was changed $\sim$ every hour).
 The nights of the 4th and Sth were generally clear but the nights of the 9th and LOth were alfected by cloud. especially so the data of the 9th.," The nights of the 4th and 5th were generally clear but the nights of the 9th and 10th were affected by cloud, especially so the data of the 9th."
 An exposure time of 600s was used on the 4th. Sth and 9th: 400s on the 10th.," An exposure time of 600s was used on the 4th, 5th and 9th: 400s on the 10th."
 The images for cach night were first. bias subtracted., The images for each night were first bias subtracted.
 Tungsten [latis taken cach night were then used. after dividing by fitted low-order polynomials. to Gatfick the images.," Tungsten flats taken each night were then used, after dividing by fitted low-order polynomials, to flatfield the images."
 Extraction of the spectra was performed using an optimal extraction method. (Horne 1986. Robertson 1986). and then wavelength calibration was done using low-order polynomial fits to the appropriate are spectra.," Extraction of the spectra was performed using an optimal extraction method (Horne 1986, Robertson 1986), and then wavelength calibration was done using low-order polynomial fits to the appropriate arc spectra."
to a low confidence level for any result drawn from statistical studies. while radio or optical incompleteness Heads to biases in the determination of the recshilt distribution of radio sources at such low-Llux levels.,"to a low confidence level for any result drawn from statistical studies, while radio or optical incompleteness leads to biases in the determination of the redshift distribution of radio sources at such low-flux levels."
 In this Paper we present an analysis of the properties and angular clustering of ~ 4100 radio sources brighter han 1l mJwv with optical counterparts brighter 6;<22., In this Paper we present an analysis of the properties and angular clustering of $\sim$ 4100 radio sources brighter than 1 mJy with optical counterparts brighter $b_J \le 22$.
 This sample was obtained by matching together objects rom the FIRST and APM surveys over 350) square degrees near the celestial equator., This sample was obtained by matching together objects from the FIRST and APM surveys over $\sim 350$ square degrees near the celestial equator.
 Despite not. having redshift measurements for these sources. their joint. radio. ;whotometric ancl morphological properties. can be used o infer extremely useful information on their nature.," Despite not having redshift measurements for these sources, their joint radio, photometric and morphological properties can be used to infer extremely useful information on their nature."
" Furthermore. for by)x21.5. we obtain a catalogue of sources which is ~100% complete in the optical band (ALA2000) and complete in the radio band. (with a completeness level rising to for radio Huxes 5,4c6g;c3 ody. Becker et al."," Furthermore, for $b_J \le
21.5$, we obtain a catalogue of sources which is $\sim 100\%$ complete in the optical band (MA2000) and complete in the radio band (with a completeness level rising to for radio fluxes $S_{1.4 {\rm
GHz}}\ge 3$ mJy, Becker et al.,"
 1995)., 1995).
 The photometric and. morphological properties of the optical identifications are then used to divide the whole sample of radio sources into two well defined populations., The photometric and morphological properties of the optical identifications are then used to divide the whole sample of radio sources into two well defined populations.
 The first population consists of Iow-redshift. radio galaxies and the second one of high-z QSO and. objects with no counterpart on the VAST plates., The first population consists of low-redshift radio galaxies and the second one of high-z QSO and objects with no counterpart on the UKST plates.
 The homogeneity and completeness of the low-z sample allows us to study. the clustering properties of these sources., The homogeneity and completeness of the low-z sample allows us to study the clustering properties of these sources.
 Lastlv. the area chosen for our analvsis coincides with some of the fields observed. in the. 2d£ Galaxy Redshift Survey (Maddox 1998. Colless. 1999).," Lastly, the area chosen for our analysis coincides with some of the fields observed in the 2df Galaxy Redshift Survey (Maddox 1998, Colless, 1999)."
 Phis sample of optical identifications of  4100 radio sources therefore provides an excellent starting point for further wicle-area spectroscopic follow-up., This sample of optical identifications of $\sim$ 4100 radio sources therefore provides an excellent starting point for further wide-area spectroscopic follow-up.
 The lavout of the paper is as follows: Section 2 brielly describes the two surveys and the data coming from them., The layout of the paper is as follows: Section 2 briefly describes the two surveys and the data coming from them.
 Section 3 presents the procedure we adopted to match radio and optical sources together. while Section 4 is devoted to the analvsis of the photometric properties of the optical identifications.," Section 3 presents the procedure we adopted to match radio and optical sources together, while Section 4 is devoted to the analysis of the photometric properties of the optical identifications."
 Section 5r examines the clustering properties of the sample ancl Section 6 summarizes our conclusions., Section 5 examines the clustering properties of the sample and Section 6 summarizes our conclusions.
 Throughout the paper we will assume Qu=0.4. ο=0.65. A=0.6.," Throughout the paper we will assume $\Omega_0=0.4$, $h_0=0.65$, $\Lambda=0.6$."
 The FIRST (Faint Images of the Racio Sky at Twenty centimetres) survey (Becker. White and Holfand. 1995) began in the spring of 1993 ancl will eventually cover sonic 10.000 square degrees of the sky in the north Galactic cap and equatorial zones.," The FIRST (Faint Images of the Radio Sky at Twenty centimetres) survey (Becker, White and Helfand, 1995) began in the spring of 1993 and will eventually cover some 10,000 square degrees of the sky in the north Galactic cap and equatorial zones."
 Phe beamesize is 5.4 aresee at 1.4 Cillz. with an rms sensitivity of typically 0.15 mv/beam.," The beam-size is 5.4 arcsec at 1.4 GHz, with an rms sensitivity of typically 0.15 mJy/beam."
 A map is produced for each field ancl sources are detected: using an elliptical Gaussian fitting procedure (White et al..," A map is produced for each field and sources are detected using an elliptical Gaussian fitting procedure (White et al.,"
 1997): the 5o source detection limit is 1 mv., 1997); the $5\sigma$ source detection limit is $\sim 1$ mJy.
 The astrometric reference [rame of the maps is accurate to 0.05 aresec. and individual sources have 90 per cent confidence error. circles of radius « 0.5 arcsec at the 3 mJy level. and 1 aresee a the survey threshold.," The astrometric reference frame of the maps is accurate to 0.05 arcsec, and individual sources have 90 per cent confidence error circles of radius $<$ 0.5 arcsec at the 3 mJy level, and 1 arcsec at the survey threshold."
 Phe surface density of objects in the catalogue is ~90 per square degree. though this is reduce to ~SO per square degree if. we combine multi-componen sources (Magliocchetti et al.," The surface density of objects in the catalogue is $\sim 90$ per square degree, though this is reduced to $\sim 80$ per square degree if we combine multi-component sources (Magliocchetti et al.,"
 1998)., 1998).
 The depth. uniformity and angular extent of the survey are excellent attributes for investigating. amongst others. the clustering properties of faint sources.," The depth, uniformity and angular extent of the survey are excellent attributes for investigating, amongst others, the clustering properties of faint sources."
 The catalogue derived. [rom this survey has been estimated to be 95 per cent complete at 2 mJ» and SO per cent complete at 1 mJy (Becker et al..," The catalogue derived from this survey has been estimated to be 95 per cent complete at 2 mJy and 80 per cent complete at 1 mJy (Becker et al.,"
. 1995)., 1995).
 We used the 5 July 2000 version of the catalogue which contains approximately 722.354 sources from the north and south Galactic caps.," We used the 5 July 2000 version of the catalogue which contains approximately 722,354 sources from the north and south Galactic caps."
" This is derived from the 1993 through 2000 observations that cover nearly TOSS square degrees of skv. and includes most of the area 1720"". 22.27€Dee<57.5% and 2120""< 3P20"" 2:85Dee<2.2"","," This is derived from the 1993 through 2000 observations that cover nearly 7988 square degrees of sky, and includes most of the area $7^h20^m \simlt {\rm RA}(2000) \simlt 17^h20^m$ , $22.2^\circ \simlt
{\rm Dec} \simlt 57.5^\circ$ and $21^h20^m \simlt {\rm RA}(2000)
\simlt 3^h20^m$ , $-2.8^\circ \simlt {\rm Dec} \simlt 2.2^\circ$."
 The APM galaxy survey has been extensively described in Alacddox ct al. (," The APM galaxy survey has been extensively described in Maddox et al., ("
1990a. 19000. 1996).,"1990a, 1990b, 1996)."
" Driellv. it is based on APM scans of 5.875.8"" UINST plates. covering about 7000 square degrees in the regions Dec <—207 in the southern Galactic cap and —11.57< Dec z2.57. 7.57< Dec 2.5"" respectively in the equatorial sep and ngp."," Briefly, it is based on APM scans of $5.8^\circ\times 5.8^\circ$ UKST plates, covering about 7000 square degrees in the regions Dec $<-20^\circ$ in the southern Galactic cap and $-17.5^\circ\le$ Dec $\le 2.5^\circ$, $-7.5^\circ\le$ Dec $
\le 2.5^\circ$ respectively in the equatorial sgp and ngp."
 A scan of a typical plate records about 300.000 images. with a limiting magnitude for image detection by£22.," A scan of a typical plate records about 300,000 images, with a limiting magnitude for image detection $b_J \simlt 22$."
 For each image. the measurcments show a positional accuracy of ~| aresec ancl an isophotal magnitude accuracy of ~0.10.2 mag.," For each image, the measurements show a positional accuracy of $\sim 1$ arcsec and an isophotal magnitude accuracy of $\sim 0.1-0.2$ mag."
 The photometry from each plate is corrected so that it is consistent with the neighbouring plates. in the way described in Maddox. ct al. (," The photometry from each plate is corrected so that it is consistent with the neighbouring plates, in the way described in Maddox et al. ("
1990b).,1990b).
. Phe limit for uniform image, The limit for uniform image
For all 27 models the majority of star clusters merge into a stable object.,For all 27 models the majority of star clusters merge into a stable object.
 The turnover iu the reg Vs. Mog space (Figure T.. 8 aud 12)) leads to degenerate states. because a relatively compact CC can produce a comparable merger object as a more massive CC having a significantly larger CC size.," The turnover in the $r_{\rm eff}$ vs. $M_{\rm encl}$ space (Figure \ref{figmassreff}, \ref{turnover} and \ref{obssims}) ) leads to degenerate states, because a relatively compact CC can produce a comparable merger object as a more massive CC having a significantly larger CC size."
 In consequence. a range of initial couclitious cau form a merger object comparable to 22119 preventing us to pinpoint the parameters of the original CC. which formed 22119.," In consequence, a range of initial conditions can form a merger object comparable to 2419 preventing us to pinpoint the parameters of the original CC, which formed 2419."
 On the other hand. the larger the range of initial conditions that end up in à 22[19-like object. the larger is the probability of creating a massive EC like 221190.," On the other hand, the larger the range of initial conditions that end up in a 2419-like object, the larger is the probability of creating a massive EC like 2419."
 As the individual star clusters of the CC formed at approximately the same time [rom molecular clouds of a galaxy. the observed abseuce of multiple stellar populations (Ripepietal.2007) is fully consistent with our model.," As the individual star clusters of the CC formed at approximately the same time from molecular clouds of a galaxy, the observed absence of multiple stellar populations \citep{ripepi} is fully consistent with our model."
 A small scatter in metallicity would also be explained. as the individual pre-cluster cloud cores of the complex could have had slightly dillerent metallicities.," A small scatter in metallicity would also be explained, as the individual pre-cluster cloud cores of the complex could have had slightly different metallicities."
 La acdcitiou. some stars from the parent galaxy. might have been captured by 22119 during its formation as demonstrated for massive star clusters like & Cen by Fellhaueretal.(2006).," In addition, some stars from the parent galaxy might have been captured by 2419 during its formation as demonstrated for massive star clusters like $\omega$ Cen by \cite{fellhauer06}."
. Observations demoustrate that massive complexes of star clusters nowadays. precomiuautly form during gravitational encouuters between late-type galaxies., Observations demonstrate that massive complexes of star clusters nowadays predominantly form during gravitational encounters between late-type galaxies.
 Youug massive CCs with masses above 106AL. and sizes of a few hundred pe. comparable to those used as initial conditions for the numerical simulations presented in Sect. 3..," Young massive CCs with masses above $^6\,M_\odot$ and sizes of a few hundred pc, comparable to those used as initial conditions for the numerical simulations presented in Sect. \ref{results},"
 have been observed in the Antennae galaxies and in 9922 (Pellerinetal.2010)., have been observed in the Antennae galaxies \citep{bastian06} and in 922 \citep{pellerin}.
. Since galaxy-galaxy mergers are anticipated to have been much more common curing early hierarchical structure formation it is expected that star formation in cluster complexes has been a sienificaut star formation mode during early cosmological epochs., Since galaxy-galaxy mergers are anticipated to have been much more common during early hierarchical structure formation it is expected that star formation in cluster complexes has been a significant star formation mode during early cosmological epochs.
 As stellar structures in the outer halo are lone lived features. renminauts of a parent galaxy or stellar tidal streams may still be observable in the Milky Way halo.," As stellar structures in the outer halo are long lived features, remnants of a parent galaxy or stellar tidal streams may still be observable in the Milky Way halo."
 Newbergetal.(2003) found. an over-cleusity of A-type stars at a distance of 83 to 85 kpe. which has a width of at least LO° and which was traced for more thau 20° on the sky.," \cite{newberg03} found an over-density of A-type stars at a distance of 83 to 85 kpc, which has a width of at least $^\circ$ and which was traced for more than $^\circ$ on the sky."
 22119 is located within this debris (on the sky aud at the same distance)., 2419 is located within this debris (on the sky and at the same distance).
 Newbergetal.(2003) argued that 22119 aud the stellar over-deusity might be associated with the Sagittarius dwarf galaxy. as 22119 lies only 13 kpe from its orbital plane.," \cite{newberg03} argued that 2419 and the stellar over-density might be associated with the Sagittarius dwarf galaxy, as 2419 lies only 13 kpc from its orbital plane."
 However. the stellar streams from Ser dwarf have mean metallicities between [Fe/H] = -0.1 near the core and -1.1 within the arms 2007).. while 22119 has a very low metallicity of [Fe/H] = —2.12 (Harris 1996).," However, the stellar streams from Sgr dwarf have mean metallicities between [Fe/H] = –0.4 near the core and –1.1 within the arms \citep{chou07}, while 2419 has a very low metallicity of [Fe/H] = –2.12 \citep{harris}. ."
. The Virgo Stellar Stream (VSS) has a metallicity of [Fe/H] = —1.86 with a large scatter of 0.10 dex (Dutlauetal.2006).. which makes it comparable to the metallicity of 22119.," The Virgo Stellar Stream (VSS) has a metallicity of [Fe/H] = –1.86 with a large scatter of 0.40 dex \citep{duffau}, which makes it comparable to the metallicity of 2419."
 cleterimined the proper motion of the VSS and calculated an orbit with a pericentric distance of 11 kpe and au apocentric distance of 89kpc., \cite{casetti} determined the proper motion of the VSS and calculated an orbit with a pericentric distance of 11 kpc and an apocentric distance of 89kpc.
 They concluded that the current, They concluded that the current
Fig.,Fig.
 11 shows the results of the uew method., \ref{plot4} shows the results of the new method.
 We lave removed the prefiltered FITS methocl [rom the plot (shown in Fig. 7)), We have removed the prefiltered FITS method from the plot (shown in Fig. \ref{plot2C}) )
 so that we may concentrate οw attention ou the faster methods. namely. the two sequence file methods previously described (tle first two bars shown in Fie. 11)).," so that we may concentrate our attention on the faster methods, namely, the two sequence file methods previously described (the first two bars shown in Fig. \ref{plot4}) ),"
 aud two uew metliods. oue each for performing SQL agaiust tlie iustruetured sequence file database (the third bar) and for performiug SQL against the structure sequence file database (the fourth bar).," and two new methods, one each for performing SQL against the unstructured sequence file database (the third bar) and for performing SQL against the structured sequence file database (the fourth bar)."
 Several patterus are observed in the plot., Several patterns are observed in the plot.
 First. uote tha it is only meaniugful to compare SQL vs. nouSQL for a given sequence file database. either structured or unstructured.," First, note that it is only meaningful to compare SQL vs. nonSQL for a given sequence file database, either structured or unstructured."
 Thus. the appropriate comuparisons to make are between bars one aud tliree or between bars two and four.," Thus, the appropriate comparisons to make are between bars one and three or between bars two and four."
 We observe that in such comparisons. the SQL method does successfully outperform the nonsSQL method in most cases. but to a lower degree than we hoped for when we implemented it. bar three shows only a small improvement over bar oue aud bar [οιw shows virtually no improvement over bay iwo.," We observe that in such comparisons, the SQL method does successfully outperform the nonSQL method in most cases, but to a lower degree than we hoped for when we implemented it, bar three shows only a small improvement over bar one and bar four shows virtually no improvement over bar two."
 Let us cousicer the results iucdicated by bars two aud four for the larger query., Let us consider the results indicated by bars two and four for the larger query.
 The noteworthy observatiou is that wile SQL shows some benelit. that benefit is lower thau originally anticipated eiven tlje 2.5x clillerence in mapper input records (FITS files). 1:3335 vs. 3885.," The noteworthy observation is that while SQL shows some benefit, that benefit is lower than originally anticipated given the 3.5x difference in mapper input records (FITS files), 13335 vs. 3885."
 13335 is the nunber of FITS files reac iu by the prefiltered sequence file method (of which 3885 were relevant and the ‘est cliscarded) while 3885 is tlie uumber of FITS files read in by the SQL method (all of which were relevant)., 13335 is the number of FITS files read in by the prefiltered sequence file method (of which 3885 were relevant and the rest discarded) while 3885 is the number of FITS files read in by the SQL method (all of which were relevant).
 Τjerefore. we can conclude that in this case the cost of considering aud cdiscardiug uunerous relevant FITS files was negligible aud likewise that he adcditioual complexity innposed yw supporting aud sing an external SQL database ollers uo bejefit.," Therefore, we can conclude that in this case the cost of considering and discarding numerous irrelevant FITS files was negligible and likewise that the additional complexity imposed by supporting and using an external SQL database offers no benefit."
 Let us consider the results indicated by bars one aud three for the larger query., Let us consider the results indicated by bars one and three for the larger query.
 Again. we note ittle provement in performance given the variance of the couidence intervals.," Again, we note little improvement in performance given the variance of the confidence intervals."
 This is surprising eiven the dramatic dillereuce in mapper input records. a [ul 26 fold difference.," This is surprising given the dramatic difference in mapper input records, a full 26 fold difference."
 The first bar 'epreseits. non-prefiterecl uustructured sequence files. thus 100.058 input records while the third or represents a 5€JL iuethod. thus 3885 input records.," The first bar represents non-prefiltered unstructured sequence files, thus 100,058 input records while the third bar represents a SQL method, thus 3885 input records."
 The corclusion is similar iu this case. that he cos of cliscardi& irrelevant. [iles is low.," The conclusion is similar in this case, that the cost of discarding irrelevant files is low."
 We might the ask: why is there a two-fold difference in performauce between the two SQL uethocs cousiceritig that they both processed exactly the same amount of data?, We might then ask: why is there a two-fold difference in performance between the two SQL methods considering that they both processed exactly the same amount of data?
 To auswer this questio] we lust investigate not merely the ununber of mapper input records but the πο of lauuched by Hadoop., To answer this question we must investigate not merely the number of mapper input records but the number of launched by Hadoop.
 These two values are rarely equal because a mapper cau be reusecl O process 1jultiple input records., These two values are rarely equal because a mapper can be reused to process multiple input records.
 In the SQL-unustructu'ed-secqueuce-file inetlod. the 3885 input FITS files were processed by 17114 mapper objects (about two FITS files each) while in the SQL-stπιοurede-sequenuce-file method. 338 mapper objects were 1sed (about eleven FITS files each).," In the SQL-unstructured-sequence-file method, the 3885 input FITS files were processed by 1714 mapper objects (about two FITS files each) while in the SQL-structured-sequence-file method, 338 mapper objects were used (about eleven FITS files each)."
 This discrepaucey is due to the way we assign FITS files to eacl mapper object., This discrepancy is due to the way we assign FITS files to each mapper object.
 For each mapper object. we assign FITS files in the same HDFS block.," For each mapper object, we assign FITS files in the same HDFS block."
 Due to the replication of blocks across HDFS. copies of each block will be stored ou multiple hosts.," Due to the replication of blocks across HDFS, copies of each block will be stored on multiple hosts."
 When possible. Hadoop will schedule mapper objects to run on one of the hosts where that inappers input blocks are located so that the files," When possible, Hadoop will schedule mapper objects to run on one of the hosts where that mapper's input blocks are located so that the files"
The observations of radio light curves give an estimate of the mass loss rate of the progenitor system of the supernova.,The observations of radio light curves give an estimate of the mass loss rate of the progenitor system of the supernova.
 The radio light curve of the type ΠΟ SN 2001gd indicates a mass loss rate of about 2-12 x 107 Mayr! (22).. which is in between the rates for typical type II and type Ib SNe.," The radio light curve of the type IIb SN 2001gd indicates a mass loss rate of about 2–12 $\times$ $10^{-5}$ $_\odot$ $^{-1}$ \citep[]{Stockdale03, Perez-Torres05}, which is in between the rates for typical type II and type Ib SNe."
 Another supernova of type IIb. SN 200115. was observed in 2001 in NGC 7424 and shows a spectral evolution similar to that of SN 1993J (?)..," Another supernova of type IIb, SN 2001ig, was observed in 2001 in NGC 7424 and shows a spectral evolution similar to that of SN 1993J \cite[]{Ryder06}."
 Evidence was found for a star of spectral type late-B through late-F at the location of SN 20011g. a possible companion of the progenitor of SN 200lig (2)..," Evidence was found for a star of spectral type late-B through late-F at the location of SN 2001ig, a possible companion of the progenitor of SN 2001ig \cite[]{Ryder06}."
 SN 2003be evolved from a type Ic supernova to a rich type IIb. to a hydrogen-poor type [be (?)..," SN 2003bg evolved from a type Ic supernova to a hydrogen-rich type IIb, to a hydrogen-poor type Ibc \cite[]{Soderberg06}."
" It was observed as a broad-lined type Hb supernova and proclaimed to be ""the first type IIb hypernova’ (22).."," It was observed as a broad-lined type IIb supernova and proclaimed to be 'the first type IIb hypernova' \citep[]{Mazzali09, Hamuy09}. ."
 The broadness of the lines indicates a high progenitor mass (2).., The broadness of the lines indicates a high progenitor mass \cite[]{Hamuy09}.
 The light curve and spectral evolution indicate the presence of a thin layer of hydrogen at time of explosion. =0.05 (2)..," The light curve and spectral evolution indicate the presence of a thin layer of hydrogen at time of explosion, $\approx 0.05$ \cite[]{Mazzali09}."
 The velocity ofthe ejecta resembles more closely the velocity of SNe type Ib than type II (2).., The velocity ofthe ejecta resembles more closely the velocity of SNe type Ib than type II \cite[]{Soderberg06}.
 This implies a compact Wolf-rayet progenitor. with a progenitor mass between 20 and 25," This implies a compact Wolf-rayet progenitor, with a progenitor mass between 20 and 25."
M... ? conclude that this event is an intermediate case between Se type IIb and type Ib., \cite{Soderberg06} conclude that this event is an intermediate case between SNe type IIb and type Ib.
 The light curve of SN 2008ax shows some differences with the light curve of SN 1993J. namely the lack of the first peak and it has slightly bluer colors (2)..," The light curve of SN 2008ax shows some differences with the light curve of SN 1993J, namely the lack of the first peak and it has slightly bluer colors \cite[]{Pastor08}."
 These features can be explained by a less massive hydrogen envelope at time of explosion. less than a few » 0.1 in comparison with the progenitor of SN 1993J (?)..," These features can be explained by a less massive hydrogen envelope at time of explosion, less than a few $\times$ 0.1 $_\odot$, in comparison with the progenitor of SN 1993J \cite[]{Crocket08}."
 Cas A is the supernova remnant of a star that exploded about 350 years ago (?).., Cas A is the supernova remnant of a star that exploded about 350 years ago \cite[]{Thorstensen01}.
 A light echo from this explosion (?) shows evidence that it was a supernova of type IIb., A light echo from this explosion \cite[]{Krause08} shows evidence that it was a supernova of type IIb.
" Direct methods to determine the mass of the progenitor star are difficult. but the ejecta mass was calculated to be 2—4 M,. and the remnant would be expected to be a neutron star with a mass between 1.5 and 2.2 M. (2).."," Direct methods to determine the mass of the progenitor star are difficult, but the ejecta mass was calculated to be 2–4 $_\odot$ and the remnant would be expected to be a neutron star with a mass between 1.5 and 2.2 $_\odot$ \cite[]{Young06}."
 This sets the mass of the star at time of its explosion at about 4-6 M..., This sets the mass of the star at time of its explosion at about 4–6 $_\odot$.
 There is no direct evidence as to wether this supernova was of the compact or extended type IIb. but the possibilty that the progenitor was a red supergiant is left open.," There is no direct evidence as to wether this supernova was of the compact or extended type IIb, but the possibilty that the progenitor was a red supergiant is left open."
 For this supernova single and binary progenitor models were calculated (?).., For this supernova single and binary progenitor models were calculated \cite[]{Young06}.
 The single star models indicated fine-tuning of the stellar wind is necessary to evolve to the specific characteristics of the supernova remnant Cas A (?).., The single star models indicated fine-tuning of the stellar wind is necessary to evolve to the specific characteristics of the supernova remnant Cas A \cite[]{Young06}.
 Besides. there is evidence that the progenitor could only have had a very short-lived Wolf-rayet phase. which is difficult to explain with single stars (22)..," Besides, there is evidence that the progenitor could only have had a very short-lived Wolf-rayet phase, which is difficult to explain with single stars \citep[]{Schure08, VanVeelen09}."
 There is no evidence for à companion star., There is no evidence for a companion star.
 Therefore à common envelope scenario was proposed. in which the two stars merge into a single star before explosion.," Therefore a common envelope scenario was proposed, in which the two stars merge into a single star before explosion."
 Observations show tentative evidence for this scenario (?).. such as the asymmetric distribution of the quasi-stationary flocculi near Cas A. which could arise from the loss of a common TThe observations put constraints on the general properties of a SN type IIb: the explosion of a supergiant with a hydrogen envelope mass between 0.1 and 0.5 M...," Observations show tentative evidence for this scenario \cite[]{Krause08}, such as the asymmetric distribution of the quasi-stationary flocculi near Cas A, which could arise from the loss of a common The observations put constraints on the general properties of a SN type IIb: the explosion of a supergiant with a hydrogen envelope mass between 0.1 and 0.5 $_\odot$."
 We consider the lower limit of the mass of the hydrogen envelope to be 0.1. rather than 0.01M... the lower limit proposed by ?..," We consider the lower limit of the mass of the hydrogen envelope to be 0.1 rather than 0.01, the lower limit proposed by \cite{Chevalier+Soderberg2010}."
 The explosion of a star with a hydrogen envelope mass smaller than 0.1. will exhibit hydrogen lines in its spectrum. but only in the early phases of the supernova.," The explosion of a star with a hydrogen envelope mass smaller than 0.1 will exhibit hydrogen lines in its spectrum, but only in the early phases of the supernova."
 In addition the light curve resembles the typical light curve of a SN type Ib., In addition the light curve resembles the typical light curve of a SN type Ib.
 Therefore such a supernova will more likely be defined as a supernova type Ib or a transitional type between Hb and Ib (see examples above. e.g. SN 2003056 and SN 2000H).," Therefore such a supernova will more likely be defined as a supernova type Ib or a transitional type between IIb and Ib (see examples above, e.g. SN 2003bg and SN 2000H)."
 ?. places the upper limit of the mass of the hydrogen envelope of the progenitor of a type Ib supernova at 0.1., \cite{Elmhamdi06} places the upper limit of the mass of the hydrogen envelope of the progenitor of a type Ib supernova at 0.1.
M... A binary progenitor is confirmed or considered likely in some cases. and the secondary has been detected as a blue supergiant in possibly two cases.," A binary progenitor is confirmed or considered likely in some cases, and the secondary has been detected as a blue supergiant in possibly two cases."
 However. no general constraints on the secondary can be set.," However, no general constraints on the secondary can be set."
 We use a version of the binary evolution code STARS originally developed by ? and later updated and described by various. authors (e.g.?2?)..," We use a version of the binary evolution code STARS originally developed by \citet{Eggleton71} and later updated and described by various authors \citep[e.g.][]{Pols+95, EggletonsBook06, Glebbeek+08}."
 The code is fully implicit and solves the equations for the structure and composition of the star simultaneously., The code is fully implicit and solves the equations for the structure and composition of the star simultaneously.
 It employs an adaptive non-Lagrangian mesh that places mesh points in regions of the star where higher resolution is required., It employs an adaptive non-Lagrangian mesh that places mesh points in regions of the star where higher resolution is required.
 This allows us to evolve stars with a reasonable accuracy using as few as 200 mesh points., This allows us to evolve stars with a reasonable accuracy using as few as 200 mesh points.
 The code therefore is fast and suitable to compute the large numbers of models needed to investigate wide initial. parameter space of binary systems (e.g.?).., The code therefore is fast and suitable to compute the large numbers of models needed to investigate wide initial parameter space of binary systems \citep[e.g.][]{DeMink+07}.
 We use nuclear reaction rates from ?» and ? and opacities from ? and ?.., We use nuclear reaction rates from \citet{Caughlan+1985} and \citet{Caughlan+Fowler1988} and opacities from \citet{Rogers+Iglesias1992} and \citet{Alexander+Ferguson1994}.
 The assumed heavy-element composition is scaled to solar abundances (?).., The assumed heavy-element composition is scaled to solar abundances \citep{Anders+Grevesse1989}.
 Convection is implemented using a diffusion approximation (?) of the mixing-length theory (2).. assuming a mixing length of 2.0 pressure scale heights.," Convection is implemented using a diffusion approximation \citep{Eggleton72} of the mixing-length theory \citep{Boehm-Vitense1958}, assuming a mixing length of 2.0 pressure scale heights."
 We use the Schwarzschil¢ criterion to determine the boundaries of the convective regions., We use the Schwarzschild criterion to determine the boundaries of the convective regions.
" Convective overshooting is taken into account using the prescription of ? with an overshooting parameter of 04, = 0.12. which was calibrated against accurate stellar data eclipsing binaries (?).."," Convective overshooting is taken into account using the prescription of \cite{Schroder97} with an overshooting parameter of $\delta_{\rm ov}$ = 0.12, which was calibrated against accurate stellar data eclipsing binaries\citep{Pols+1997}. ."
 In terms of the pressure scale height. as the overshooting parameter is commonly defined in other stellar evolution codes. this value approximately compares to a= 0.25.," In terms of the pressure scale height, as the overshooting parameter is commonly defined in other stellar evolution codes, this value approximately compares to $\alpha_{\rm ov} \approx 0.25$ ."
determined. among other things. bx the nuclear partition function.,"determined, among other things, by the nuclear partition function."
 In turn. the equilibrium nuclear composition is determined by (he competition between beta decay. and. electron capture.," In turn, the equilibrium nuclear composition is determined by the competition between beta decay and electron capture."
 If the partition functions used to estimate the composition at a given } do not match with the partition functions used to calculate the weak rates. the caleulated equilibrium of the svtem will not be correct.," If the partition functions used to estimate the composition at a given $Y_e$ do not match with the partition functions used to calculate the weak rates, the calculated equilibrium of the sytem will not be correct."
 The authors acknowledge helpful correspondence wilh G. Martinez-Pinedo and Ix. Langanke regarding (heir calculations., The authors acknowledge helpful correspondence with G. Martinez-Pinedo and K. Langanke regarding their calculations.
 We also thank Rob Iollnana for useful discussions regarding the rp-process., We also thank Rob Hoffman for useful discussions regarding the -process.
 This work was partially supported by the DOE Program for Scientilic Discovery through Advanced Computing (5c3DAC) al UCSD and LLNL. and by NSF grant. al UCSD., This work was partially supported by the DOE Program for Scientific Discovery through Advanced Computing (SciDAC) at UCSD and LLNL and by NSF grant PHY-00-0099499 at UCSD.
" A portion of (this work was performed under (he auspices of the U.S. Department of Energy by University of California Lawrence Livermore Laboratory. wider contract W-7405-ENG--48,", A portion of this work was performed under the auspices of the U.S. Department of Energy by University of California Lawrence Livermore Laboratory under contract W-7405-ENG-48.
"remains even for a change of ten times the standard deviation in opposite directions for f, and f» and for arbitrary changes in epoch of phase zero.",remains even for a change of ten times the standard deviation in opposite directions for $f_1$ and $f_2$ and for arbitrary changes in epoch of phase zero.
 The result is clearly robust to observational errors. but it is not clear what physical conclusions may be derived. from. this result.," The result is clearly robust to observational errors, but it is not clear what physical conclusions may be derived from this result."
 The behaviour is certainly not random and must have a physical basis., The behaviour is certainly not random and must have a physical basis.
" Note that for simple trigonometric products. 6, will always be zero."," Note that for simple trigonometric products, $\phi_r$ will always be zero."
 Finally. we note that the amplitudes of the combination modes relative to the amplitudes of their. parents can be compared with values detected. in the star Tau. (?)..," Finally, we note that the amplitudes of the combination modes relative to the amplitudes of their parents can be compared with values detected in the star Tau \citep{BregerLenz2008}."
 They agree to a factor of two or better. suggesting that 90700322. is not unusual in this regard. just more accurately studied because of theAepler data.," They agree to a factor of two or better, suggesting that 9700322 is not unusual in this regard, just more accurately studied because of the data."
 In addition to. the quintuplet structure around the two dominant modes. another quintuplet with dillerent properties is present in WLC 9700822 (see the listing of fy to fs in refequin)), In addition to the quintuplet structure around the two dominant modes another quintuplet with different properties is present in KIC 9700322 (see the listing of $f_4$ to $f_8$ in \\ref{quin}) ).
 The average spacing between the frequencies in this quintuplet is slightly smaller than the rotational frequeney boys. D), The average spacing between the frequencies in this quintuplet is slightly smaller than the rotational frequency $^{-1}$ vs. $^{-1}$ ).
 This makes this quintuplet different from. the quintuplet structures found. around the two dominant modes. which exhibit a spacing that corresponds exactly to the rotation frequency.," This makes this quintuplet different from the quintuplet structures found around the two dominant modes, which exhibit a spacing that corresponds exactly to the rotation frequency."
" Moreover. the distribution of amplitudes within the third quintuplet is. fundamentally. dillerent το the patterns around f, and fo."," Moreover, the distribution of amplitudes within the third quintuplet is fundamentally different to the patterns around $f_1$ and $f_2$ ."
 The given characteristics support an interpretation of the quintuplet as an / = 2 moce., The given characteristics support an interpretation of the quintuplet as an $l$ = 2 mode.
 The location of the quintuplet near the centre in between the radial fundamental ancl first. overtone mode rules out pure acoustic character., The location of the quintuplet near the centre in between the radial fundamental and first overtone mode rules out pure acoustic character.
 Consequently. he observed quintuplet consists of mixed. modes. with considerable kinetic οποίον contribution from the eravity-mode cavity.," Consequently, the observed quintuplet consists of mixed modes with considerable kinetic energy contribution from the gravity-mode cavity."
 For such modes theory predicts. a smaller (and more svnimetrical) rotational splittings compared. to acoustic modes due to cdillerent values of the Ledoux constant Cj., For such modes theory predicts a smaller (and more symmetrical) rotational splittings compared to acoustic modes due to different values of the Ledoux constant $C_{nl}$.
 Using the framework of second. order theory (7) we determined. the equatorial rotation rate which oovides the best fit of the observed quintuplet with an (=2 multiplet., Using the framework of second order theory \citep{wdgd1992} we determined the equatorial rotation rate which provides the best fit of the observed quintuplet with an $\ell=2$ multiplet.
 The best results were obtained for an equatorial rotation rate of +., The best results were obtained for an equatorial rotation rate of $^{-1}$.
 This is only slighth higher han the observed esin/ value of |. and therefore indicates a near-equator-on-view.," This is only slightly higher than the observed $v \sin i$ value of $^{-1}$, and therefore indicates a near-equator-on-view."
 The Ledoux constant. Cy. of the (=2 quintuplet is 0.164.," The Ledoux constant, $C_{nl}$, of the $\ell=2$ quintuplet is 0.164."
" For. quadrupole modes C',j; ranges between 0.2 for pure gravity modes to smaller values for acoustic modes.", For quadrupole modes $C_{nl}$ ranges between $\approx$ 0.2 for pure gravity modes to smaller values for acoustic modes.
 With (1 - ο) = 0.836 this leads to a rotational frequency. £44= of around 0.16 +.," With (1 - $C_{nl}$ ) = 0.836 this leads to a rotational frequency, $\nu_{rot}=\frac{\Omega}{2\Pi}$, of around 0.16 $^{-1}$."
 Consequently. this theoretical result ADconfirms the interpretation of fs as a rotational feature. and. of. the quintuplet as / 2 2 modes.," Consequently, this theoretical result confirms the interpretation of $f_3$ as a rotational feature and of the quintuplet as $l$ = 2 modes."
 Further support is provided. by the fact that we see various combinations of the quintuplet with fi and. fo., Further support is provided by the fact that we see various combinations of the quintuplet with $f_1$ and $f_2$.
 Moreover. the location of the quintuplet allows us to determine the extent of overshooting from the convective core.," Moreover, the location of the quintuplet allows us to determine the extent of overshooting from the convective core."
 In the given model we obtained a.=0.13 but the uncertainties elaborated in Section 5.1. currenthy prevent an accurate determination., In the given model we obtained $\alpha_{ov}=0.13$ but the uncertainties elaborated in Section \ref{sec:radmodes} currently prevent an accurate determination.
 A remarkable aspect of the star is the fact that so. few pulsation modes are excited with amplitudes of pppm or larger., A remarkable aspect of the star is the fact that so few pulsation modes are excited with amplitudes of ppm or larger.
 1n the interior of an evolved 0 SSct. star. even high-frequency p modes behave like high-order g modes.," In the interior of an evolved $\delta$ Sct star, even high-frequency $p$ modes behave like high-order $g$ modes."
 The large number of spatial oscillations of these moces in the deep interior leads to severe radiative damping., The large number of spatial oscillations of these modes in the deep interior leads to severe radiative damping.
 As a result. nonracdial modes are increasingly damped for more massive OSSct stars. which explains why high-amplitucded SSct stars pulsate in mostly radial modes and why in even more massive classical Cepheics nonradial modes are no longer visible.," As a result, nonradial modes are increasingly damped for more massive $\delta$ Sct stars, which explains why $\delta$ Sct stars pulsate in mostly radial modes and why in even more massive classical Cepheids nonradial modes are no longer visible."
 In general. we do not expect the frequencies in the 9 SScet stars observed byΔορίο to be regularly spaced because. unlike grounc-basecl photometry. the observed pulsation modes are not limited to small spherical harmonic degree. ἐν," In general, we do not expect the frequencies in the $\delta$ Sct stars observed by to be regularly spaced because, unlike ground-based photometry, the observed pulsation modes are not limited to small spherical harmonic degree, $l$."
 For the very low amplitudes detected by.fvepler we may expect to see a large number of small-aumplitude modes with high /., For the very low amplitudes detected by we may expect to see a large number of small-amplitude modes with high $l$.
 The observed amplitudes decreases very slowly. with { and. all things being equal. a large number of modes with high / might be expected to be seen in 0 SSct and other stars (?)..," The observed amplitudes decreases very slowly with $l$ and, all things being equal, a large number of modes with high $l$ might be expected to be seen in $\delta$ Sct and other stars \citep{Balona1999}."
 The 8 SSct stars 550844 (7). and 1174936 (7) observed byCotoT show many hundreds of eloselv-spaced frequencies and may be examples of high-degree modes., The $\delta$ Sct stars 50844 \citep{Poretti2009} and 174936 \citep{Hernandez2009} observed by show many hundreds of closely-spaced frequencies and may be examples of high-degree modes.
 The relatively small number of independent frequencies detected in 90700322 stands in strong contrast to the two stars observed by., The relatively small number of independent frequencies detected in 9700322 stands in strong contrast to the two stars observed by.
Cofto. lt should be noted that. unlike many dSSet stars observed. byAepéer. 99700322. does. not. have any frequencies in the range INKICnormally seen in 5 DDor stars.," It should be noted that, unlike many $\delta$ Sct stars observed by, 9700322 does not have any frequencies in the range normally seen in $\gamma$ Dor stars."
 The only strong [frequencies in this range are a few combination frequencies., The only strong frequencies in this range are a few combination frequencies.
 Although we have identified significant frequencies below +. it is not. possible at this stage to verily whether these are cue to the star or instrumental artefacts.," Although we have identified significant frequencies below $^{-1}$, it is not possible at this stage to verify whether these are due to the star or instrumental artefacts."
 At present. we do not understand why low frequencies arc present in so many ὁ SScet stars.," At present, we do not understand why low frequencies are present in so many $\delta$ Sct stars."
 ltegularities in the frequency spacing clue to combination modes have already been observed from. the erounc even in low amplitude ὁ SSct stars., Regularities in the frequency spacing due to combination modes have already been observed from the ground even in low amplitude $\delta$ Sct stars.
 An example is the star T'Fau (2).., An example is the star Tau \citep{BregerLenz2008}.
 22 of ?. demonstrates that all the observed regularities outside the 513 + ranee are caused by combination moces., 2 of \cite{BregerLenzPamyatnykh2009} demonstrates that all the observed regularities outside the $5-13$ $^{-1}$ range are caused by combination modes.
 For combination nioces the frequency spacing must be absolutely regular within the limits of measurability., For combination modes the frequency spacing must be absolutely regular within the limits of measurability.
 This is found for 99700322., This is found for 9700322.
 ALD is grateful to LE. L. Robinson and M. Montgomery. [or helpful discussions., MB is grateful to E. L. Robinson and M. Montgomery for helpful discussions.
 This investigation has been supported by the Austrian Fonds zur Fórrderung der wissenschaltlichen Forschung through project P? 21830-N16., This investigation has been supported by the Austrian Fonds zur Förrderung der wissenschaftlichen Forschung through project P 21830-N16.
 LAB which to acknowledge financial support from the South African Astronomical Observatory., LAB which to acknowledge financial support from the South African Astronomical Observatory.
 AAP and. PL acknowledge partial financial support fromthe Polish MNiSW grant. No., AAP and PL acknowledge partial financial support fromthe Polish MNiSW grant No.
 Ν N203 379 636., N N203 379 636.
" This work has been supported by the ""Lendüllet program of the Llungarian Academy of Sciences and Hungarian OTIX grant. Ix83790.", This work has been supported by the `Lendüllet' program of the Hungarian Academy of Sciences and Hungarian OTKA grant K83790.
 The authors wish to thank theAepíer team for their generosity in allowing the data to be released to the, The authors wish to thank the team for their generosity in allowing the data to be released to the
 The authors wish to thank theAepíer team for their generosity in allowing the data to be released to the , The authors wish to thank the team for their generosity in allowing the data to be released to the
 The authors wish to thank theAepíer team for their generosity in allowing the data to be released to the W, The authors wish to thank the team for their generosity in allowing the data to be released to the
 The authors wish to thank theAepíer team for their generosity in allowing the data to be released to the We, The authors wish to thank the team for their generosity in allowing the data to be released to the
 The authors wish to thank theAepíer team for their generosity in allowing the data to be released to the Wep, The authors wish to thank the team for their generosity in allowing the data to be released to the
 The authors wish to thank theAepíer team for their generosity in allowing the data to be released to the Wepl, The authors wish to thank the team for their generosity in allowing the data to be released to the
 The authors wish to thank theAepíer team for their generosity in allowing the data to be released to the Weple, The authors wish to thank the team for their generosity in allowing the data to be released to the
 The authors wish to thank theAepíer team for their generosity in allowing the data to be released to the Wepler, The authors wish to thank the team for their generosity in allowing the data to be released to the
"In this paper, we have investigated the role that the choice of the initial seed black hole mass function at high redshift (z~ 18) plays in the determination of observed properties of local quiescent SMBHs.","In this paper, we have investigated the role that the choice of the initial seed black hole mass function at high redshift $z \sim 18$ ) plays in the determination of observed properties of local quiescent SMBHs."
" While the errors on mass determinations of local black holes are large at the present time, definite trends with host galaxy properties are observed."," While the errors on mass determinations of local black holes are large at the present time, definite trends with host galaxy properties are observed."
 The tighest correlation appears to be between the BH mass and the velocity dispersion of the host spheroid., The tighest correlation appears to be between the BH mass and the velocity dispersion of the host spheroid.
" Starting with the ab-initio black hole seed mass function computed in the context of direct formation of central objects from the collapse of pre-galactic discs in high redshift halos, we follow the assembly history to late times using a Monte-Carlo merger tree approach."," Starting with the ab-initio black hole seed mass function computed in the context of direct formation of central objects from the collapse of pre-galactic discs in high redshift halos, we follow the assembly history to late times using a Monte-Carlo merger tree approach."
 Key to our calculation of the evolution and build-up of mass is the prescription that we adopt for determining the precise mass gain during a merger., Key to our calculation of the evolution and build-up of mass is the prescription that we adopt for determining the precise mass gain during a merger.
" Motivated by the phenomenological observation of MauοςVe, we assume that this proportionality carries over to the gas mass accreted in each step."," Motivated by the phenomenological observation of $M_{\rm BH} \propto V_{\rm c}^5$, we assume that this proportionality carries over to the gas mass accreted in each step."
" With these prescriptions, a range of predictions can be made for the mass function of blackholes at high and low z, and the integrated mass density of black holes, all of which are observationally determined."," With these prescriptions, a range of predictions can be made for the mass function of blackholes at high and low $z$, and the integrated mass density of black holes, all of which are observationally determined."
" We evolve 3 models, designated model A, B and C which correspond to increasing efficiencies respectively for the formation of seeds at high redshift."," We evolve 3 models, designated model A, B and C which correspond to increasing efficiencies respectively for the formation of seeds at high redshift."
 T'hese models are compared to one in which the seeds are remnants of Population III stars., These models are compared to one in which the seeds are remnants of Population III stars.
 It is important to note here that one major uncertainty prevents us from making more concrete predictions: the unknown metal enrichment history of the Universe., It is important to note here that one major uncertainty prevents us from making more concrete predictions: the unknown metal enrichment history of the Universe.
 Key to the implementation of our models is the choice of redshift at which massive seed formation is quenched., Key to the implementation of our models is the choice of redshift at which massive seed formation is quenched.
 The direct seed formation channel described here ceases to operate once the Universe has been enriched by metals that have been synthesized after the first generation of stars have gone supernova., The direct seed formation channel described here ceases to operate once the Universe has been enriched by metals that have been synthesized after the first generation of stars have gone supernova.
 Once metals are available in the Inter-Galactic Medium gas cooling is much more efficient and hydrogen in either atomic or molecular form is no longer the key player., Once metals are available in the Inter-Galactic Medium gas cooling is much more efficient and hydrogen in either atomic or molecular form is no longer the key player.
 In this work we have assumed this transition redshift to be z—15., In this work we have assumed this transition redshift to be $z = 15$.
" The efficiency of MBH formation and the transition redshift are somehow degenerate (e.g., a model with Q—1.5 and enrichment redshift z—12 is halfway between model A and model B); if other constraints on this redshift were available we could considerably tighten our predictions."," The efficiency of MBH formation and the transition redshift are somehow degenerate (e.g., a model with $Q=1.5$ and enrichment redshift $z=12$ is halfway between model A and model B); if other constraints on this redshift were available we could considerably tighten our predictions."
 Below we list our predictions and compare how they fare with respect to current observations., Below we list our predictions and compare how they fare with respect to current observations.
 The models investigated here clearly differ in predictions at the low mass end of the black hole mass function., The models investigated here clearly differ in predictions at the low mass end of the black hole mass function.
" With future observational sensitivity in this domain, these models can be distinguished."," With future observational sensitivity in this domain, these models can be distinguished."
 Our model for the formation of relatively high-mass black hole seeds in high-z halos has direct influence on the black hole occupation fraction in galaxies at z=0., Our model for the formation of relatively high-mass black hole seeds in $z$ halos has direct influence on the black hole occupation fraction in galaxies at $z=0$.
 This effect is more pronounced for low mass galaxies., This effect is more pronounced for low mass galaxies.
 We find that a significant fraction of low-mass galaxies might not host a nuclear black hole., We find that a significant fraction of low-mass galaxies might not host a nuclear black hole.
" This is in very good agreement with the shape of the My,—σ relation determined recently from an observational census (an HST ACS survey) of low mass galaxies in the Virgo cluster reported by Ferrarese et al. (", This is in very good agreement with the shape of the $M_{\rm bh} - \sigma$ relation determined recently from an observational census (an HST ACS survey) of low mass galaxies in the Virgo cluster reported by Ferrarese et al. (
2006).,2006).
" The models studied here (with different black hole seed formation efficiency) are distinguishable at the low mass end of the BH mass function, while at the high mass end the effect of initial seeds appears to be sub-dominant."," The models studied here (with different black hole seed formation efficiency) are distinguishable at the low mass end of the BH mass function, while at the high mass end the effect of initial seeds appears to be sub-dominant."
" While current data in the low mass regime is scant (??),, future instruments and surveys are likely to probe this region of parameter space with significantly higher sensitivity."," While current data in the low mass regime is scant \citep{Barth2004,Greene2007}, future instruments and surveys are likely to probe this region of parameter space with significantly higher sensitivity."
" All our models predict that low surface brightness, bulge-less galaxies with high spin parameters (i.e. large discs) are systems where MBH formation is least probable."," All our models predict that low surface brightness, bulge-less galaxies with high spin parameters (i.e. large discs) are systems where MBH formation is least probable."
 One of the key caveats of our picture is that it is unclear if the differences produced by different seed models on observables at z=0 might be compensated or masked by BH fueling modes at earlier epochs., One of the key caveats of our picture is that it is unclear if the differences produced by different seed models on observables at $z = 0$ might be compensated or masked by BH fueling modes at earlier epochs.
" There could be other channels for BH growth that dominate at low redshifts like minor mergers, dynamical instabilities, accretion of molecular clouds, tidal disruption of stars."," There could be other channels for BH growth that dominate at low redshifts like minor mergers, dynamical instabilities, accretion of molecular clouds, tidal disruption of stars."
" The decreased importance of the merger driven scenario is patent from observations of low-redshift AGN, which are for the large majority hosted by undisturbed galaxies (e.g.,?,andreferencestherein) in low-density environments (e.g.,?).."," The decreased importance of the merger driven scenario is patent from observations of low-redshift AGN, which are for the large majority hosted by undisturbed galaxies \citep[e.g.,][and references therein] {Pierce2007} in low-density environments \citep[e.g.,][]{Li2006}."
" However, the feasibility and efficiency of some alternative channels are still to be proven (for example, about the efficiency of feeding from large scale instabilities see discussion in ????))."," However, the feasibility and efficiency of some alternative channels are still to be proven (for example, about the efficiency of feeding from large scale instabilities see discussion in \citealt{KingPringle2007,shlosman89,goodman03,collin99}) )."
" In any event, while these additional channels for BH growth can modify the detailed shape of the mass function of MBHs, or of the luminosity"," In any event, while these additional channels for BH can modify the detailed shape of the mass function of MBHs, or of the luminosity"
Nam The high euerev πιο events collected by the ALACRO apparatus at the average depth of 3600 1iw.e.,.7 cm The high energy muon events collected by the MACRO apparatus at the average depth of 3600 m.w.e.
 represent one of the most exteusive records of such kind of data., represent one of the most extensive records of such kind of data.
 The series of these high-cnerey unmons can be to search for time variations of periodic aud of stochastic characters. as it was doue extensively by using arrival times of EAS |1]..," The series of these high-energy muons can be to search for time variations of periodic and of stochastic characters, as it was done extensively by using arrival times of EAS \cite{first}."
 These variations in the underground amon fux may be due to differcut causes of ealactic. solar and terrestrial origin.," These variations in the underground muon flux may be due to different causes of galactic, solar and terrestrial origin."
 The conunon problem for this type of searches is to determine whether an observed effect has occurred by chance of if it signals a departure frou a pure raucous nmon arrival., The common problem for this type of searches is to determine whether an observed effect has occurred by chance of if it signals a departure from a pure random muon arrival.
 MACRO was a multipurpose modular apparatus with 6 supermodiules with scintillator detectors. limited streamer tubes aud unclear track detectors |2].. and studied atinospherie neutrinos [3].. aspects of CR plivsics and astroplivsics |1].. searched. for GUT Magnetic Monopoles and other exotica [5|]..," MACRO was a multipurpose modular apparatus with 6 supermodules with scintillator detectors, limited streamer tubes and nuclear track detectors \cite{MACRO}, and studied atmospheric neutrinos \cite{neu}, aspects of CR physics and astrophysics \cite{crmu}, searched for GUT Magnetic Monopoles and other exotica \cite{MM}."
 Some interruptions of different kinds occurred during data taking. either randomly (6.8. power outages}. or regularly (6.8. maintenance). so appropriate statistical methods have to be applied aud particular care should be used in choosing periods of stationary conuditious.," Some interruptions of different kinds occurred during data taking, either randomly (e.g. power outages), or regularly (e.g. maintenance), so appropriate statistical methods have to be applied and particular care should be used in choosing periods of stationary conditions."
 Iu the following we discuss the results of the searches for periodic variations aud for time clustering of nmon events., In the following we discuss the results of the searches for periodic variations and for time clustering of muon events.
" For this analysis we considered data recorded by the streamer tube svstei in the time muterval November 1991-Maxy. 2000 aud selected the data with the following criteria: - un duration lounger than 1 hour: - streamer tube efficieucies of wires aud strips larger than andτον, respectively. for each module: - all 6 super-iuodules in acquisition: - acquisition dead time snaller than for the whole detector."," For this analysis we considered data recorded by the streamer tube system in the time interval November 1991-May 2000 and selected the data with the following criteria: - run duration longer than 1 hour; - streamer tube efficiencies of wires and strips larger than and, respectively, for each module; - all 6 super-modules in acquisition; - acquisition dead time smaller than for the whole detector."
 The total uuuber of rus surviving these cuts was 6920 for a total nuuber of 3.5-10° mon events., The total number of runs surviving these cuts was 6920 for a total number of $3.5 \cdot 10^7$ muon events.
"(which is in the northern hemisphere), while is directed upward elsewhere (all the other sites are located in the southern hemisphere).","(which is in the northern hemisphere), while is directed upward elsewhere (all the other sites are located in the southern hemisphere)."
 We can notice also a sign reversal at the frequencies v~99 GHz and v~160 GHz., We can notice also a sign reversal at the frequencies $\nu \simeq 99$ GHz and $\nu \simeq 160$ GHz.
 The first null is originated by the 118 GHz line wing crossing the 50-70 GHz lines wings., The first null is originated by the 118 GHz line wing crossing the 50-70 GHz lines wings.
 The null at higher frequency is originated adding the contribution from lines with L>9., The null at higher frequency is originated adding the contribution from lines with $L \geq 9$.
" Adding lines with higher L the null frequency shifts to lower values, converging to the final value shown in Fig. 2.."," Adding lines with higher $L$ the null frequency shifts to lower values, converging to the final value shown in Fig. \ref{pol_comparison}. ."
" The coefficients Ao, Αι and Ag (see eq."," The coefficients $A_0$, $A_1$ and $A_2$ (see eq."
" 6 and 7), which assume different values line by line, seem to play an important role in this effect."," 6 and 7), which assume different values line by line, seem to play an important role in this effect."
" Anyway, as discussed in section ??,, it should be stressed that above ~120 GHz the model is not yet validated, and calculation could be not accurate."," Anyway, as discussed in section \ref{acc_temp}, it should be stressed that above $\sim$ 120 GHz the model is not yet validated, and calculation could be not accurate."
 We then built the maps of the polarized signal as seen at the various sites in local alt-azimuthal coordinates., We then built the maps of the polarized signal as seen at the various sites in local alt-azimuthal coordinates.
" Maps of the polarized signal, at 90 GHz, seen at the selected sites are shown in Figures 3,, 4,, 5 and 6.."," Maps of the polarized signal, at 90 GHz, seen at the selected sites are shown in Figures \ref{dc}, \ref{atacama}, \ref{sp} and \ref{tg}."
 Both positive and negative values indicate that both right and left handed circular polarization is present., Both positive and negative values indicate that both right and left handed circular polarization is present.
 We find V—0 in the maps (red contour) at the angular positions where is orthogonal to the line of sight., We find $V=0$ in the maps (red contour) at the angular positions where is orthogonal to the line of sight.
 Elevation cuts along the North-South direction and azimuth cuts at constant elevation of 45? of these maps are presented respectively in Figures 7 and 8.., Elevation cuts along the North-South direction and azimuth cuts at constant elevation of $^{\circ}$ of these maps are presented respectively in Figures \ref{cut_elev_NS} and \ref{cut_az_45}.
 Observing in directions different from the zenith the actual signal depends also on the atmospheric thickness., Observing in directions different from the zenith the actual signal depends also on the atmospheric thickness.
 This effect is visible in the elevation scans that show a zenith-secant law dependence., This effect is visible in the elevation scans that show a zenith-secant law dependence.
 The spatial structure of these maps is dominated by the combination of magnetic field direction and atmospheric thickness in the line of sight., The spatial structure of these maps is dominated by the combination of magnetic field direction and atmospheric thickness in the line of sight.
" Each map covers a surface, which has a typical scale length of ~200 km, at an altitude of ~30 km, corresponding to a multipole n~200."," Each map covers a surface, which has a typical scale length of $\sim 200$ km, at an altitude of $\sim 30$ km, corresponding to a multipole $n \sim 200$."
 The magnetic field components at these multipoles (or even higher) are negligible with respect to the longer wavelength components and not larger than the accuracy of the IGRF-2010 model (~ 10 nT)., The magnetic field components at these multipoles (or even higher) are negligible with respect to the longer wavelength components and not larger than the accuracy of the IGRF-2010 model $\sim$ 10 nT).
 The effect due to the variation of the atmospheric thickness is combined in the maps with the direction of the magnetic field., The effect due to the variation of the atmospheric thickness is combined in the maps with the direction of the magnetic field.
" At Dome C, where the magnetic field is almost vertical, the two effects compensate each other and the signal is flat on a large part of the visible sky."," At Dome C, where the magnetic field is almost vertical, the two effects compensate each other and the signal is flat on a large part of the visible sky."
" Conversely at Atacama, where the magnetic field is almost horizontal, the two effects combine increasing the gradient along the visible sky."," Conversely at Atacama, where the magnetic field is almost horizontal, the two effects combine increasing the gradient along the visible sky."
" The typical signal variation, at90 GHz, for North-South elevation scans (down to 60? from the zenith) is of the order"," The typical signal variation, at90 GHz, for North-South elevation scans (down to $^{\circ}$ from the zenith) is of the order"
warn that they are likely metallicity dependent (Westbelow )..,"warn that they are likely metallicity dependent \citep[see also Section 3.3
below]{west04, bootem}. ."
 Figure 23. shows the r versus r> ess diagram for the AL dwarfs in the SDSS DR sample., Figure \ref{fig:hess} shows the $r$ versus $r-z$ Hess diagram for the M dwarfs in the SDSS DR7 sample.
 The gaps in color space reflect the non-unifoiii smupliug due to the spectral targetiug algorithm., The gaps in color space reflect the non-uniform sampling due to the spectral targeting algorithm.
 To quantifv the effect of SDSS spectroscopic targeting aleorithlius on the median colors. we conducted the following test.," To quantify the effect of SDSS spectroscopic targeting algorithms on the median colors, we conducted the following test."
 We begau with the straw man assumption that the ΕΕ”. A loud rt color-spectral type relations reported in Table 20 are represcutative of the underlying stellar »pulatiou.," We began with the straw man assumption that the median $r-i$, $i-z$ and $r-z$ color-spectral type relations reported in Table \ref{tab:color} are representative of the underlying stellar population."
 We then used these quantities. along with he observed color spread in cach spectral type bin. to construct a svuthetic stellar population with a unifoniii distribution iu both maeuituce and spectral type.," We then used these quantities, along with the observed color spread in each spectral type bin, to construct a synthetic stellar population with a uniform distribution in both magnitude and spectral type."
 We hen applied photometric cuts to the smooth underlying distribution. replicating the SDSS spectroscopic selection and producing a svuthetic “observed” sample that is Πο] structured in color-magnitude space.," We then applied photometric cuts to the smooth underlying distribution, replicating the SDSS spectroscopic selection and producing a synthetic “observed” sample that is highly structured in color-magnitude space."
 The median colors aud spreads were then calculated as a function of spectral type from this svuthetic sample auc compared o the measured colors iu Table 2.., The median colors and spreads were then calculated as a function of spectral type from this synthetic sample and compared to the measured colors in Table \ref{tab:color}.
 Iu general. the colors of the svuthetic stars agree to within 0.02 mags of the observed. suuple (an agreement that would not ο expected if the uuderlviug population were different roni the observed SDSS stars ) aud we conclude that the colors preseuted iu Table 20 are not severely affected by he spectroscopic selection.," In general, the colors of the synthetic stars agree to within 0.02 mags of the observed sample (an agreement that would not be expected if the underlying population were different from the observed SDSS stars ) and we conclude that the colors presented in Table \ref{tab:color} are not severely affected by the spectroscopic selection."
 Figure E. shows the Wa activity fraction as a function of AL dwarf spectral type for stars with data quality sufficient to measure Tea cussion (uid GOODPIIOT =1 and WDAL 20)., Figure \ref{hafrac} shows the $\alpha$ activity fraction as a function of M dwarf spectral type for stars with data quality sufficient to measure $\alpha$ emission (and GOODPHOT =1 and WDM =0).
 Our results are in good agreement with previous studies (Westetal.2004.WoOs).. but with much lower uncertainties due to the 259.000 NI cavarts used to generate Figure |. (7952 are Πα active and 51.031 are Τα inactive).," Our results are in good agreement with previous studies \citep[W08]{west04}, but with much lower uncertainties due to the $\sim$ 59,000 M dwarfs used to generate Figure \ref{hafrac} (7952 are $\alpha$ active and 51,034 are $\alpha$ inactive)."
 There are a few πα] differeuces between Figure { and the Πα activity fractions reported by WWo0s that are likely due to the changes in some of the spectral types (see above) and the different Calactic sieltlues included in the DR? sample., There are a few small differences between Figure \ref{hafrac} and the $\alpha$ activity fractions reported by \nocite{west08}W W08 that are likely due to the changes in some of the spectral types (see above) and the different Galactic sightlines included in the DR7 sample.
 Because of the lower SNR in the blue portion of AL dwarfs. fewer stars were available to measure activity for the higher order Balmer aud Call Is. ciuuission lines.," Because of the lower SNR in the blue portion of M dwarfs, fewer stars were available to measure activity for the higher order Balmer and CaII K emission lines."
 Table 3 gives the munber of active aud inactive stars for cach cussion line indicator that had the necessary quality for our activity analysis., Table \ref{activity} gives the number of active and inactive stars for each emission line indicator that had the necessary quality for our activity analysis.
 As discussed in WWOs. the activity faction is highly dependent on the location of the samples iu the Galaxy.," As discussed in \nocite{west08}W W08, the activity fraction is highly dependent on the location of the samples in the Galaxy."
 Because AL dwarts have finite activity lifetimes aud are dynamically heated in the Galactic disk as they age. their activity state is correlated with position in the Galaxy.," Because M dwarfs have finite activity lifetimes and are dynamically heated in the Galactic disk as they age, their activity state is correlated with position in the Galaxy."
 One of the reasons that the M. dwarfs in our sample have much simaller activity fractions than those studied nearby (ee.Dawleyetal.1996:Gizis2000).. is that the SDSS volume is much larger than those used In previous catalogs aud is concentrated on the north ealactic cap. vieldiug a auch older stella population.," One of the reasons that the M dwarfs in our sample have much smaller activity fractions than those studied nearby \citep[e.g.][]{hawley96,gizis00}, is that the SDSS volume is much larger than those used in previous catalogs and is concentrated on the north galactic cap, yielding a much older stellar population."
 This is particularly clear for the carly-tvpe AI dwarts. which have median distances greater than 500 pc. aud therefore ages that are considerably older than their short activity lifetimes (~ 1-2 Cir)," This is particularly clear for the early-type M dwarfs, which have median distances greater than 500 pc, and therefore ages that are considerably older than their short activity lifetimes $\sim$ 1-2 Gyr)."
 We reiterate the warning frou W WS that activity fractious in AL charts must be discussed in the proper Galactic context., We reiterate the warning from \nocite{west08}W W08 that activity fractions in M dwarfs must be discussed in the proper Galactic context.
 The various cussion lines measured iu our spectra are formed at slightly different locations im the chromosphere. suggesting— that the Wf.streneth of one Cluission line may not necessarily predict the strenetl of another.," The various emission lines measured in our spectra are formed at slightly different locations in the chromosphere, suggesting that the strength of one emission line may not necessarily predict the strength of another."
 We therefore examined how the various activity induced ciuission lines trace each other as a function of spectral type and absolute distance from the Galactic plane (a proxy for age)., We therefore examined how the various activity induced emission lines trace each other as a function of spectral type and absolute distance from the Galactic plane (a proxy for age).
 Figure Ὁ shows the Te (diamouds) aud I> (asterisks) activity fractious for ΑΟΧΙ dwarfs as afunction of Galactic height., Figure \ref{betafrac} shows the $\alpha$ (diamonds) and $\beta$ (asterisks) activity fractions for M2-M7 dwarfs as afunction of Galactic height.
 Oulv the IIo activity fractions for stars that could have detected, Only the $\alpha$ activity fractions for stars that could have detected
containing I2|A galaxies appear dissimilar to those hosting elliptical galaxies: the elliptical. model is rejected. by the average-Dalmer (or H8) data with a confidence of 99.7%ss (98.0%). based on the value of the chi-squared statistic.,"containing E+A galaxies appear dissimilar to those hosting elliptical galaxies; the elliptical model is rejected by the average-Balmer (or $\delta$ ) data with a confidence of $99.7\%$ $98.0\%$ ), based on the value of the chi-squared statistic."
 As à final method. of investigating the environment of PdPGRS LE|A galaxies. we obtained the Supercosmos Sky Survey (SSS) photometric catalogues for the regions surrounding cach 15|A galaxy.," As a final method of investigating the environment of 2dFGRS E+A galaxies, we obtained the Supercosmos Sky Survey (SSS) photometric catalogues for the regions surrounding each E+A galaxy."
" The approximate magnitude limits of the photographic plates which are scanned to produce the SSS are b,=22.5 and rp=21.5.", The approximate magnitude limits of the photographic plates which are scanned to produce the SSS are $\bj = 22.5$ and $\rf = 21.5$.
 We downloaded: SSS object catalogues using the web interface specifying a circular extraction with radius 5 arcmin., We downloaded SSS object catalogues using the web interface specifying a circular extraction with radius 5 arcmin.
 Using the 2dkCRS }y-band luminosity function (Norberg et 22002). we can deduce the apparent magnitude 6)(2) corresponding to an absolute magnitude Adj at the red:shift z of a sample LE}A galaxy.," Using the 2dFGRS $\bj$ -band luminosity function (Norberg et 2002), we can deduce the apparent magnitude $\bj^*(z)$ corresponding to an absolute magnitude $M_b^*$ at the redshift $z$ of a sample E+A galaxy."
 We then define a ‘bright neighbour as a nearby SSS galaxy. with by<by(2)|1., We then define a `bright neighbour' as a nearby SSS galaxy with $\bj < \bj^*(z) + 1$.
 We define a ‘faint neighbour’ as a nearby S ealaxy with b3(2)|1<6)«22.5., We define a `faint neighbour' as a nearby SSS galaxy with $\bj^*(z) + 1 < \bj < 22.5$.
 The presence of a nearby γιοί or ‘faint’ neighbour may indicate. respectively. an imminent major or nilnor merger.," The presence of a nearby `bright' or `faint' neighbour may indicate, respectively, an imminent major or minor merger."
 Using the known I2|.X galaxy. redshift we can convert angular separations into transverse physical separations., Using the known E+A galaxy redshift we can convert angular separations into transverse physical separations.
 For each LE|A galaxy we derived: Note that cach of these quantities is insensitive to the ellects of peculiar velocities., For each E+A galaxy we derived: Note that each of these quantities is insensitive to the effects of peculiar velocities.
 We compared these statistics to those determined for a sample of 1000 randomdls-drawn 2dECHIUS ealaxies (restricted. to redshifts 0.002<2«0.25 to eliminate stellar contamination)., We compared these statistics to those determined for a sample of 1000 randomly-drawn 2dFGRS galaxies (restricted to redshifts $0.002 < z < 0.25$ to eliminate stellar contamination).
 The results for the two 151A galaxy. catalogues are cisplaved in reffielocab ane 10.. with the results for the random sample overplotted in cach case.," The results for the two E+A galaxy catalogues are displayed in \\ref{figlocab} and \ref{figlochd}, with the results for the random sample overplotted in each case."
 We used the Ixolmogorov-Smürnov. (Ίντο). statistical test to ascertain whether the local environments. of the EVA galaxy samples were drawn from a cilferent parent distribution. to the local environments of the random 3011 sample., We used the Kolmogorov-Smirnov (K-S) statistical test to ascertain whether the local environments of the E+A galaxy samples were drawn from a different parent distribution to the local environments of the random 2dFGRS sample.
 Let p be the probability that the K-S statistic could exceed the observed value by random chance. if the two distributions were drawn from the same parent distribution.," Let $p$ be the probability that the K-S statistic could exceed the observed value by random chance, if the two distributions were drawn from the same parent distribution."
 The values ofp for the distributions of nearest faint neighbour. nearest bright neighbour and local surface density. for the average-Dalmoer (118) LE|A catalogue. were 0.58 (0.35). 0.32 (0.11) and 0.21 (0.06).," The values of $p$ for the distributions of nearest faint neighbour, nearest bright neighbour and local surface density, for the average-Balmer $\delta$ ) E+A catalogue, were $0.58$ $0.35$ ), $0.32$ $0.11$ ) and $0.21$ $0.06$ )."
instantancously towards an observer indicated at the top of the page. making an angle with OS of where ο is the speed of light and rj the light evlinder radius.,"instantaneously towards an observer indicated at the top of the page, making an angle with OS of where c is the speed of light and $r_{lc}$ the light cylinder radius."
" AX κος 2M at which all incident radiation is assumed to be absorbed. is also rotating at a rate © and located at a distance r7, and making a permanent angle of O with OS.", A site $A^{*}$ at which all incident radiation is assumed to be absorbed is also rotating at a rate $\Omega$ and located at a distance $r_{*}$ and making a permanent angle of $\Theta$ with OS.
 If at a time {550 15 is positioned so that its radiation (subject to aberration) is directed towards the observer. then the radius vector of the absorber A will at that moment make an angle 0 with the direction to the observer. given by Since (6.©) is the angle the sources postion vector makes with the observer's line of sight. (2) defines the equation of the critical spiral.," If at a time t=0 S is positioned so that its radiation (subject to aberration) is directed towards the observer, then the radius vector of the absorber $A^{*}$ will at that moment make an angle $\theta$ with the direction to the observer, given by Since $(\theta-\Theta)$ is the angle the source's postion vector makes with the observer's line of sight, (2) defines the equation of the critical spiral."
 We need to know for which coordinates (r.O) of S will the radiation from S be absorbed. before it reaches the observer (io we seek to know from which position or positions on the spiral will the radiation be absorbed hy A*)., We need to know for which coordinates $\Theta$ ) of S will the radiation from S be absorbed before it reaches the observer (ie we seek to know from which position or positions on the spiral will the radiation be absorbed by A*).
 Hepresented in Fig 2. this requires that 21 has rotated to a new position A dn the same time that the radiation has travelled. [rom 8 to 21.," Represented in Fig 2, this requires that $A^{*}$ has rotated to a new position $A^{'}$ in the same time that the radiation has travelled from S to $A^{'}$."
" Lf e; the angle. between the new vector OLY and the t=0 vector for 8S. is assumed to be small (as it will be if the aberration elfect is small). then the distance travelled by the raciation before absorption will be (r, or) and o is ENὃνgiven in. the small angleC approximation. by The crucial condition that A* rotates to 4 in the same time. T. that the raciation travels from S to V is then which vieles from (4) and. substituting from (2). we obtain a quadratic for r: =O The two solutions ror. of equation(7) represent two possible locations at which an emitter S can be obscured by 2U as it rotates."," If $\phi{`}$, the angle between the new vector $OA{'}$ and the t=0 vector for S, is assumed to be small (as it will be if the aberration effect is small), then the distance travelled by the radiation before absorption will be $r_{*}-r$ ) and $\phi{'}$ is given in the small angle approximation by hence from (1) The crucial condition that A* rotates to $A{'}$ in the same time, T, that the radiation travels from S to $A{'}$ is then which yields from (4) and, substituting from (2), we obtain a quadratic for r: )=0 The two solutions $r_{1},r_{2}$ of equation(7) represent two possible locations at which an emitter S can be obscured by $A^{*}$ as it rotates."
" These solutions satisfy and for both these roots to be positive we additionally require that rere>O. or Given that 8,,,, must be less than unity. we have our inst constraint on the position of the absorber. namely that if raclation is to be absorbed then it must be emitted when he absorbers radial vector bears an angle less than one radian (about 57°) to the observer's line of sight. and in oactice much less."," These solutions satisfy and for both these roots to be positive we additionally require that $r_{1}r_{2}>0$, or Given that $\theta_{max}$ must be less than unity, we have our first constraint on the position of the absorber, namely that if radiation is to be absorbed then it must be emitted when the absorber's radial vector bears an angle less than one radian (about $57^{o}$ ) to the observer's line of sight, and in practice much less."
 The dilference (iethe approximate distance) between he two locations is so that and the physical condition that ¢ is real. together with (9) above. establish the upper and lower limits on 8 as where ie only when ÀA* is between these two angular positions can radiation be absorbed.," The difference (i.e.the approximate distance) between the two locations is so that and the physical condition that $\epsilon$ is real, together with (9) above, establish the upper and lower limits on $\theta$ as where ie only when A* is between these two angular positions can radiation be absorbed."
 This defines the duration. of Phase 1. at the conclusion of which «0 and the two possible solutions for the position of S coalesce to the position I in Fig 1.," This defines the duration of Phase 1, at the conclusion of which $\epsilon\rightarrow{0}$ and the two possible solutions for the position of S coalesce to the position E in Fig 1."
 The upper limit is only appropriate when rj7rs and rS$0., The upper limit is only appropriate when $r_{1}\rightarrow{r_{*}}$ and $r_{2}\rightarrow{0}$.
" Phe more interesting case here is when r, and ro are close to each other and A is located at 6,;,.", The more interesting case here is when $r_{1}$ and $r_{2}$ are close to each other and $A^{*}$ is located at $\theta_{min}$.
" Phrough Phase 1. as 6 decreases from 06,,,,. the positions of r, and rs eraclually move along the spiral towards each other. meeting at the lower limit of (13) and fixing the location of E in Fig lon the critical spiral."," Through Phase 1, as $\theta$ decreases from $\theta_{max}$, the positions of $r_{1}$ and $r_{2}$ gradually move along the spiral towards each other, meeting at the lower limit of (13) and fixing the location of E in Fig 1 on the critical spiral."
 Hence in general a narrow extended absorbing region centred. racially or azimuthallv on A* will al any moment obscure two strips of the dotted. spiral on either side of the midpoint. as indicated in Fig 1.," Hence in general a narrow extended absorbing region centred radially or azimuthally on A* will at any moment obscure two strips of the dotted spiral on either side of the midpoint, as indicated in Fig 1."
as the Spearman's correlation cocllicient for the relations. support this conclusion.,"as the Spearman's correlation coefficient for the relations, support this conclusion."
 First. note that the regions with the highest (PALL S sana)/160 jum ratios within some of these ealaxies are found outside the nucleus.," First, note that the regions with the highest (PAH 8 $\mu$ m)/160 $\mu$ m ratios within some of these galaxies are found outside the nucleus."
 As can be seen in ligure 1.. the regions with enhanced. (PALL S sana)/160 jum ratios may correspond to spiral structure. as is most clearly seen in NGC 3081) and. NGC 6046.," As can be seen in Figure \ref{f_map}, the regions with enhanced (PAH 8 $\mu$ m)/160 $\mu$ m ratios may correspond to spiral structure, as is most clearly seen in NGC 3031 and NGC 6946."
 In. NGC 41725. the inner ring has the highest (PALL S sam)/160 pam ratio. not the nucleus.," In NGC 4725, the inner ring has the highest (PAH 8 $\mu$ m)/160 $\mu$ m ratio, not the nucleus."
 The intrinsic scatter in the fits versus surface brightness is generally lower than for those for the fits versus racius., The intrinsic scatter in the fits versus surface brightness is generally lower than for those for the fits versus radius.
 Moreover. the Spearman's correlation coelficients for the relations between the (PAIL S fam)/160 jim ratio and 160 pm surface brightness generally has a higher absolute value than the corresponding correlation coellicients for the relation between the (PATE S sam)/160 jim ratio and radius.," Moreover, the Spearman's correlation coefficients for the relations between the (PAH 8 $\mu$ m)/160 $\mu$ m ratio and 160 $\mu$ m surface brightness generally has a higher absolute value than the corresponding correlation coefficients for the relation between the (PAH 8 $\mu$ m)/160 $\mu$ m ratio and radius."
 These results suggest that the ratio may be more strongly dependent on 160 sam surface brightness than radius., These results suggest that the ratio may be more strongly dependent on 160 $\mu$ m surface brightness than radius.
 Since the (PALES pim)/160 pim ratio may be dependent on dust heating. we also examine how the ratio is related to the 24 pim/160. jm ratio.," Since the (PAH 8 $\mu$ m)/160 $\mu$ m ratio may be dependent on dust heating, we also examine how the ratio is related to the 24 $\mu$ m/160 $\mu$ m ratio."
 Ehe 24 sam band. which traces 2100 Ix hot dust emission. increases faster than other infrared band. (including the PALL S and. 160 jam bands) as the illuminating radiation field increases (Dalectal.2001: 2007)..," The 24 $\mu$ m band, which traces $\gtrsim100$ K hot dust emission, increases faster than other infrared band (including the PAH 8 and 160 $\mu$ m bands) as the illuminating radiation field increases \citep{dhcsk01, ld01, dl07}. ."
 Phe 160 fia band is approximately directly. proportional to the TII ας. as discussed. in Section. 2.1..," The 160 $\mu$ m band is approximately directly proportional to the TIR flux, as discussed in Section \ref{s_data_band}."
 Therefore. the 24 μα160. sam ratio should be a reasonable indicator of dust heating.," Therefore, the 24 $\mu$ m/160 $\mu$ m ratio should be a reasonable indicator of dust heating."
 Lf the mass fraction of PALIS remains constantand if the (PALL S pm)/160 yam ratio is dependent on dust heating. then the (PALL S pam)/160 jum ratio should monotonically increase as the 24 pm160. jim ratio increases. although the slope may not necessarily be constant.," If the mass fraction of PAHs remains constantand if the (PAH 8 $\mu$ m)/160 $\mu$ m ratio is dependent on dust heating, then the (PAH 8 $\mu$ m)/160 $\mu$ m ratio should monotonically increase as the 24 $\mu$ m/160 $\mu$ m ratio increases, although the slope may not necessarily be constant."
 This comparison is similar to the direct comparison between the PALL S and 24 jun bands performed in Section 3.. but the normalisation with the 160 yaa band removes variations related to dust surface clensitv.," This comparison is similar to the direct comparison between the PAH 8 and 24 $\mu$ m bands performed in Section \ref{s_comp_pah24}, but the normalisation with the 160 $\mu$ m band removes variations related to dust surface density."
 Figure 7. shows how the (PALL 8 jum)/160. pm. ratio varies with the 24 j/im/160 jun ratio within 45 aresec regions in the sample galaxies., Figure \ref{f_pah160vs24160} shows how the (PAH 8 $\mu$ m)/160 $\mu$ m ratio varies with the 24 $\mu$ m/160 $\mu$ m ratio within 45 arcsec regions in the sample galaxies.
 Many of the 45 aresee regions that are relatively weak in 24 jim emission tend to show a tight) correspondence between the (PALL S sam)/160 im and 24 jm/160 pm ratios. but some regions with enhanced 24 pim emission appear far to the right of these curves.," Many of the 45 arcsec regions that are relatively weak in 24 $\mu$ m emission tend to show a tight correspondence between the (PAH 8 $\mu$ m)/160 $\mu$ m and 24 $\mu$ m/160 $\mu$ m ratios, but some regions with enhanced 24 $\mu$ m emission appear far to the right of these curves."
 The most spectacular example is NGC 3938. where the ratios for most of theregions closely follow a linear relation but theLIL region on the east side of the dise falls far to the right of this relation.," The most spectacular example is NGC 3938, where the ratios for most of theregions closely follow a linear relation but the region on the east side of the disc falls far to the right of this relation."
 The relation between the two ratios on the left side of the panels in Figure 7 suggest that the (PALL S yaa)/160 yam ratio is dependent on dust heating. but the outliers on the right in these panels show PALL S jin emission is not enhanced in areas with strong cust heating such asLLL regions.," The relation between the two ratios on the left side of the panels in Figure \ref{f_pah160vs24160} suggest that the (PAH 8 $\mu$ m)/160 $\mu$ m ratio is dependent on dust heating, but the outliers on the right in these panels show PAH 8 $\mu$ m emission is not enhanced in areas with strong dust heating such as regions."
 Further interpretation of these results is presented in the next section., Further interpretation of these results is presented in the next section.
 For another perspective on the nature of the variation in the (PALL8. μαι)100 jim ratio. we examined how the 160 jamPLR and (PATES pm)ELI ratios vary as a function of “PIR. surface brightness.," For another perspective on the nature of the variation in the (PAH8 $\mu$ m)/160 $\mu$ m ratio, we examined how the 160 $\mu$ m/TIR and (PAH 8 $\mu$ m)/TIR ratios vary as a function of TIR surface brightness."
 These are displaved inFigures S ancl 9.. with slopes for the best fit lines given in Table 7 and Table S. ," These are displayed inFigures \ref{f_160vstir} and \ref{f_pahvstir}, , with slopes for the best fit lines given in Table \ref{t_160vstir} and Table \ref{t_pahvstir}. ."
As stated in Section 2.1.. the 160 panyELI ratio should. gradually decrease as dust. temperatures. increase," As stated in Section \ref{s_data_band}, , the 160 $\mu$ m/TIR ratio should gradually decrease as dust temperatures increase"
Following:. the first⋅ discovery. of⋅ an extrasolar planet. around ; - ⋅↱≻↓↓⋖⋅⋏∙≟↿∖↳∖↓⋜↧∙∖⇁∪↓⋅⊾∖↽≺≥⋯⋅⇂∪∠↓≤⋗≤⋗⋅≻⊐⋡⋡↿↓↥⋖⋅⊔⊔⊔↓∣⋡∢⋅↓⋅∪⇂⋖⋅⇀∖∪↓≻↓⋜⋯∢⋅↥≱∖ known has risen⊀ to 429.,"Following the first discovery of an extrasolar planet around 51 Peg \citep{b20}, the number of exoplanets known has risen to 429."
" Although most of⋅ them are giant⊀ planets. the improvements: in. observational. techniques. have ensured that planets with. masses less than 15M,- have started being.. detected: with. racial. velocity. survey (c.g. Lovisetal.(2006):UciryBontils(2007):Ucdrv.etal.(2007):Alavorct (2009))) ancl eravitational. microlensinesing ssurvey (Beaulier(Beietal.2005)2005)."," Although most of them are giant planets, the improvements in observational techniques have ensured that planets with masses less than $15 M_{\oplus}$ have started being detected with radial velocity survey (e.g., \citet{b22,b23,b24,b25,b21}) ) and gravitational microlensing survey \citep{b26}."
. Although most of extrasolar planets so Lar discovered are giant planets. several statistical models for planetary erowth presented in the last vears suggest that a large number of small plancts who fail to have enough mass to start the eas accretion. onto the core exists.. Uda&Lin.2004:Miguel&Brunini2009:Morcdasinietal. 2009).. and has still not been able to be discovered. (Mordasini 2009).," Although most of extrasolar planets so far discovered are giant planets, several statistical models for planetary growth presented in the last years suggest that a large number of small planets who fail to have enough mass to start the gas accretion onto the core exists \citep{b9,b2,b28}, and has still not been able to be discovered \citep{b27}."
. At the time. several projects are in progress to detect --- ↿⋖⊾↓⊓⊾⊳∖↿↓⋯↓↓≻↓⋜⋯∢⋅⇂⊳∖⊳∖∖⋎⋖⊾⋖⋅⇀∖↓≻∢⊾≼∼⇂⇂↓⋯↿↿↓∐⋅∙∖⊔↓⋜↧∙∖∐⊔∠⇂⊔↓∪⊓⋅∟⋜⊔∣↓↕−. ] e ∙⊲∙size planets in a close-future. but today. the sample is not enough and we also have to rely on what we know fron our own Solar] System.. and through computational. models of⋅ planetary. formation.," At the time, several projects are in progress to detect terrestrial planets, we expect that they may find more Earth-size planets in a close-future, but today, the sample is not enough and we also have to rely on what we know from our own Solar System, and through computational models of planetary formation."
"⋅⊀ MEPhis evidence. supports the standard. scenario.. where terrestrialAPpsEN planets""n iuxaro formed).1 throughBn then next.""T cilcHEUAPnnt stages: 1) agelomeration of dust. particles through physical setung the protoplanetary 2) planetesimalcollisions andformation fromin grains in a thin midplanedisc. (CGoldreich&Ward1973:WoeidenschillingCuzzi1993).. 3) runaway (e.g.. Wokubo&Ida (1996))) and. oligarchic σα©&-Alakinoeye]1993:[S11EUWokubodIca1998). accumulationnOn)©1 of planctesimals to form and 4) Brant impact stage. where the embryos formed by protoplanetsoligarchic growth collide with one another to form planets (Wetherill1985)."," This evidence supports the standard scenario, where terrestrial planets are formed through the next different stages: 1) agglomeration of dust particles through physical collisions and setting in the protoplanetary disc, 2) planetesimal formation from grains in a thin midplane \citep{b31,b32}, 3) runaway (e.g., \citet{b33}) ) and oligarchic \citep{b4,b15} accumulation of planetesimals to form protoplanets and 4) giant impact stage, where the embryos formed by oligarchic growth collide with one another to form planets \citep{b30}."
. The final stage of terrestrial planetary formation is the particular importance as it has a deep ellect. on the final characteristics of the planets: mass. orbital ancl spin parameters.," The final stage of terrestrial planetary formation is the particular importance as it has a deep effect on the final characteristics of the planets: mass, orbital and spin parameters."
 After this stage of planetary formation. the spin," After this stage of planetary formation, the spin"
w=0 80 that ¢ can be independent ofr. and we must allow 4 to be either a function of x (latitude) only. or a function of z (radius) only.,"$w=0$ so that $v$ can be independent of $x$, and we must allow $u$ to be either a function of $x$ (latitude) only, or a function of $z$ (radius) only."
 There is evidence (Dikpati et al 2005) that the latitude eradient of rotation is more important (han the radial gradient in the flux transport dvnanmos that best simulate solar cycles. so in (his study we restrict ourselves (o consideration of the latitude gradient.," There is evidence (Dikpati et al 2005) that the latitude gradient of rotation is more important than the radial gradient in the flux transport dynamos that best simulate solar cycles, so in this study we restrict ourselves to consideration of the latitude gradient."
 The difusivity 7 and the a-effect are also taken to be independent ofr., The diffusivity $\eta$ and the $\alpha$ -effect are also taken to be independent of $x$.
 Equations (3) aud (4) then reduce to There are a varietv of boundary conditions that could be chosen lor the top and bottom of the infinite plane laver., Equations (3) and (4) then reduce to There are a variety of boundary conditions that could be chosen for the top and bottom of the infinite plane layer.
 Consistent with solar conditions. we take (he bottom to be a perlect conductor. and therefore require A=0 there.," Consistent with solar conditions, we take the bottom to be a perfect conductor, and therefore require $A=0$ there."
 B in that case is determined internally., $B$ in that case is determined internally.
 For the top there are four plausible alternatives: perfect conductor (Al=0 again)knsulator with no forcing (5=0 and Ji matched to potential field above): forcing in potential at. top (A= Àj) and B=0 or determined internally., For the top there are four plausible alternatives: perfect conductor $A=0$ again);insulator with no forcing $B=0$ and $A$ matched to potential field above); forcing in potential at top $A=A_F$ ) and $B=0$ or determined internally.
 We make choices among these boundary conditions when we derive the 1- and 2-Iaver equations in the next section., We make choices among these boundary conditions when we derive the 1- and 2-layer equations in the next section.
 In preparation for these derivations. we simplify the dynamo equations further by taking ou=s.v.0.1 all independent of x. and assuming all variables have solutions of the form elke90 eanations (5) and (6) reduce to Belore developing 1. 2 and 3-Iaver clvnamo equations in detail. we first describe schematically What variables and parameters are retained in (hese cases.," In preparation for these derivations, we simplify the dynamo equations further by taking ${\partial u \over \partial x}=s,v,\alpha,\eta $ all independent of x, and assuming all variables have solutions of the form $e^{i(kx-{\omega}t)}$ equations (5) and (6) reduce to Before developing 1, 2 and 3-layer dynamo equations in detail, we first describe schematically what variables and parameters are retained in these cases."
 These are summarized graphically, These are summarized graphically
line signatures of ACN accretion.,line signatures of AGN accretion.
 Optically dull ACNs make up ~15% of luminous (Ly>LOY ere/s) point sources in deep X-rav surveys. and ~25% of those at 2<1 CDuuupctal.2009a:Trouilleet2009:Yanal. 2010).," Optically dull AGNs make up $\sim$ of luminous $L_X>10^{42}$ erg/s) point sources in deep X-ray surveys, and $\sim$ of those at $z<1$ \citep{tru09a,tro09,yan10}."
. Three major paracdiguis exist to explain their A-rav brightuess aud optical dullness: (1) obscuration of both narrow aud broad emüsson lines. (2) dilution bv host galaxy starlight. aud (3) a plysically distinct accretioun flow.," Three major paradigms exist to explain their X-ray brightness and optical dullness: (1) obscuration of both narrow and broad emission lines, (2) dilution by host galaxy starlight, and (3) a physically distinct accretion flow."
 We explore the evidence for aud against each paradigm in turn., We explore the evidence for and against each paradigm in turn.
 The standard ACN unified model has had great success Tn using obscuration to explain the differences between optically selected Type 1 (broad-Iue) aud Type 2 (narrow-line) ACNs (τοις&Degelhuan1988:Au-touucci 1993).," The standard AGN unified model has had great success in using obscuration to explain the differences between optically selected Type 1 (broad-line) and Type 2 (narrow-line) AGNs \citep{kro88,ant93}."
. Iu the simplest interpretation of this model. all AGNs have the broad optical eiuission Hines and strong UV/optical continua of Type 1. ACNs. but along certain lines of sight these features are obscured within a dusty “torus” a few parsecs frou the black hole.," In the simplest interpretation of this model, all AGNs have the broad optical emission lines and strong UV/optical continua of Type 1 AGNs, but along certain lines of sight these features are obscured within a dusty “torus” a few parsecs from the black hole."
 Iu au obscured object the narrow cussion lines τοπ visible because they are excited bevond the obscuring nmateral. and so the lack of a BLR aud weaker optical coutimmuun of Type 2 AGNs are attributed to obscuration.," In an obscured object the narrow emission lines remain visible because they are excited beyond the obscuring material, and so the lack of a BLR and weaker optical continuum of Type 2 AGNs are attributed to obscuration."
 Similarly. Comastrietal.(2002) and Civanoctal.(2007)— sueeested that optically dull," Similarly, \citet{com02} and \citet{civ07} suggested that optically dull"
Tipp. E.B.Jenkins? Weak Lyinan alpha (να) absorbers that are now detectable in IIST/STIS quasar spectra are predicted to arise in diffuse non-equilibrium larec-scale structures. analogous to high-redshift forest absorbers|1]: this has beeu preliminarily confirmed by first comparisous ofthe statistical properties of aabsorbers with sinmlatious|2].,", E. B. $^5$ } Weak Lyman alpha ) absorbers that are now detectable in HST/STIS quasar spectra are predicted to arise in diffuse non-equilibrium large-scale structures, analogous to high-redshift forest absorbers; this has been preliminarily confirmed by first comparisons of the statistical properties of absorbers with simulations."
. At high redshift. absorbers with Lotcin? are already euriched to a level of [C/II].=—2.5l.," At high redshift, absorbers with $\nh\approx
10^{14.5}\cdunits$ are already enriched to a level of $_\odot \approx
-2.5$."
. Tt is of great interest to study similar absorbers at low redshift. to determine the rate at which metals have been injected into the IGAL since 2—3.," It is of great interest to study similar absorbers at low redshift, to determine the rate at which metals have been injected into the IGM since $z\sim 3$."
" The results have implications for wucerstanding carly galaxy formation. associated blowout. :id πείς,"," The results have implications for understanding early galaxy formation, associated blowout, and winds."
 To: study absorbers equivalent: to τμιzm+1Tn2 atic-3. requires. oue to exaniiuc. absorbers with. Nyy1=MNLO?’ytcm57 at zzm0 |1]., To study absorbers equivalent to $\nh\approx 10^{14.5}\cdunits$ at $z=3$ requires one to examine absorbers with $\nh\approx 10^{13}\cdunits$ at $z\approx 0$ .
. Unfortunately.v. siuce," Unfortunately, since"
. Unfortunately.v. siuce.," Unfortunately, since"
example. in the SAIC (Glattetal.2008).,"example, in the SMC \citep{gl08}."
. At such a time. the study by would have revealed little variation in IIB morphology among the halo GC's and. hence. little or no insight into the formation of the Galaxy.," At such a time, the study by \citet{sz78} would have revealed little variation in HB morphology among the halo GCs and, hence, little or no insight into the formation of the Galaxy."
 Only the (truly. peculiar GC's. such as NGC 2808 (Dalessandroοἱal.2010).. might have shown a sizable variation.," Only the truly peculiar GCs, such as NGC 2808 \citep{da10}, might have shown a sizable variation."
 On the other hand. turn the elock forward ~2 Gyr and all but the most metal-rich Galactic GCs would have blue IIDs.," On the other hand, turn the clock forward $\sim2$ Gyr and all but the most metal-rich Galactic GCs would have blue HBs."
 IIST/ACS photometry of 6 Galactic halo GCs. IC 4499. NGC 6426. NGC 7006. Palomar 15. Pyxis. and Ruprecht 106. were presented.," HST/ACS photometry of 6 Galactic halo GCs, IC 4499, NGC 6426, NGC 7006, Palomar 15, Pyxis, and Ruprecht 106, were presented."
 The resulüing CMDs were used to derive ages via isochrone fitting., The resulting CMDs were used to derive ages via isochrone fitting.
 The age of Palomar 5 was derived using the same set of models and the combined photometry of Grillmair&S$qmüth(2001) ancl Stetson(2000)., The age of Palomar 5 was derived using the same set of models and the combined photometry of \citet{gr01} and \citet{st00}.
. The ages and metallieities of these 7 GCs were added to the AMIR of Dotteretal.(2010)., The ages and metallicities of these 7 GCs were added to the AMR of \citet{do10}.
. The total inuuber of homogeneously studied GCs in the sample is 68. excluding those 7 GCs imaged bv the ACS Survev of Galactic GCs that are known to harbor multiple. clistinet stellar populations.," The total number of homogeneously studied GCs in the sample is 68, excluding those 7 GCs imaged by the ACS Survey of Galactic GCs that are known to harbor multiple, distinct stellar populations."
 Divided at Ree=8 κρο the inner Galaxy GCs exhibit a pattern of rapid chemical enrichment spanning two orders of magnitude in metallicity over a timescale of 1-2 Gyr.," Divided at $\rgc = 8$ kpc, the inner Galaxy GCs exhibit a pattern of rapid chemical enrichment spanning two orders of magnitude in metallicity over a timescale of 1-2 Gyr."
 The outer Galaxy. GCs exhibit a much slower chemical enrichment that is evocative of dwarl galaxy chemical evolution ancl is reasonably well matehecl by the semi-analvtie GC formation models of Muratov&Gneclin(2010)., The outer Galaxy GCs exhibit a much slower chemical enrichment that is evocative of dwarf galaxy chemical evolution and is reasonably well matched by the semi-analytic GC formation models of \citet{mg10}.
. Data presented in this paper were obtained from the Multimission Archive at the Space Telescope Science Institute (ALAST)., Data presented in this paper were obtained from the Multimission Archive at the Space Telescope Science Institute (MAST).
 Support for this work (proposal GO-11536) was provided by NASA through a grant. trom the Space Telescope Science Institute. which is operated by the Association of Universities for Research in Astronomy. Inc.. under NASA contract NÀS5-26555.," Support for this work (proposal GO-11586) 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, Inc., under NASA contract NAS5-26555."
 This research has made use of NASAs Astrophysics Data System Bibliographic Service as well as the SIMDBAD database. operated at CDS. Strasbourg. France.," This research has made use of NASA's Astrophysics Data System Bibliographic Service as well as the SIMBAD database, operated at CDS, Strasbourg, France."
No significant correlation is found between the Karlsson relation and cither of the observed redshift distributions.,No significant correlation is found between the Karlsson relation and either of the observed redshift distributions.
" However. there does appear o be a significant correlation between the redshift distribution observed for the QSOs near NGC 6212 and the redshift distribution predicted bv he zi, relation. for No — 3."," However, there does appear to be a significant correlation between the redshift distribution observed for the QSOs near NGC 6212 and the redshift distribution predicted by the $_{iQ}$ relation, for $N$ = 3."
 No evidence of a correlation was found for N-values of 1. 2. 4. or higher.," No evidence of a correlation was found for $N$ -values of 1, 2, 4, or higher."
 A significant correlation is also found vetween the distribution of 571 quasar redshifts discussed by Iarlsson aud the redshifts predicted w othe ze relation (for No = 2 aud 3)., A significant correlation is also found between the distribution of 574 quasar redshifts discussed by Karlsson and the redshifts predicted by the $_{iQ}$ relation (for $N$ = 2 and 3).
 For the QSOs near NGC 6212 there is also evidence that here is indeed a redshift compoucu due to the cosinological distance of that galaxy prescut iu the data., For the QSOs near NGC 6212 there is also evidence that there is indeed a redshift component due to the cosmological distance of that galaxy present in the data.
 However. the optiunun value appears to be closer to z. = 0.02 sugeesting that the observed redshift of NCC 6212 nav also contain a small intrinsic redshift component.," However, the optimum value appears to be closer to $_{c}$ = 0.02 suggesting that the observed redshift of NGC 6212 may also contain a small intrinsic redshift component."
 We conclude on the basis of these results. and the points discussed above. that although the Iarlsson relation may predict some values tji align with peaks in the observed redshift distributions. there is little evidence to sugeest that this equation is related to. or has any bearing on. he production of iutrinsic redshifts in quasars.," We conclude on the basis of these results, and the points discussed above, that although the Karlsson relation may predict some values that align with peaks in the observed redshift distributions, there is little evidence to suggest that this equation is related to, or has any bearing on, the production of intrinsic redshifts in quasars."
Sunvaev 1973: Novikov Thorne 1974: Page Thorne 974) is one in which the accretion disk is geometrically-thin and raciativelv-ellicient. extends from the radius of marginal stability to large μασ. and with an iron line emissivity that racks the underlying dissipation (Itevnolds Nowak 2003: tevnolds ct al.,"Sunyaev 1973; Novikov Thorne 1974; Page Thorne 1974) is one in which the accretion disk is geometrically-thin and radiatively-efficient, extends from the radius of marginal stability to large radii, and with an iron line emissivity that tracks the underlying dissipation (Reynolds Nowak 2003; Reynolds et al."
 2003)., 2003).
 Applying such a line profile to the GC 4593 data in the case of a near-extreme Ixerr. black ole (with dimensionless spin parameter e=0.998) results in an upper limit to the equivalent width of90eV., Applying such a line profile to the NGC 4593 data in the case of a near-extreme Kerr black hole (with dimensionless spin parameter $a=0.998$ ) results in an upper limit to the equivalent width of $99\eV$.
 These are significantly less than the values expected from theoretical reflection moclels (200eV for solar abundances: e.g. Matt. Fabian Reynolds 1997 and references therein) or observed in the Sevlert galaxy AICG6-30-15 (~400eV: Fabian et al.," These are significantly less than the values expected from theoretical reflection models $\sim 200\eV$ for solar abundances; e.g., Matt, Fabian Reynolds 1997 and references therein) or observed in the Seyfert galaxy MCG–6-30-15 $\sim 400\eV$; Fabian et al."
 2002)., 2002).
 Thus. there appears to be a significant absence of spectral features from a relativistic accretion disk.," Thus, there appears to be a significant absence of spectral features from a relativistic accretion disk."
 As discussed in the Introduction. the black hole accretion paradigm of AGN is very well established and supported by a significant body of evidence.," As discussed in the Introduction, the black hole accretion paradigm of AGN is very well established and supported by a significant body of evidence."
 Phus. the results of Section beg us to turn the question around: why are we not seeing the X-ray reflection signatures of a relativistic aceretion disk. given that we believe such a disk exists and is responsible for all of the AGN emissions that we observe?," Thus, the results of Section \ref{sec:results} beg us to turn the question around: why are we not seeing the X-ray reflection signatures of a relativistic accretion disk, given that we believe such a disk exists and is responsible for all of the AGN emissions that we observe?"
 A straightforward solution to this problem would be to hypothesize sub-solar abundances of iron in the black hole aceretion clisk., A straightforward solution to this problem would be to hypothesize sub-solar abundances of iron in the black hole accretion disk.
 Lf! the light5 elements are present in cosmic abundance: (Anclers Gerevesse 1989). one needs iron to be under-abundant bv a factor of 3 (he. Zee< 0.3) in order to reduce the (cold) broad [uorescent emission. below the LOOceVY leve as required. by our data (e.g. see Revnolcs. Fabian. Inoue 1995).," If the light elements are present in cosmic abundance (Anders Grevesse 1989), one needs iron to be under-abundant by a factor of 3 (i.e., $Z_{\rm Fe}< 0.3$ ) in order to reduce the (cold) broad fluorescent emission below the eV level as required by our data (e.g., see Reynolds, Fabian Inoue 1995)."
 If the light. elements are.overabiumndant. the enhanced photoelectrie absorption further decreases the iron line equivalent. width.," If the light elements are, the enhanced photoelectric absorption further decreases the iron line equivalent width."
 A light element overabundance by a [actor of two reduces the required iron underabundance to only (Le. Zpé« 0.7).," A light element overabundance by a factor of two reduces the required iron underabundance to only (i.e., $Z_{\rm Fe}< 0.7$ )."
 However. it woulcl be surprising if the solution to the lack of a broad iron line was simply an underabundance of iron. given the highly evolved nature of stellar populations in the nuclei of galaxies such as NGC 4593.," However, it would be surprising if the solution to the lack of a broad iron line was simply an underabundance of iron, given the highly evolved nature of stellar populations in the nuclei of galaxies such as NGC 4593."
" With this motivation. we explore modifications. of the ""canonical"" line models discussed above ancl show that it is. in fact. rather easy to bury relativistic spectral features in the noise even i£ they are present at the level associated with cosmic abundance material."," With this motivation, we explore modifications of the “canonical” line models discussed above and show that it is, in fact, rather easy to bury relativistic spectral features in the noise even if they are present at the level associated with cosmic abundance material."
 Wilms et al. (, Wilms et al. (
2001) ancl Hevnolds et al. (,2001) and Reynolds et al. (
2004). have analyzed the IEEPIC-pn spectrum of MC?6-30-15 in its Deep Minimum state and found it to possessvery broadened X-rav rellection features.,2004) have analyzed the EPIC-pn spectrum of MCG–6-30-15 in its Deep Minimum state and found it to possess broadened X-ray reflection features.
 On the basis of these data. they suggest that the accretion disk in ALCG6-30-15 is being torqued by interactions with the central spinning black hole. producing a dissipation that is very centrally concentrated.," On the basis of these data, they suggest that the accretion disk in MCG–6-30-15 is being torqued by interactions with the central spinning black hole, producing a dissipation that is very centrally concentrated."
 ὃν emploving the disk models of gol Ixrolik. (2001). ltevnolds ct al. (," By employing the disk models of Agol Krolik (2001), Reynolds et al. ("
2004) suggest that the Deep Minimum state of ALCG6-30-15 corresponds to a torque-dominated (cinlinite-ellicienev: see Agol Ixrolik 2001) accretion disk viewed at an inclination of 30.—40°,2004) suggest that the Deep Minimum state of MCG–6-30-15 corresponds to a torque-dominated (“infinite-efficiency”; see Agol Krolik 2001) accretion disk viewed at an inclination of $30-40^\circ$.
 There is little evidence for an obscuring molecular torus in AICG6-30-15 (e.g... see Lee et al.," There is little evidence for an obscuring molecular torus in MCG–6-30-15 (e.g., see Lee et al."
 2002) and so. in principle. the central accretion clisk could be viewed at any angle.," 2002) and so, in principle, the central accretion disk could be viewed at any angle."
 Ht is interesting to note that if ALCG6-80-15 were observed at a high inclination with only moderate signal-to-noise. the Deep Minimum iron line would be so broad as to be undetectable against the noise of the continuum.," It is interesting to note that if MCG–6-30-15 were observed at a high inclination with only moderate signal-to-noise, the Deep Minimum iron line would be so broad as to be undetectable against the noise of the continuum."
 To explore whether this is also the case in NGC 4593. we add to the spectral fit a cold. iron line with a profile corresponding to an infinite-ellicieney accretion disk around a near-extremoe Ixerr black hole (α= 0.998).," To explore whether this is also the case in NGC 4593, we add to the spectral fit a cold iron line with a profile corresponding to an infinite-efficiency accretion disk around a near-extreme Kerr black hole $a=0.998$ )."
 The inclination and normalization of the line were left as free parameters., The inclination and normalization of the line were left as free parameters.
 This leads to only a slight improvement in the goodness of fit over the simple power-law plus narrow line model (Ay?—3 for 2 additional degrees of freedom)., This leads to only a slight improvement in the goodness of fit over the simple power-law plus narrow line model $\Delta\chi^2=3$ for 2 additional degrees of freedom).
 Since this line is so broad. the upper limit on the equivalent width is Ίρις 250eV.," Since this line is so broad, the upper limit on the equivalent width is $W_{\rm
 broad}<250\eV$ ."
 Thus. à broad line with the strength expected from a cosmic abundance accretion clisk is consistent with these data if the emissivity profile is very centrally concentrated.," Thus, a broad line with the strength expected from a cosmic abundance accretion disk is consistent with these data if the emissivity profile is very centrally concentrated."
 The prominence of X-ray reflection features can be significantly reduced. by ionization of the aceretion disk, The prominence of X-ray reflection features can be significantly reduced by ionization of the accretion disk
evolved morphologically because they began to interact at a earlier tine?,evolved morphologically because they began to interact at a earlier time?
 The multiple evidence of interactions in the CC. both iu the optical and NIR. suggest these galaxies are clearly not iu equilibrimu.," The multiple evidence of interactions in the CGs, both in the optical and NIR, suggest these galaxies are clearly not in equilibrium."
 Therefore. compact eroups could not have formed that long ago iu the past.," Therefore, compact groups could not have formed that long ago in the past."
 For the pairs of galaxies. we may casily assuiue that the interactions are relatively receut.," For the pairs of galaxies, we may easily assume that the interactions are relatively recent."
 These ealaxies formed im low deusity cuvirouments aud it took an IInubble time to two of them for iicet and interact., These galaxies formed in low density environments and it took an Hubble time to two of them for meet and interact.
 This iuterpretation is consistent with our observations., This interpretation is consistent with our observations.
 Iu the NPCs. the stellar populations in the central region of the ealaxies with ciffercut morplologics secu vouuger. in geueral. than in the IICCs.," In the KPGs, the stellar populations in the central region of the galaxies with different morphologies seem younger, in general, than in the HCGs."
 This is consistent with the spectroscopic evidence. which shows a higher level of star formation in the IKPCs compared to the ICCs.," This is consistent with the spectroscopic evidence, which shows a higher level of star formation in the KPGs compared to the HCGs."
 Oue key differcuce between the IICCs aud θές seems to be that iu the CGs. interactions are happening under dry conditions. confriuius what we observed iu Coziol Plauchu-Fravu (2007).," One key difference between the HCGs and KPGs seems to be that in the CGs, interactions are happening under dry conditions, confirming what we observed in Coziol Plauchu-Frayn (2007)."
 The fact that we do not see such phenomenon in the WKPGs could be due to interactions iu hese systems are very recent (this is supported w spectroscopy)., The fact that we do not see such phenomenon in the KPGs could be due to interactions in these systems are very recent (this is supported by spectroscopy).
 In the CCs. a first round of interactions would have produced the nuuerous carly-type galaxies we now observe.," In the CGs, a first round of interactions would have produced the numerous early-type galaxies we now observe."
 Possibly when hey formed. the CGs would have experieuced a nore active phase of star formation (aud ACN).," Possibly when they formed, the CGs would have experienced a more active phase of star formation (and AGN)."
 But now that the eas is exhausted. the galaxies iu Cs. assimuiue mereie orbits. can onlv interact uuder dev couditious.," But now that the gas is exhausted, the galaxies in CGs, assuming merging orbits, can only interact under dry conditions."
 For the NPCs. we do uot show what will be their future.," For the KPGs, we do not know what will be their future."
 Possibly those are systems with Ligh energv orbits. that will interact again onlv after an extremely lounge time jas passed.," Possibly those are systems with high energy orbits, that will interact again only after an extremely long time has passed."
 In the case of the ς the evidence of dev interaction conditions would thus be an evidence that galaxies in these systems are now in qncreing orbits., In the case of the CGs the evidence of dry interaction conditions would thus be an evidence that galaxies in these systems are now in merging orbits.
 Cousequeuthy. their formation cannot be that far iu the past.," Consequently, their formation cannot be that far in the past."
 Qur analysis sugeests that pairs of galaxies are vouug structures: the ealaxics in pairs formed and speut most of their life iu relative isolation and are just beginning to interact., Our analysis suggests that pairs of galaxies are young structures: the galaxies in pairs formed and spent most of their life in relative isolation and are just beginning to interact.
 This behavior would be typical of low deusity euviromneuts. or what is found at the periphery of large scale structures.," This behavior would be typical of low density environments, or what is found at the periphery of large scale structures."
 Ou the other hand. galaxies in Cs are obviously more evolved morphologically aud have suffered nmiore interactions.," On the other hand, galaxies in CGs are obviously more evolved morphologically and have suffered more interactions."
 ILowever. based on he abundant evidence of interactions iu both he NIR aud the optical. these svstems cannot xo iu equilibrium.," However, based on the abundant evidence of interactions in both the NIR and the optical, these systems cannot be in equilibrium."
 In particular. the evidence for drv duteractions in CCCs is consistent with the wpothesis that galaxies are in imereie orbits.," In particular, the evidence for dry interactions in CGs is consistent with the hypothesis that galaxies are in merging orbits."
 Consequeutlv. ως cannot be extremely old.," Consequently, CGs cannot be extremely old."
 Cosmologically speaking the difference in formation iue between pairs and C'CGs may be relatively small., Cosmologically speaking the difference in formation time between pairs and CGs may be relatively small.
 That is. the two phenomena are probably vpical of the formation of structures in low density environments and consequently their respective formation represcuts relatively recent events compared to the formation of larger aud nuüore nissive structures.," That is, the two phenomena are probably typical of the formation of structures in low density environments and consequently their respective formation represents relatively recent events compared to the formation of larger and more massive structures."
 According to this interpretation. one would not expect svstenis like local CC to exist at high redshifts.," According to this interpretation, one would not expect systems like local CGs to exist at high redshifts."
 CC may have formed in the past. but these would have been nmch more massive than what we find today aud such svstems would have be expected to merge with others to form cluster of galaxies (Cozioletal.2009).," CGs may have formed in the past, but these would have been much more massive than what we find today and such systems would have be expected to merge with others to form cluster of galaxies \citep{coziol09}."
. We thank the CATT of San Pedro Mürrtir for the observing tine eiven on the 2.limn telescope to realizo this project and all the personel of the observatory for their support., We thank the CATT of San Pedro Márrtir for the observing time given on the 2.1m telescope to realize this project and all the personel of the observatory for their support.
 We also thank an anonviuous referee for iuportant conuueuts and sugeestious., We also thank an anonymous referee for important comments and suggestions.
 I. P. F. acknowledges to Drs., I. P. F. acknowledges to Drs.
 II. Audernach and J. M. Islas-Islas for their valuable feedback., H. Andernach and J. M. Islas-Islas for their valuable feedback.
" This research has made use of: 1) SAOTmage DS9. developed by Sinithsouian Astrophysical Observatory: 2) TOPCAT software provided by the UN's AstroGrid Virtual Observatory Project. which is funded by the Science aud Technology Facilities Council and through the EUs Framework 6 programme: 3) Data products from the Two Micron. All Skv. Survey, which is a joint project of the University of Massachusetts and the Dufraved Processing aud Analysis Center/California Tustitute of Technology. fuuded by the National Aeronautics and Space Adiuiuistration aud the National Scicuce Foundation: I) Funudiug for the SDSS aud SDSS-II has been provided bv the Alfred P. Sloan Foundation. the Participating"," This research has made use of: 1) SAOImage DS9, developed by Smithsonian Astrophysical Observatory; 2) TOPCAT software provided by the UK's AstroGrid Virtual Observatory Project, which is funded by the Science and Technology Facilities Council and through the EU's Framework 6 programme; 3) Data products from the Two Micron All Sky Survey, which is a joint project of the University of Massachusetts and the Infrared Processing and Analysis Center/California Institute of Technology, funded by the National Aeronautics and Space Administration and the National Science Foundation; 4) Funding for the SDSS and SDSS-II has been provided by the Alfred P. Sloan Foundation, the Participating"
2011).,.
" A similar extension of the fast mode solutions to frequencies higher than w,; is shown in Fig. 8..", A similar extension of the fast mode solutions to frequencies higher than $\omega_{ci}$ is shown in Fig. \ref{whamp3}.
" We observe frequency gaps in the dispersion curves caused by the ion Bernstein modes, known to develop in hot plasmas at quasi-perprendicular angles of propagation Li&Habbal (2001)))."," We observe frequency gaps in the dispersion curves caused by the ion Bernstein modes, known to develop in hot plasmas at quasi-perprendicular angles of propagation (e.g., \cite{li01}) )."
" The more oblique Bernstein modes(e.g., are the least damped as shown in Fig.", The more oblique Bernstein modes are the least damped as shown in Fig.
 7 (lower panel)., \ref{damping} (lower panel).
" This result clearly invalidates the magnetosonic dispersion w~kC, of the classical-whistler mode at ω>wei found in the previous section from the hot two-fluid theory.", This result clearly invalidates the magnetosonic dispersion $ \omega\sim kC_s$ of the classical-whistler mode at $\omega >\omega_{ci}$ found in the previous section from the hot two-fluid theory.
" Figure 7 shows clearly that, while the Bernstein and the fast modes are less damped than the Alfvénn-whistler modes near kp;~1, both the Bernstein and the fast magnetosonic modes become heavily damped at kp; 3."," Figure \ref{damping} shows clearly that, while the Bernstein and the fast modes are less damped than the Alfv\'enn-whistler modes near $k\rho_i\sim1$, both the Bernstein and the fast magnetosonic modes become heavily damped at $k\rho_i\gtrsim 3$ ."
 The Alfvénn-whistler modes appear to be the less damped2 ones at scales kp;>3., The Alfvénn-whistler modes appear to be the less damped ones at scales $k\rho_i\geq3$.
" Furthermore, among these modes the most oblique one is the least damped, which is the KAW mode at (kg=89.99? in Fig. 7.."," Furthermore, among these modes the most oblique one is the least damped, which is the KAW mode at $\theta_{\bf kB}=89.99^\circ$ in Fig. \ref{damping}. ."
 In Fig., In Fig.
 9 we tested the validity of these results for different high 6; values and with 7T;=Τε., \ref{whamp4} we tested the validity of these results for different high $\beta_i$ values and with $T_i=T_e$.
 We see clearly that the conclusions above remain valid and that the ratio T;/T. does not modify much the results but slight changes of the dispersion curves near the harmonics of ions., We see clearly that the conclusions above remain valid and that the ratio $T_i/T_e$ does not modify much the results but slight changes of the dispersion curves near the harmonics of ions.
 This can be seen for instance by comparing Fig., This can be seen for instance by comparing Fig.
 9 (middle panel) to Fig., \ref{whamp4} (middle panel) to Fig.
" 3 which was plotted for the same 6;=2.9 but with 7T;=5T,.", \ref{whamp1} which was plotted for the same $\beta_i=2.9$ but with $T_i=5T_e$.
" It is important however to note that when 8; is becoming smaller (i.e. decreasing toward 1) the damping of Alfvénn-whistler modes by cyclotron resonances become important for some angles of propagation, as can be seen in Fig."," It is important however to note that when $\beta_i$ is becoming smaller (i.e. decreasing toward 1) the damping of Alfvénn-whistler modes by cyclotron resonances become important for some angles of propagation, as can be seen in Fig."
 9 (top , \ref{whamp4} (top panel).
We tested other values of 8; shown and found panel).that the conclusions above remain(not valid., We tested other values of $\beta_i$ (not shown here) and found that the conclusions above remain valid.
" here)Moreover, we found that higher values of 8; allow extending the modes to even smaller scales 8;=8.25, the highest tested value, the modes were (forfound to extend up to kp;~ 30)."," Moreover, we found that higher values of $\beta_i$ allow extending the modes to even smaller scales (for $\beta_i=\beta_e=25$, the highest tested value, the modes were found to extend up to $k\rho_i\sim 30$ )."
" On the contrary, for 8;~1 we found that none of the Alfvénn-whistler modes (when 80°«0< can propagate at w>cw. because they are strongly 90°)damped at wcwa."," On the contrary, for $\beta_i\sim 1$ we found that none of the Alfvénn-whistler modes (when $80^\circ\leq\theta<90^\circ$ ) can propagate at $\omega>\omega_{ci}$ because they are strongly damped at $\omega\sim \omega_{ci}$."
" Now one can ask the question: which of these plasma modes, fast, Bernstein, KAW or whistler, is likely to carry the energy cascade of turbulence down to the dissipation scales in the limited range of SW parameters studied here?"," Now one can ask the question: which of these plasma modes, fast, Bernstein, KAW or whistler, is likely to carry the energy cascade of turbulence down to the dissipation scales in the limited range of SW parameters studied here?"
" While all these modes might contribute to the energy cascade in the SW,based on lineardamping rates and assuming that (quasi-)linear theory is applicable to small scale SW fluctuations, we can conclude that the KAW branch (tw« w,;) is more likely"," While all these modes might contribute to the energy cascade in the SW,based on lineardamping rates and assuming that (quasi-)linear theory is applicable to small scale SW fluctuations, we can conclude that the KAW branch $\omega<\omega_{ci}$ ) is more likely"
should allow one to use the co-adding approach with very long integration times (e.g. for all the scans of a given star in a night).,should allow one to use the co-adding approach with very long integration times (e.g. for all the scans of a given star in a night).
 We are currently developing such an approach., We are currently developing such an approach.
 We show the results of an astrometric fit to 45 minutes of data (taken over the course of TO minutes of observation) in Figures 4 and 6..," We show the results of an astrometric fit to 45 minutes of data (taken over the course of 70 minutes of observation) in Figures \ref{fig:fit}
 and \ref{fig:error}."
 With 1769 scans used. we find the residual delay errors to be well modeled by a Gaussian distribution with a full-width at hall-maximum of 0.160 jim. In order (o characterize the residuals. and in particular determine if they could be considered to be independent. we plot the Allan. variance of the residuals in Figure 5..," With 1769 scans used, we find the residual delay errors to be well modeled by a Gaussian distribution with a full-width at half-maximum of 0.160 $\mu$ m. In order to characterize the residuals, and in particular determine if they could be considered to be independent, we plot the Allan variance of the residuals in Figure \ref{fig:allan}. ."
 The Allan variance (Thompsonetal.2001) at lag / is given bv where Af’ is the total number of data points., The Allan variance \citep{tms01} at lag $l$ is given by where $M'$ is the total number of data points.
 As can be seen in the figure. the residuals are white oul to lags of more than 500 scans. implying a final astrometric precision of LO µας.," As can be seen in the figure, the residuals are white out to lags of more than 500 scans, implying a final astrometric precision of 10 $\mu$ as."
 We list the results [rom 4 nights of observation in Table 1.., We list the results from 4 nights of observation in Table \ref{tab:res}. .
" The I-sigma error region (found by plotting the V?=\7,,,,4-NA? contour. Press et al. 1992))"," The 1-sigma error region (found by plotting the $\chi^2 = \chi^2_{min}
+ \Delta\chi^2$ contour, Press et al. \nocite{press92}) )"
 is hiehly elliptical with the major axis oriented roughly parallel to the RÀ. axis., is highly elliptical with the major axis oriented roughly parallel to the R.A. axis.
 Such error ellipses are to be expected in single-baseline interferometric data. which has limited sensitivity in the direction perpendicular to the baseline.," Such error ellipses are to be expected in single-baseline interferometric data, which has limited sensitivity in the direction perpendicular to the baseline."
 It should be noted however that for sullicientlv long observations. Earth-rotation will provide an orthogonal baseline.," It should be noted however that for sufficiently long observations, Earth-rotation will provide an orthogonal baseline."
 The major and minor axes of the uncertainty ellipse are easily found by diagonalizing the covariance matrix: (he magnitude of the uncertainty in the direction of (he minor axis was 8.3 µας for the 10 August data and 12 pas for the 11. August data: consistent with the delay residuals., The major and minor axes of the uncertainty ellipse are easily found by diagonalizing the covariance matrix: the magnitude of the uncertainty in the direction of the minor axis was 8.3 $\mu$ as for the 10 August data and 12 $\mu$ as for the 11 August data; consistent with the delay residuals.
 The uncertainty in the major axis direction was 144 and 143. jas respectively., The uncertainty in the major axis direction was 144 and 143 $\mu$ as respectively.
 Figure 6 shows the four errorellipses superimposed., Figure \ref{fig:error} shows the four errorellipses superimposed.
 We fit (he measured differential declinations toa, We fit the measured differential declinations toa
Detailed observations of the interstellar mec (ISAT) or even (the intergalactic medium have hiehliehted the need to provide a description (hat accounts for the turbulent. pressure. thermal pressure. magnetic fields. rotation. and even stars themselves.,"Detailed observations of the interstellar medium (ISM) or even the intergalactic medium have highlighted the need to provide a description that accounts for the turbulent pressure, thermal pressure, magnetic fields, rotation, and even stars themselves."
 A universal theory. describing the complex structures of (he ISM. is far from complete and remains a challenge for the future.," A universal theory, describing the complex structures of the ISM, is far from complete and remains a challenge for the future."
 However. by dividing the ISM according to its properties. il is possible to present satisfactory (theories for particular tvpes of ISM.," However, by dividing the ISM according to its properties, it is possible to present satisfactory theories for particular types of ISM."
 Gravitation. cooling. turbulence. and magnetic fields produce variations in (le properties of the ISM such as the clensity and," Gravitation, cooling, turbulence, and magnetic fields produce variations in the properties of the ISM such as the density and"
the available data.,the available data.
 A moclel in which the single-tompcrature dust. clouds discussed here are replaced. by a temperature distribution will probably be required. to account. for the redshift istribution of the SCUBA ealaxies., A model in which the single-temperature dust clouds discussed here are replaced by a temperature distribution will probably be required to account for the redshift distribution of the SCUBA galaxies.
 When the two spiral galaxies at 2oκ0.5 are replaced by EROs ab olI (Smail ct 11999). the agreement. between the 35-Ix. hierarchical. prediction and the observed redshift clistribution is rather satisfactory.," When the two spiral galaxies at $z < 0.5$ are replaced by EROs at $z>1$ (Smail et 1999), the agreement between the 35-K hierarchical prediction and the observed redshift distribution is rather satisfactory."
 In all the hierarchical mocels. despite strong negative evolution of the mass-to-light ratio of mergers with increasing redshift. most of the detected. galaxies are expected to [ie at. redshifts less than 5. and so will be accessible to. multi-waveband: study using S-m class telescopes.," In all the hierarchical models, despite strong negative evolution of the mass-to-light ratio of mergers with increasing redshift, most of the detected galaxies are expected to lie at redshifts less than 5, and so will be accessible to multi-waveband study using 8-m class telescopes."
 When final reliable identifications and. redshifts for submiüllimetre-selected galaxies are available. this information will be crucial for refining the hierarchical moclel.," When final reliable identifications and redshifts for submillimetre-selected galaxies are available, this information will be crucial for refining the hierarchical model."
 The discussion. has so far centred. on the properties of merging galaxies as observed through their dust emission in the mid-infrared. far-infrared ancl millimetre/submillimetre wavebands.," The discussion has so far centred on the properties of merging galaxies as observed through their dust emission in the mid-infrared, far-infrared and millimetre/submillimetre wavebands."
 Here. we assume the same forms of evolution of both the merger cllicieney parameter. and the activity parameter (fo)+ that were required to account. for the data in the far-infrared) ancl submillimetre wavebancls in the previous section. but make predictions in the near-infrared. optical ancl ultraviolet wavebands.," Here we assume the same forms of evolution of both the merger efficiency parameter $x$ and the activity parameter $(F\sigma)^{-1}$ that were required to account for the data in the far-infrared and submillimetre wavebands in the previous section, but make predictions in the near-infrared, optical and ultraviolet wavebands."
 In. particular. we investigate the 35-Ix. model. in which the recdshift distribution of SCUBA galaxies is in best agreement. with observations.," In particular, we investigate the 35-K model, in which the redshift distribution of SCUBA galaxies is in best agreement with observations."
 Subject to the uncertain fraction. of the. Iuminosity of these galaxies that is assumed to be powered. by star formation activity. we predict the integrated: background radiation intensity from the near-infrared {ο ionizing ultraviolet wavebands. and the counts of galaxies with SEDs that are dominated by evolved stars in the near-infrarecl A- band. and by voung stars in the optical B-band.," Subject to the uncertain fraction of the luminosity of these galaxies that is assumed to be powered by star formation activity, we predict the integrated background radiation intensity from the near-infrared to ionizing ultraviolet wavebands, and the counts of galaxies with SEDs that are dominated by evolved stars in the near-infrared $K$ -band, and by young stars in the optical $B$ -band."
" By requiring that the A- and. D-band counts are reproduced. accurately. we estimate both the fraction of all energy. released. in mergers that is reprocessed by dust sf and the normalization of the activity parameter at =0. (Lo),tin the optical waveband."," By requiring that the $K$ - and $B$ -band counts are reproduced accurately, we estimate both the fraction of all energy released in mergers that is reprocessed by dust $A$ and the normalization of the activity parameter at $z=0$, $(F\sigma)_0^{-1}$ in the optical waveband."
 For a discussion of the evolution of faint galaxies and their stellar populations see Ellis (1997)., For a discussion of the evolution of faint galaxies and their stellar populations see Ellis (1997).
 The counts of galaxies in the A-band at a [Bux density Si can be predicted by assuming the forms of the merger ellicieney. ο). as listed. in 11. an SED typical of evolved: stars [5 (Charlot. Worthey Bressan 1996). the PressSchechter mass function (equation 1) and the mass-to-light ratio of evolved stellar populations fait.," The counts of galaxies in the $K$ -band at a flux density $S_K$ can be predicted by assuming the forms of the merger efficiency $x(z)$, as listed in 1, an SED typical of evolved stars $f^K_\nu$ (Charlot, Worthey Bressan 1996), the Press–Schechter mass function (equation 1) and the mass-to-light ratio of evolved stellar populations $R_{\rm ML}$."
 “Phe SED was calculated. using a 9.25-Gvr old BruzualCharlot instantaneous burst model with a Salpeter ALP., The SED was calculated using a 9.25-Gyr old Bruzual–Charlot instantaneous burst model with a Salpeter IMF.
 Upper anl lower mass limits of 0.1 and MAL. were assumed for the IME., Upper and lower mass limits of 0.1 and $_\odot$ were assumed for the IMF.
 Note that the form. of the evolved stellar spectrum derived is almost independent. of the exact. values of the upper and lower mass limits assumed., Note that the form of the evolved stellar spectrum derived is almost independent of the exact values of the upper and lower mass limits assumed.
 Ehe A-band count with ὃν ensuring that the predicted. counts match the observed. A-band. counts. a suitable form. of the mass-to-light ratio yi is determined as a function of redshift.," The $K$ -band count with By ensuring that the predicted counts match the observed $K$ -band counts, a suitable form of the mass-to-light ratio $R_{\rm ML}$ is determined as a function of redshift."
 Phe mass in this ratio is defined as the mass of the dark matter haloes of galaxies. taken from the PressSchechter function (equation 1). and the luminosity is the bolometric Luminosity of the evolved stellar population in the galaxies.," The mass in this ratio is defined as the mass of the dark matter haloes of galaxies, taken from the Press–Schechter function (equation 1), and the luminosity is the bolometric luminosity of the evolved stellar population in the galaxies."
 In order to reproduce the observed. D-band counts. the counts derived. for the evolved. and merging. components ave added together. as shown in H41(b).," In order to reproduce the observed $B$ -band counts, the counts derived for the evolved and merging components are added together, as shown in 11(b)."
 Phe redshift dependence of the mass-to-light ratio is the same as that of the luminosity density of evolved stars. which depends on the SER at all earlier epochs.," The redshift dependence of the mass-to-light ratio is the same as that of the luminosity density of evolved stars, which depends on the SFR at all earlier epochs."
 A factor of 1|z is included in the denominator to mimic the effects of passive stellar evolution (Longair 1998)., A factor of $1+z$ is included in the denominator to mimic the effects of passive stellar evolution (Longair 1998).
 At z0 a form of the mass-to-light. ratio.," At $z=0$ a form of the mass-to-light ratio,"
formation rates and high-metallicities (Fig.4)).,formation rates and high-metallicities \ref{fig:sfrZrelation}) ).
" The first two phases of the relation between star formation and dilution are associated with before, during, and after the first pericentre passage and then close to final coalescence of the merging galaxies, the last phase in the metallicity evolution is always associated with well after their first pericentre passage and is near or after coalescence."," The first two phases of the relation between star formation and dilution are associated with before, during, and after the first pericentre passage and then close to final coalescence of the merging galaxies, the last phase in the metallicity evolution is always associated with well after their first pericentre passage and is near or after coalescence."
 This final state of the interaction is true whether one considers mergers or simply flybys., This final state of the interaction is true whether one considers mergers or simply flybys.
" While the correlation between SFR (at the peak of intensity) and metallicity dilution is quite remarkable, we expect a larger scatter when plotting the dilution as a function of the physical distance of galaxies in the pairs, especially for those at the smallest separations."," While the correlation between SFR (at the peak of intensity) and metallicity dilution is quite remarkable, we expect a larger scatter when plotting the dilution as a function of the physical distance of galaxies in the pairs, especially for those at the smallest separations."
" This is because the interacting pair can be in a number of different evolutionary states at any given distance, except for perhaps the smallest separations."," This is because the interacting pair can be in a number of different evolutionary states at any given distance, except for perhaps the smallest separations."
" For example, a galaxy at moderate separations of several 10s of kpc could be in an initial encounter and before first pericentre passage when the circumnuclear metallicity is largely unchanged."," For example, a galaxy at moderate separations of several 10s of kpc could be in an initial encounter and before first pericentre passage when the circumnuclear metallicity is largely unchanged."
 Or it could be just after the first pericentre passage when its circumnuclear gas could be highly diluted (Fig.5))., Or it could be just after the first pericentre passage when its circumnuclear gas could be highly diluted \ref{fig:sepZrelation}) ).
" This causes a range of possible metallicity dilutions, from Z/Ziso=1 to z/ziso~0.5 (Fig.5,, right panel), with an average value of z/Zis9=0.75."," This causes a range of possible metallicity dilutions, from $z/z_{iso}=1$ to $z/z_{iso}\sim0.5$ \ref{fig:sepZrelation}, right panel), with an average value of $z/z_{iso}=0.75$."
" At the smallest separations, say less than 10-20 kpc, but still distinguishable systems, galaxies are likely to be either near the pericentre passage or near final coalescence."," At the smallest separations, say less than 10-20 kpc, but still distinguishable systems, galaxies are likely to be either near the pericentre passage or near final coalescence."
" This leads to both a strong dilution and, on average, lower circumnuclear metallicity (see Fig.5))."," This leads to both a strong dilution and, on average, lower circumnuclear metallicity (see \ref{fig:sepZrelation}) )."
" At very small separations, less than about 4 kpc, galaxies can be near first pericentre passage or near coalescence which in either case have relatively large dilution and relatively low-metallicity, or can have coalesced and be strong enriched by supernovae ejecta and have relatively high-metallicity."," At very small separations, less than about 4 kpc, galaxies can be near first pericentre passage or near coalescence which in either case have relatively large dilution and relatively low-metallicity, or can have coalesced and be strong enriched by supernovae ejecta and have relatively high-metallicity."
 We show this in Fig., We show this in Fig.
" 5 where in the smallest separation bin, if we include galaxy pairs with projected separations of less than 4 kpc we find a net average increase in the metallicity (i.e., z >1, with an average value of z/ziso~ 1.2) while if we exclude such systems, we find a net decrease in the circumnuclear metallicity (z/z;;9~ 0.7)."," \ref{fig:sepZrelation} where in the smallest separation bin, if we include galaxy pairs with projected separations of less than 4 kpc we find a net average increase in the metallicity (i.e., $z>$ 1, with an average value of $z/z_{iso} \sim 1.2$ ) while if we exclude such systems, we find a net decrease in the circumnuclear metallicity $z/z_{iso} \sim 0.7$ )."
 This indicates the predominance of the metal enriched mergers at small projected separations., This indicates the predominance of the metal enriched mergers at small projected separations.
 Kewleyetal.(2006) have shown that galaxy pairs with small projected separations (4kpch7!«€s<20 kpch-!) and strong bursts of star formation have metallicities lower than the comparable field galaxies., \citet{kewley06} have shown that galaxy pairs with small projected separations $4 \;\mathrm{kpc \; h^{-1}} \le s\le 20 \;\mathrm {kpc \; h^{-1}}$ ) and strong bursts of star formation have metallicities lower than the comparable field galaxies.
" Note that in this analysis we are considering the physical (3D) distance between the two galaxies, and not their projected separation."," Note that in this analysis we are considering the physical (3D) distance between the two galaxies, and not their projected separation."
 Considering the effects of projection would likely contribute to increasing the scatter in the relation between metallicity and separation., Considering the effects of projection would likely contribute to increasing the scatter in the relation between metallicity and separation.
 Our simulations do not track the evolution of individual elements directly in the code., Our simulations do not track the evolution of individual elements directly in the code.
" However, to track their evolution in an approximate way, we can use some simple arguments to give an estimate of the amount of iron and « elements, like oxygen, released in the ISM by type I and II SNe."," However, to track their evolution in an approximate way, we can use some simple arguments to give an estimate of the amount of iron and $\alpha$ elements, like oxygen, released in the ISM by type I and II SNe."
" At a given time, t, during the evolution of a galaxy, the rate of oxygen mass released in the ISM by SNe II can be estimated by:"," At a given time, $t$, during the evolution of a galaxy, the rate of oxygen mass released in the ISM by SNe II can be estimated by:"
values inCDMS=.. ios is gathered directly. auc Ay is calculated. from the listed. Εως values: Ey=Blow0.0331xr. where v is the line (requeney in GHz.,"values in, $\mu^2S$ is gathered directly, and $E_{\rm u}$ is calculated from the listed $E_{\rm low}$ values: $E_{\rm u}=E_{\rm low}+0.0334\times\nu$, where $\nu$ is the line frequency in GHz."
 Then au upper limit ou tle NoH temperature can be calculated [rom Usine the integrated inteusity of the J=3-2 line and the 2o upper limit on the J=1-3 line. we obtain a 20 upper limit ou the NoH rotational temperature of <26 Ik. This is consistent with the predictions that NeH should become abuncaut below the CO lreeze-out temperature of 20 Ix. but does not provide any stronger coustraints.," Then an upper limit on the $_2$ $^+$ temperature can be calculated from Using the integrated intensity of the J=3–2 line and the $\sigma$ upper limit on the J=4–3 line, we obtain a $\sigma$ upper limit on the $_2$ $^+$ rotational temperature of $<$ 26 K. This is consistent with the predictions that $_2$ $^+$ should become abundant below the CO freeze-out temperature of 20 K, but does not provide any stronger constraints."
 [tis also consistent with the NoH J=1-0 line reported by Dutreyetal.(2007)., It is also consistent with the $_2$ $^+$ J=1–0 line reported by \citet{Dutrey07}.
. We estimate the line [lux [roin their Figure 1 to be «0.1 Jy kins +., We estimate the line flux from their Figure 1 to be $<$ 0.1 Jy km $^{-1}$ .
 Using this limit and the inteeratecl intensity of the J—3-2 liue results in a lower limit on the NoH rotational temperature of 11 Ix. More seusitive observations of the ΝΟΗ lines would be valuable to constrain this excitation temperature fijer., Using this limit and the integrated intensity of the J=3–2 line results in a lower limit on the $_2$ $^+$ rotational temperature of 11 K. More sensitive observations of the $_2$ $^+$ lines would be valuable to constrain this excitation temperature further.
 In the »euce of more detailed informatiou. we calculate LTE colt densities of midplane lous for a range of temperatures. from 10 to 20Ix. The 20 Ix boundary is set by laboratory experiments ou CO aud No freeze-out — above this temperature a majority of CO and Ne are maintained in the gas-phase quickly destroying any formed Ην 2006).," In the absence of more detailed information, we calculate LTE column densities of midplane ions for a range of temperatures, from 10 to 20K. The 20 K boundary is set by laboratory experiments on CO and $_2$ freeze-out – above this temperature a majority of CO and $_2$ are maintained in the gas-phase quickly destroying any formed $_3^+$ \citep{Bisschop06}."
. The CO suowline was recently reported to correspond to a freeze-out temperature of 19 Ix. consistent. with predictions (Qietal.2011," The CO snowline was recently reported to correspond to a freeze-out temperature of 19 K, consistent with predictions \citep{Qi11}."
 Iu general. columau deusities are calculated most accurately by coustructiug a self-consisteut cheimmical-physical disk mocel and applying a radiative transfer code to predict line enuission profiles hat cau be compared directly to observations.," In general, column densities are calculated most accurately by constructing a self-consistent chemical-physical disk model and applying a radiative transfer code to predict line emission profiles that can be compared directly to observations."
 Without a detection of the HeD line. however. it is difficult to coustrain where the HoD — emission arises in the disk.," Without a detection of the $_2$ $^+$ line, however, it is difficult to constrain where the $_2$ $^+$ emission arises in the disk."
 We therefore opt for a classical LTE approach (e.g.vanDishoecketal.2003).. calculated lor the range of reasonable disk midplane emperatures (see 'e[secanukdplaue;emp))/lopullimilsonthediskaveraged/N(o HaD Ἰ.," We therefore opt for a classical LTE approach \citep[e.g.][]{vanDishoeck03}, calculated for the range of reasonable disk midplane temperatures (see \\ref{sec:midplane_temp}) ) to put limits on the disk averaged$o$ $_2$ $^+$ )."
 A significant complication for the LTE calculation is that orfhe (ο) and pare (p) ΗΕ can yellave as separate species at low temperatures (Floweretal.20014:Sipila2010) resulting in a nou-thermal partitiouiug between the o and p states.," A significant complication for the LTE calculation is that $ortho$ $o$ ) and $para$ $p$ ) $_2$ $^+$ can behave as separate species at low temperatures \citep{Flower04,Sipila10} resulting in a non-thermal partitioning between the $o$ and $p$ states."
 To caleulate the partition fuuction lor o-H»3D . we use a two-level approximation. which assumes no trausitious between the o and p H3D . and that transitious between the ground state aud first excited state clominate within each," To calculate the partition function for$o$ $_2$ $^+$ , we use a two-level approximation, which assumes no transitions between the $o$ and $p$ $_2$ $^+$ , and that transitions between the ground state and first excited state dominate within each"
"at D,~6.5kpe. so that the angular Einstein radius becomes fp-320µας. the Einstein radius becomes re=DL2.0AU and its projection tothe source distance becomes rp=(DifDi)rg~500 R..","at $D_\rmn{L} \sim 6.5~\mbox{kpc}$, so that the angular Einstein radius becomes $\theta_\rmn{E} \sim\,320~\mu\mbox{as}$, the Einstein radius becomes $r_\rmn{E} = D_\rmn{L}\,\theta_\rmn{E} \sim 2.0~\mbox{AU}$ and its projection tothe source distance becomes $r_\rmn{E}' = (D_\rmn{S}/D_\rmn{L})\,r_\rmn{E} \sim 500~R_{\odot}$ ."
" With 4,~25fH. for a KS giant. a source size parameter p,=0.05 is therefore adopted."," With $R_\star \sim 25~R_{\odot}$ for a K5 giant, a source size parameter $\rho_\star = 0.05$ is therefore adopted."
 In order to make an optimal case for observing. a rather low proper motion pooGkmsà corresponding to a lens velocity of e~55kms relative to the source has been chosen. which vields an event time- fe=55d.," In order to make an optimal case for observing, a rather low proper motion $\mu \sim\,6~\mbox{km}\,\mbox{s}^{-1}$, corresponding to a lens velocity of $v \sim 55~\mbox{km}\,\mbox{s}^{-1}$ relative to the source has been chosen, which yields an event time-scale $t_\rmn{E} = 55~\mbox{d}$."
 The impact parameter is chosen as vg=0.015. corresponding to a peak magnification of a point source of οί nu ," The impact parameter is chosen as $u_0 = 0.015$, corresponding to a peak magnification of a point source of $A_0 \sim 70$ ."
Configuration. II has been chosen to be similar to the parameters of the observed event MACHO 1995-BLG-30 (?).. for which the source is an even larger star. namely an M4 giant.," Configuration II has been chosen to be similar to the parameters of the observed event MACHO 1995-BLG-30 \citep{MACHO:9530}, for which the source is an even larger star, namely an M4 giant."
" According to the obtained model parameters. let us adopt 0.05. p,= 0.075. and the event time-scale fe=35d."," According to the obtained model parameters, let us adopt $u_0 = 0.05$ , $\rho_\star = 0.075$ , and the event time-scale $t_\rmn{E} = 35~\mbox{d}$."
" For kpe and D,~6.5 kpe. the appropriate stellar radius of 2,~A. is obtained for AL~0.7AL... so that rp~SOO Re."," For $D_\rmn{S} \sim\,8.5~\mbox{kpc}$ and $D_\rmn{L} \sim\,6.5~\mbox{kpc}$ , the appropriate stellar radius of $R_\star \sim\,60~R_{\sun}$ is obtained for $M \sim\,0.7~M_{\sun}$, so that $r_\rmn{E}' \sim\,800~R_{\sun}$ ."
 These choices yield rg~2.4)AU and 6p.~460jas. so tha the proper motion becomes yp~15kms! and the relative lens velocity ise~140kms," These choices yield $r_\rmn{E} \sim\,2.9~\mbox{AU}$ and $\theta_\rmn{E} \sim\,460~\mu\mbox{as}$, so that the proper motion becomes $\mu \sim\,15~\mbox{km}\,\mbox{s}^{-1}$ and the relative lens velocity is $v \sim\,140~\mbox{km}\,\mbox{s}^{-1}$."
 With Ady= |O2and\f=|34 for an M4 giant. an extinction of τς=0.85 yields the baseline magnitude fi...=12.3.," With $M_V = +0.2$ and $V-I = +3.4$ for an M4 giant, an extinction of $A_I = 0.85$ yields the baseline magnitude $I_\rmn{base} = 12.3$."
 Binary lenses are likely to cause an asymmetry to the ligh curve. Which however can be arbitrarily small and can even vanish or some configurations.," Binary lenses are likely to cause an asymmetry to the light curve, which however can be arbitrarily small and can even vanish for some configurations."
 In order to study the maximal impact of binarity on the measurement of limb darkening. configurations have been chosen that preserve the symmetry.," In order to study the maximal impact of binarity on the measurement of limb darkening, configurations have been chosen that preserve the symmetry."
 Therefore. let us assume hat both lens objects have the same mass and consider a source rajectory parallel to the line connecting their angular positions.," Therefore, let us assume that both lens objects have the same mass and consider a source trajectory parallel to the line connecting their angular positions."
 In general. the light received from the source is blended with additional light from other unresolved sources (that are not affected by microlensing) or from the lens star. which is quantified by the blend ratio y=fpffs. where P5 denotes the source flux and £1 denotes the background (blend) flux.," In general, the light received from the source is blended with additional light from other unresolved sources (that are not affected by microlensing) or from the lens star, which is quantified by the blend ratio $g = F_\rmn{B}/F_\rmn{S}$, where $F_\rmn{S}$ denotes the source flux and $F_\rmn{B}$ denotes the background (blend) flux."
 Since the choice of equal lens masses implies that the lens objects are M dwarfs of mass Al/2— or M/2~0.35M. . their contribution to the total lightalls well belowthe systematic error barseven at the observed /-," Since the choice of equal lens masses implies that the lens objects are M dwarfs of mass $M/2 \sim 0.18~M_\odot$ or $M/2 \sim 0.35~M_\odot$ , their contribution to the total lightfalls well belowthe systematic error barseven at the observed $I$ -baseline."
 Therefore. blending isneglected with the choice g= 0.," Therefore, blending isneglected with the choice $g = 0$ ."
 Finally. a limb-darkening coefficient Py=0.5 has been adopted or both configurations.," Finally, a limb-darkening coefficient $\Gamma_I = 0.5$ has been adopted for both configurations."
distance @ from the centre of foreground galaxies.,distance $\theta$ from the centre of foreground galaxies.
 Although in the extreme locality there will be a larger flexion than usual. the mean signal averaged around the galaxy will still drop rapidly with 8.," Although in the extreme locality there will be a larger flexion than usual, the mean signal averaged around the galaxy will still drop rapidly with $\theta$."
 However. this will not be true for the flexion in the annulus.," However, this will not be true for the flexion in the annulus."
 This will respond to any density fluctuations within the annulus., This will respond to any density fluctuations within the annulus.
 Therefore. in any annuli with non-negligible substructure. even if the mean flexion is small (as the mean gradient of density is small). the flexion variance will remain comparatively large.," Therefore, in any annuli with non-negligible substructure, even if the mean flexion is small (as the mean gradient of density is small), the flexion variance will remain comparatively large."
 It is his behaviour that we will use to constrain substructure on galactic scales., It is this behaviour that we will use to constrain substructure on galactic scales.
 At this point one may ask whether flexion variance is the best tool for our task: wouldn't flexion correlation functions in annuli provide more information?, At this point one may ask whether flexion variance is the best tool for our task; wouldn't flexion correlation functions in annuli provide more information?
 The question could be informed by experience in cosmic shear studies. where shear correlation 'uncetions provide more finesse than shear variance in cells.," The question could be informed by experience in cosmic shear studies, where shear correlation functions provide more finesse than shear variance in cells."
 However. in our present case a correlation function does no seem to be helpful.," However, in our present case a correlation function does not seem to be helpful."
 Since the correlation function in question would be in annuli around a foreground galaxy. it constitutes a form of galaxy-galaxy-galaxv lensing (2): it i8 a three-point statistic.," Since the correlation function in question would be in annuli around a foreground galaxy, it constitutes a form of galaxy-galaxy-galaxy lensing \citep{2005A&A...432..783S}; it is a three-point statistic."
 In order to use this to measure substructures. it is necessary for two background sources to be close to the same substructure as wel as to the foreground galaxy: this rarely happens. leading to low signal-to-noise.," In order to use this to measure substructures, it is necessary for two background sources to be close to the same substructure as well as to the foreground galaxy; this rarely happens, leading to low signal-to-noise."
 On the other hand. the flexion variance in annuli only involves two points. a foreground-background pair. with much greater signal-to-noise as we shall see.," On the other hand, the flexion variance in annuli only involves two points, a foreground-background pair, with much greater signal-to-noise as we shall see."
 In this section we describe N-body simulations which we will use to demonstrate the utility of flexion variance as a probe of substructure., In this section we describe N-body simulations which we will use to demonstrate the utility of flexion variance as a probe of substructure.
 We use the cosmological ACDM simulation already presented in ?.., We use the cosmological $\Lambda$ CDM simulation already presented in \citet{2005MNRAS.364..367D}.
 The simulation was run using PRUGRAY (?).. with cosmological parameters: (Qu.Q4.σε.fh)Ξ(0.268.0.732.0.7.0.71). and a box of size Ling=90 Mpc. with 300° particles.," The simulation was run using PkdGRAV \citep{2001PhDT........21S}, with cosmological parameters: $(\Omega_{\rm m}, \Omega_{\Lambda}, \sigma_8, h) = (0.268, 0.732, 0.7, 0.71)$, and a box of size $L_{\rm box}=90$ Mpc, with $300^3$ particles."
 The initial conditions were generated with GRAFIC? (?).., The initial conditions were generated with GRAFIC2 \citep{2001ApJS..137....1B}.
" From the simulation volume. we extracted four Milky Way sized halos at a mass resolution of m,=5.7«LOYAL.: their virial masses are 2.1.1.5.1.2.1.3]10/7AL."," From the simulation volume, we extracted four Milky Way sized halos at a mass resolution of $m_p=5.7\times 10^5 M_\odot$; their virial masses are $[2.1, 1.5, 1.2, 1.3]\times10^{12}M_\odot$."
 While we are therefore very limited in our number of lenses (to three projections of each of four high-resolution galaxies). we will find that this is sufficient to give the initial indicative results required by this paper.," While we are therefore very limited in our number of lenses (to three projections of each of four high-resolution galaxies), we will find that this is sufficient to give the initial indicative results required by this paper."
 As in ?.. the subhalos inside each “Milky Way” at redshift 2=0 were identitied using theAHF algorithm (2)..," As in \citet{2008MNRAS.389.1041R}, , the subhalos inside each `Milky Way' at redshift $z=0$ were identified using the algorithm \citep{2004MNRAS.351..399G}."
 We considered all subhalos with 50 particles and assigned particles to the smallest structure they appear in so that each particle was counted only once., We considered all subhalos with $>50$ particles and assigned particles to the smallest structure they appear in so that each particle was counted only once.
 In some cases we will remove substructure: this is achieved by subtracting all particles not assigned to the main halo., In some cases we will remove substructure; this is achieved by subtracting all particles not assigned to the main halo.
 An example halo. with and without substructure. is shown in Figure I.," An example halo, with and without substructure, is shown in Figure \ref{fig:nbody}."
 The 3D numerical simulations discussed above represent the density field using discrete particles., The 3D numerical simulations discussed above represent the density field using discrete particles.
 We transform these into convergence maps by projecting the particles along particular spatial directions and placing the particles onto a 2D 1024x1024 grid. which we carry out using the IDL cloud-in-cell routine available as part of The IDL Astronomy User's.," We transform these into convergence maps by projecting the particles along particular spatial directions and placing the particles onto a 2D 1024x1024 grid, which we carry out using the IDL cloud-in-cell routine available as part of The IDL Astronomy User's."
Library... We produce three projection maps for each 3D halo (by projecting along the X. y or z axis).," We produce three projection maps for each 3D halo (by projecting along the x, y or z axis)."
 We investigated a number of techniques for filtering the mass maps., We investigated a number of techniques for filtering the mass maps.
 This is important because the finite number of simulation particles introduces shot noise into the 2D maps: this can compete with the substructure signal we are investigating., This is important because the finite number of simulation particles introduces shot noise into the 2D maps; this can compete with the substructure signal we are investigating.
 Here we show the results obtained using Multiscale Entropy Filtering (MEF) (2).., Here we show the results obtained using Multiscale Entropy Filtering (MEF) \citep{2006A&A...451.1139S}.
 This provides superior performance to a simple Guassian filter. as it reduces the shot noise while preserving the density fluctuations. as we shall demonstrate.," This provides superior performance to a simple Guassian filter, as it reduces the shot noise while preserving the density fluctuations, as we shall demonstrate."
 For this purpose weuse the routines, For this purpose weuse the routines
"corresponding to ESR size are general, Figure 12a shows correlation measures and Figure 12b displays a diagram indicating the ESR separations that relate to the lags for the cross- and auto-correlations in Figure 12a..","corresponding to ESR size are general, Figure \ref{Fig:SizeCorrelation} shows correlation measures and Figure \ref{Fig:CorrExplanation} displays a diagram indicating the ESR separations that relate to the lags for the cross- and auto-correlations in Figure \ref{Fig:SizeCorrelation}."
" The left panel in Figure 12a shows the HP filter light curve auto-correlation and the right panel shows the cross-correlation of HP filter light curves with VP filter light curves; all with ϕ=0,0. “", The left panel in Figure \ref{Fig:SizeCorrelation} shows the HP filter light curve auto-correlation and the right panel shows the cross-correlation of HP filter light curves with VP filter light curves; all with $\phi = \theta_f = 0$. “
"Auto” indicates auto-correlation, and “X” indicates cross-correlation.","Auto” indicates auto-correlation, and “X” indicates cross-correlation."
" Overlaying each plot is either a dashed line that corresponds to the time the ESR takes to travel its diameter (2rz+wz) or a dotted line that relates to the radius (rg+we/2), both defined graphically in Figure 12b.."," Overlaying each plot is either a dashed line that corresponds to the time the ESR takes to travel its diameter $2r_E+w_E$ ) or a dotted line that relates to the radius $r_E+w_E/2$ ), both defined graphically in Figure \ref{Fig:CorrExplanation}."
" There is good agreement between both the lines and the correlation peaks; if one HP wing crosses a caustic, it is highly probable the second will too once the ESR travels the length of its diameter."," There is good agreement between both the lines and the correlation peaks; if one HP wing crosses a caustic, it is highly probable the second will too once the ESR travels the length of its diameter."
" Likewise, the perpendicular VP wings will also cross the caustic once the ESR travels the length of its diameter."," Likewise, the perpendicular VP wings will also cross the caustic once the ESR travels the length of its diameter."
" The curvature of the wings at larger inner radii and widths means that the correlation length is lower than the ESR diameter, shown by the weakened correlation between dashed lines and peaks in the bottom three panels of Figure 12a.."," The curvature of the wings at larger inner radii and widths means that the correlation length is lower than the ESR diameter, shown by the weakened correlation between dashed lines and peaks in the bottom three panels of Figure \ref{Fig:SizeCorrelation}."
 More detailed calculations relating to the spatial mean of wing intensity will improve the agreement., More detailed calculations relating to the spatial mean of wing intensity will improve the agreement.
" However, it is evident that the size of an ESR according to our model is discernible from the relevant correlation curves."," However, it is evident that the size of an ESR according to our model is discernible from the relevant correlation curves."
 In this paper we have presented a numerical study of the influence of gravitational microlensing upon the electron scattering region thought to exist in the inner regions of quasars., In this paper we have presented a numerical study of the influence of gravitational microlensing upon the electron scattering region thought to exist in the inner regions of quasars.
" Focusing upon the quadruply imaged quasar, Q2237+0305, we considered a fiducial model with a 0.08 pc inner radius and a 0.32 pc width consistent with recent observations (Taniguchi&Anabuki1999)."," Focusing upon the quadruply imaged quasar, Q2237+0305, we considered a fiducial model with a $0.08$ pc inner radius and a $0.32$ pc width consistent with recent observations \citep{Taniguchi:99}."
" As seen in Figure 5,, the magnification of the two polarised annulus components can also be well correlated if the relative angle between filter and caustics is an odd multiple of 7/4."," As seen in Figure \ref{Fig:RotationDistributionsB}, the magnification of the two polarised annulus components can also be well correlated if the relative angle between filter and caustics is an odd multiple of $\pi/4$."
" It was also discovered that, by orienting a polarising filter and defining an axis of polarisation as parallel to the magnification map’s shear axis, larger ESR dimensions allow for an increased level of events where one polarisation is magnified more than the other."," It was also discovered that, by orienting a polarising filter and defining an axis of polarisation as parallel to the magnification map's shear axis, larger ESR dimensions allow for an increased level of events where one polarisation is magnified more than the other."
" By considering the cross- and auto-correlation of the resulting light curves, it is seen that periodicities in the microlensing magnification maps are imprinted on the light curve, providing further clues to the ESR scattering geometry, with the size of the ESR being discernible from the caustic-crossing auto-correlation curve of the polarised “wings” that are parallel to the caustics, or the cross-correlation between the magnification of both polarisations."," By considering the cross- and auto-correlation of the resulting light curves, it is seen that periodicities in the microlensing magnification maps are imprinted on the light curve, providing further clues to the ESR scattering geometry, with the size of the ESR being discernible from the caustic-crossing auto-correlation curve of the polarised “wings” that are parallel to the caustics, or the cross-correlation between the magnification of both polarisations."
" Polarimetric observations during a known microlensing event would provide more information than constant polarimetric monitoring, which is observationally more expensive."," Polarimetric observations during a known microlensing event would provide more information than constant polarimetric monitoring, which is observationally more expensive."
" Q2237+0305 is regularly monitored photometrically (Udalskietal.2006,eg.theOGLEproject) and a photometric microlensing event could “trigger” an intense polarimetric study [such overrides have been proposed for Q2237+0305 in the past (e.g.Webster 2004)]], which would directly probe the structure of the ESR."," Q2237+0305 is regularly monitored photometrically \citep[e.g. the OGLE project]{Udalski:2006} and a photometric microlensing event could “trigger” an intense polarimetric study [such overrides have been proposed for Q2237+0305 in the past \citep[e.g.][]{Webster:2004}] ], which would directly probe the structure of the ESR."
" It must be noted that, to obtain long term variability statistics, photometric and subsequent polarimetric monitoring will be required on decade to century time- even to observe the characteristic correlation lags indicated in Figure 12a for the small-scale fiducial model."," It must be noted that, to obtain long term variability statistics, photometric and subsequent polarimetric monitoring will be required on decade to century time-scales, even to observe the characteristic correlation lags indicated in Figure \ref{Fig:SizeCorrelation} for the small-scale fiducial model."
" Nevertheless, a monitoring campaign on that order of"," Nevertheless, a monitoring campaign on that order of"
twilight sky images requires careful timine: if they are taken too carly the detector saturates. taken too lae then there is iusufficient sienal to form a high signal-to-noise flat.,"twilight sky images requires careful timing; if they are taken too early the detector saturates, taken too late then there is insufficient signal to form a high signal-to-noise flat."
 For our pixel eteudue (1.G10* cniD str) the optinnuuaH time. for+ taking+ Abend images ds between suuset or suse and G-deeree twilight., For our pixel etendue $1.6 \times 10^{-7}$ $^2$ str) the optimum time for taking $K_s$ -band images is between sunset or sunrise and 6-degree twilight.
 A field at high ealactic latitucle is suitable to avoid stars., A field at high galactic latitude is suitable to avoid stars.
" The telescope is jogeed between c""ach exposure so lat stars occupy differeut pixels i1 subsequent picures.", The telescope is jogged between each exposure so that stars occupy different pixels in subsequent pictures.
 Any pixels contziuiug obvkons stars are oejvon ZCLO weight and the frames ire conibined into a weighted nea where outvine pixels (fain stars or cosniic rays) are rejectcc Lit their deviation from the mean exceeds three sigma., Any pixels containing obvious stars are given zero weight and the frames are combined into a weighted mean where outlying pixels (faint stars or cosmic rays) are rejected if their deviation from the mean exceeds three sigma.
 Iu a flat normalized so that he median pixe value is 1.0 the vest fit Cassian las au rus of 0.ML (Fig. 12))., In a flat normalized so that the median pixel value is 1.0 the best fit Gaussian has an rms of 0.04 (Fig. \ref{flatfieldhistogram}) ).
 Tho‘co percent of the jxels fall in a ail with low res)onse outside of his cistribution. aud the array cau be considered operable.," Three percent of the pixels fall in a tail with low response outside of this distribution, and the array can be considered operable."
 Under the assumption that the variance iu the camera signal is the suu of a coustaut reac noise and Poisson fiuctuations it is straightforware to measure the read noise and gain., Under the assumption that the variance in the camera signal is the sum of a constant read noise and Poisson fluctuations it is straightforward to measure the read noise and gain.
 The data required for this measurement are a sequence of exposures of increase inteeration time while the camera views a source of spatially uniforiii ihunination., The data required for this measurement are a sequence of exposures of increasing integration time while the camera views a source of spatially uniform illumination.
 A pair of exposures d8 acquires at cach inteeration time., A pair of exposures is acquired at each integration time.
 The mean signal leve was incastured by subtracting a bias frame aux the variance computed by differencing the two framessince the uniformity of flat field is very hieh the frames were not flat fielded., The mean signal level was measured by subtracting a bias frame and the variance computed by differencing the two frames—since the uniformity of flat field is very high the frames were not flat fielded.
 Α straight-line. least squares fit to the linear mean-variance relation viclds the readout noise as the intercept and the gain as the slope.," A straight-line, least squares fit to the linear mean-variance relation yields the readout noise as the intercept and the gain as the slope."
 The results of this analysis vield 30 per data number (DN). aud à rend noise of TO e rms.," The results of this analysis yield 30 $^-$ per data number (DN), and a read noise of 70 $^-$ rms."
 The majority of this is detector noise. since if one grounds he iuput to the data acquisition svstem the variation is on the order of 30 rmm.," The majority of this is detector noise, since if one grounds the input to the data acquisition system the variation is on the order of 30 $^-$ rms."
 A science gerade PICNIC array ds expected to have a read noise of about 10 nus., A science grade PICNIC array is expected to have a read noise of about 10 $^-$ rms.
 Despite this elevated detector noise. the data are background lanited iu al but the shortest exposures.," Despite this elevated detector noise, the data are background limited in all but the shortest exposures."
 The array saturates at 22.000 DN (660.000 } aud shows < nonlinearity at up to of this level.," The array saturates at 22,000 DN (660,000 $^-$ ) and shows $<$ nonlinearity at up to of this level."
 The efficiency. of the camera is expected to be high., The efficiency of the camera is expected to be high.
 The detector quantuui efficiency is over most of the operating wavelength range., The detector quantum efficiency is over most of the operating wavelength range.
" The average filter trausmissions are (7). UL). and (K,) between the halfpower pots."," The average filter transmissions are $J$ ), $H$ ), and $K_s$ ) between the half-power points."
 The protected Al coating for the primary M4. las au infrared reflectivity of and the reflectivity of bare Al on the secondary. Mo. is probably similar.," The protected Al coating for the primary, $M_1$, has an infrared reflectivity of, and the reflectivity of bare Al on the secondary, $M_2$, is probably similar."
 The quartz Dewar window has a transmission of[4.. giving a predicted camera efficiencv of .," The quartz Dewar window has a transmission of, giving a predicted camera efficiency of ."
". The system efficicney, including the telescope and atinosphere. determined frou observation of standard stars is GI). (FT). aud (fy)."," The system efficiency, including the telescope and atmosphere, determined from observation of standard stars is $J$ ), $H$ ), and $K_s$ )."
" The camera zero points in magnitudes on a Veea based scale. defined as are listed iu Table L.. where DN 1 is the count rate in for a star of magnitude sg,The regionNGC 2021 (Orion D. W12) is a"," The camera zero points in magnitudes on a Vega based scale, defined as are listed in Table \ref{zeropoints}, where DN $^{-1}$ is the count rate in for a star of magnitude $m_{star}$ .The regionNGC 2024 (Orion B, W12) is a"
innvestigating the radial distribution of iron and other metals.,nvestigating the radial distribution of iron and other metals.
"flux ¢ in pixels within the grid cell, therefore A large statistical sample could be obtained by collecting magnetometer images at different times.","flux $\phi$ in pixels within the grid cell, therefore A large statistical sample could be obtained by collecting magnetometer images at different times."
" For each ἶσομ, this allows a computation of the probability distribution of net signed flux, P(|®ceu|)."," For each $l_{cell}$, this allows a computation of the probability distribution of net signed flux, $P(|\Phi_{cell}|)$."
 The model results indicate that this would yield a power law distribution of flux concentrations for a range of lcci., The model results indicate that this would yield a power law distribution of flux concentrations for a range of $l_{cell}$.
 The cutoffs at both small and large values of flux may depend on ει., The cutoffs at both small and large values of flux may depend on $l_{cell}$.
" If the cutoff at large scales increases with [..;;, data collapse methods could be used to rescale the flux according to the cell size, so that all the distributions would coincide on the rescaled plot."," If the cutoff at large scales increases with $l_{cell}$, data collapse methods could be used to rescale the flux according to the cell size, so that all the distributions would coincide on the rescaled plot."
" We expect a similar behavior to obtain by neglecting the sign of the flux, i.e. the probability distribution for net unsigned flux within a grid cell, P(S>..1;|ó|), may also be scale-free."," We expect a similar behavior to obtain by neglecting the sign of the flux, i.e. the probability distribution for net unsigned flux within a grid cell, $P(\sum_{cell}|\phi|)$, may also be scale-free."
 The magnetic concentrations themselves do not comprise a network., The magnetic concentrations themselves do not comprise a network.
" They must be joined by magnetic fields, or flux tubes, in order to do so."," They must be joined by magnetic fields, or flux tubes, in order to do so."
 In fact ? also measured statistical properties of the magnetic flux network in the quiet-Sun that allow comparisons with our model., In fact \citet{close} also measured statistical properties of the magnetic flux network in the quiet-Sun that allow comparisons with our model.
" They found that (1) for any concentration strength there are a wide range of possible connections; (2) concentrations show a preference towards connecting to nearby opposite polarity concentrations; (3) despite the vast number of possible connections, the bulk of the flux is often divided so that most of it goes to one opposite polarity concentration."," They found that (1) for any concentration strength there are a wide range of possible connections; (2) concentrations show a preference towards connecting to nearby opposite polarity concentrations; (3) despite the vast number of possible connections, the bulk of the flux is often divided so that most of it goes to one opposite polarity concentration."
 Similar results are found in numerical simulations of our model., Similar results are found in numerical simulations of our model.
" In fact, all of this behavior is explained by the existence of three different scale-free distributions."," In fact, all of this behavior is explained by the existence of three different scale-free distributions."
 The first is the distribution of concentration sizes (node degree) as discussed previously., The first is the distribution of concentration sizes (node degree) as discussed previously.
" T'he second and third are new statistical quantities we introduce to characterize networks: (a) the amount of flux connecting a pair of concentrations, or the strength of the link between a pair of nodes, and (b) the number of distinct concentrations (nodes) linked to a given one."," The second and third are new statistical quantities we introduce to characterize networks: (a) the amount of flux connecting a pair of concentrations, or the strength of the link between a pair of nodes, and (b) the number of distinct concentrations (nodes) linked to a given one."
" These are both found to be power laws in the model, with different indices."," These are both found to be power laws in the model, with different indices."
 We also discuss the distribution of the lengths of flux tubes and compare with observations of Close et al., We also discuss the distribution of the lengths of flux tubes and compare with observations of Close et al.
 This allows a calibration of the length unit in the model to a physical length on the photosphere., This allows a calibration of the length unit in the model to a physical length on the photosphere.
 'Two opposite polarity footpoints can be connected to each other by any number of loops., Two opposite polarity footpoints can be connected to each other by any number of loops.
 The number of loops connecting a pair of footpoints is defined as kpair., The number of loops connecting a pair of footpoints is defined as $k_{pair}$.
 This is the strength of the link between two nodes., This is the strength of the link between two nodes.
 Measuring this value over all footpoint pairs gives the distribution shown in Figure 7.., Measuring this value over all footpoint pairs gives the distribution shown in Figure \ref{kpair}. .
" Power law behavior occurs for both models, i.e. The critical index a is equal to 2.25+0.1 and 2+0.1 for models I and II respectively."," Power law behavior occurs for both models, i.e. The critical index $\alpha$ is equal to $2.25\pm0.1$ and $2\pm0.1$ for models I and II respectively."
" The index α2*y, as the number of loops connecting any pair offootpoints cannot exceed the number of loops"," The index $\alpha\geq\gamma$, as the number of loops connecting any pair offootpoints cannot exceed the number of loops"
By azuuuthal averaging to the maximum extent possible for a given sequence of observations. it is possible to determine the mean radial dependence of the emissiou intensity iutegrated over he line profile.,"By azimuthal averaging to the maximum extent possible for a given sequence of observations, it is possible to determine the mean radial dependence of the emission intensity integrated over the line profile."
 This was done ouly for exposures obtained under good photometric conditious., This was done only for exposures obtained under good photometric conditions.
 The iuteusitv of the scattered light at the radial distance + is written as μμ)., The intensity of the scattered light at the radial distance $r$ is written as $I_{scatt}(r)$.
 We scale the iutensitv to the measure stellar flux £. measured with the very wide slit., We scale the intensity to the measured stellar flux $I_*$ measured with the very wide slit.
 Z. was measured by integrating the ou-star spectrum flux over the whole spectral interval (512 pixels corresponding to 1À))., $I_*$ was measured by integrating the on-star spectrum flux over the whole spectral interval (512 pixels corresponding to ).
 This total flux was then divided by the wavelength range aud the integration time., This total flux was then divided by the wavelength range and the integration time.
 The resulting Z. is thus expressed iu counts per per second., The resulting $I_*$ is thus expressed in counts per per second.
 Note that Mauron et al. (, Note that Mauron et al. (
1981) used a similar definition iu averaging the flux over20À.. while Meuron (1990). Tloneveutt et al. (,"1984) used a similar definition in averaging the flux over, while Mauron (1990), Honeycutt et al. ("
1980). and Bernat et al. (,"1980), and Bernat et al. ("
1978) used the photospheric flux in ccentered onT6900A.,1978) used the photospheric flux in centered on.
". There are a number of sources of uncertainty affecting the determüunation of the Nhotospherie flux and in the με)ff. ratio. mostly secause Of possible loss of plotous intercepted by the slit jaws. and changes of transparency and seeiug coucditionus )etween exposures,"," There are a number of sources of uncertainty affecting the determination of the photospheric flux and in the $I_{scatt}(r)/I_*$ ratio, mostly because of possible loss of photons intercepted by the slit jaws, and changes of transparency and seeing conditions between exposures."
 The derived nonualized intensity of he emussion. {μετοι is shown iu Fie.," The derived normalized intensity of the emission, $I_{scatt}(r)/I_*$, is shown in Fig."
 5 for impact xuneters Rout to 50 arcsec from the star., \ref{fig6} for impact parameters $R$ out to 50 arcsec from the star.
 The mean xofile is consistent with a power law of slope -2.36 4 0.03 (least-square fit to the data)., The mean profile is consistent with a power law of slope -2.36 $\pm$ 0.03 (least-square fit to the data).
to) the CO emission width.,to) the CO emission width.
 If the widths of both diagnostics are dominated by Neplerian rotation. this suggests that the [Nell] emission arises from larger disk radii on average than the CO emission.," If the widths of both diagnostics are dominated by Keplerian rotation, this suggests that the [NeII] emission arises from larger disk radii on average than the CO emission."
 This is consistent. in general. with models of the ionization and thermal structure of T Tauri disks (e.g.. Glassgold et 22007: Meijerink οἱ 22008).," This is consistent, in general, with models of the ionization and thermal structure of T Tauri disks (e.g., Glassgold et 2007; Meijerink et 2008)."
 The observed [NelI] line profiles show little evidence for an origin in a disk photoevaporative flow., The observed [NeII] line profiles show little evidence for an origin in a disk photoevaporative flow.
 This is particularly so in the case of GM Aur. where the central component of the [Nell] enission is narrow enough to plausibly arise in a photoevaporative flow. but the prolile is centered αἱ the stellar velocity rather than showing the 5—IOkms! blueshilt that is predicted for emission arising in a photoevaporative flow (Alexander 2008).," This is particularly so in the case of GM Aur, where the central component of the [NeII] emission is narrow enough to plausibly arise in a photoevaporative flow, but the profile is centered at the stellar velocity rather than showing the $5-10\kms$ blueshift that is predicted for emission arising in a photoevaporative flow (Alexander 2008)."
 Photoevaporative flows may be present in these systems but contribute a small fraction of the [NelI] emission., Photoevaporative flows may be present in these systems but contribute a small fraction of the [NeII] emission.
 More recent work bv Pascucci Sterzik (2009) finds stronger evidence for photoevaporative flows in other [NelI|-emiting systems., More recent work by Pascucci Sterzik (2009) finds stronger evidence for photoevaporative flows in other [NeII]-emitting systems.
 The equivalent width of the [Nel] emission we detect is less than that of the spectrally unresolved [Nel] feature in theSpitzer spectra of the same sources., The equivalent width of the [NeII] emission we detect is less than that of the spectrally unresolved [NeII] feature in the spectra of the same sources.
 Variability in the [Nell] emission or the mid-infrared. continuum. a spatially extended. [Nel] component. or a very (spectrally) broad [Nel] component might account for the difference in the equivalent widths.," Variability in the [NeII] emission or the mid-infrared continuum, a spatially extended [NeII] component, or a very (spectrally) broad [NeII] component might account for the difference in the equivalent widths."
 Further work is needed to understand the origin of this discrepancy., Further work is needed to understand the origin of this discrepancy.
 These results illustrate the ability of high resolution spectroscopy to probe the origin of the Νο) emission from T Tauri stars., These results illustrate the ability of high resolution spectroscopy to probe the origin of the [NeII] emission from T Tauri stars.
 Further measurements of resolved [Nel] line profiles. for both classical T Tauri stars and transition objects. are needed to determine whether the results obtained here apply to the majority of [NelI|-emitting T Tauri stars.," Further measurements of resolved [NeII] line profiles, for both classical T Tauri stars and transition objects, are needed to determine whether the results obtained here apply to the majority of [NeII]-emitting T Tauri stars."
 In addition. hieher sensitivitv line profiles than those reported here would be useful to probe the origin of the [Nell] emission in detail.," In addition, higher sensitivity line profiles than those reported here would be useful to probe the origin of the [NeII] emission in detail."
 We thank Nathan Crockett for his help reducing the CO spectrum of AA Tat aud Colette Salvk lor making available her published CO spectrum of GM. Aur., We thank Nathan Crockett for his help reducing the CO spectrum of AA Tau and Colette Salyk for making available her published CO spectrum of GM Aur.
 We thank the Gemini stall for (heir support of TEXES observations on Gemini North., We thank the Gemini staff for their support of TEXES observations on Gemini North.
 The development of TEXES was supported by grants from the NSF ancl the NASA/USRA SOFIA project., The development of TEXES was supported by grants from the NSF and the NASA/USRA SOFIA project.
 Mocification ol TEXES [or use on Gemini was supported by Gemini Observatory., Modification of TEXES for use on Gemini was supported by Gemini Observatory.
 Observations with TEXES were supported by NSF erant AST-0607312., Observations with TEXES were supported by NSF grant AST-0607312.
 Financial support for the work of JRN and GWD was provided by the NASA Origins of Solar Systems program (NNIIOTAGOLI) and the NASA Astrobiology Institute under Cooperative Agreement CCAN-02-OSS-02 issued through the Office of Space Science., Financial support for the work of JRN and GWD 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 by the Lile and Planets Astrobiology Center (LAPLACE)., This work was also supported by the Life and Planets Astrobiology Center (LAPLACE).
 Basic research in infrared astronomy at the Naval Research Laboratory is supported by 6.1 base fueling., Basic research in infrared astronomy at the Naval Research Laboratory is supported by 6.1 base funding.
" MJI acknowledges support [rom NSF erant AST-0708074 and NASA erant. NNGOLGG92G, This work is based on observations", MJR acknowledges support from NSF grant AST-0708074 and NASA grant NNG04GG92G. This work is based on observations
parameters as the typical electron. momentum py. electron dispersion in momentum App. pitch-angle boundary of the loss-cone ας (or 4= cosa). and the loss-cone boundary width Aw.,"parameters as the typical electron momentum $p_{\mathrm{b}}$, electron dispersion in momentum $\Delta p_{\mathrm{b}}$, pitch-angle boundary of the loss-cone $\alpha_{\mathrm{c}}$ (or $\mu_{\mathrm{c}}=\cos\alpha_{\mathrm{c}}$ ), and the loss-cone boundary width $\Delta\mu_{\mathrm{c}}$."
 An example of the distribution function (3)) can be seen at Fig., An example of the distribution function \ref{fb}) ) can be seen at Fig.
 2aa. For ας=0. we obtain an isotropic ring-like distribution.," \ref{ev_b_V}a a. For $\alpha_{\mathrm{c}}=0$, we obtain an isotropic ring-like distribution."
 The initial energy density of plasma waves is assumed to equal the level of thermal oscillations: where Ky 1s the Boltzmann constant and Το is the effective plasma temperature., The initial energy density of plasma waves is assumed to equal the level of thermal oscillations: where $k_{\mathrm{B}}$ is the Boltzmann constant and $T_0$ is the effective plasma temperature.
" In addition. it is assumed that the energy density of plasma waves cannot fall below the thermal level (4)) during the process of wave/particle evolution. due to spontaneous radiation,"," In addition, it is assumed that the energy density of plasma waves cannot fall below the thermal level \ref{W0}) ) during the process of wave/particle evolution, due to spontaneous radiation."
" In. all simulations. we assume that the nonthermal distribution function (3)) has App/pp=0.2 and Aj,=0.2."," In all simulations, we assume that the nonthermal distribution function \ref{fb}) ) has $\Delta p_{\mathrm{b}}/p_{\mathrm{b}}=0.2$ and $\Delta\mu_{\mathrm{c}}=0.2$."
 The thermal component of the plasma (if present) is described by a maxwellian distribution with temperature of 10° K. The initial temperature of plasma waves Ty equals 10° K. The remaining parameters of the considered simulation models are given in Table 1: they were chosen in order to explore the influence of various factors on the process of wave/particle evolution., The thermal component of the plasma (if present) is described by a maxwellian distribution with temperature of $10^6$ K. The initial temperature of plasma waves $T_0$ equals $10^6$ K. The remaining parameters of the considered simulation models are given in Table \ref{tab1}; they were chosen in order to explore the influence of various factors on the process of wave/particle evolution.
 In all cases. the plasma density is relatively low. so that the ratio c/c varies in the range from 107 to 107: the total electron concentration 7 is calculated using the plasma frequency wp.," In all cases, the plasma density is relatively low, so that the ratio $\omega_{\mathrm{p}}/\omega_{\mathrm{B}}$ varies in the range from $10^{-3}$ to $10^{-1}$ ; the total electron concentration $n$ is calculated using the plasma frequency $\omega_{\mathrm{p}}$."
 The relative concentration of the energetic electrons 5/7 varies from 107 to I., The relative concentration of the energetic electrons $n_{\mathrm{b}}/n$ varies from $10^{-4}$ to 1.
 Firstly. we consider in detail an example of coevolution of the electron distribution and plasma waves.," Firstly, we consider in detail an example of coevolution of the electron distribution and plasma waves."
" We assume here that a low-temperature thermal component ts absent (2,= 7).", We assume here that a low-temperature thermal component is absent $n_{\mathrm{b}}=n$ ).
 Investigations of the dispersion relations for weakly-relativistic plasmas (Robinson 1986.. 1987)) have shown that the wave dispersion (for the waves propagating across the magnetic field. [cos9|« 1) becomes similar to that in vacuum 1f the typical electron speed vy satisfies the condition (a/c)2(wp/wyy.," Investigations of the dispersion relations for weakly-relativistic plasmas (Robinson \cite{rob86}, \cite{rob87}) ) have shown that the wave dispersion (for the waves propagating across the magnetic field, $|\cos\theta|\ll 1$ ) becomes similar to that in vacuum if the typical electron speed $\varv_{\mathrm{e}}$ satisfies the condition $(\varv_{\mathrm{e}}/c)\gtrsim (\omega_{\mathrm{p}}/\omega_{\mathrm{B}})^2$."
 Such a requirement is satisfied. e.g.. for the particle energy Ey23 keV and c/cg<0.1.," Such a requirement is satisfied, e.g., for the particle energy $E_{\mathrm{b}}\gtrsim 3$ keV and $\omega_{\mathrm{p}}/\omega_{\mathrm{B}}\lesssim 0.1$."
 Thus we assume that the wave refraction index equals unity both for the ordinary and extraordinary. modes., Thus we assume that the wave refraction index equals unity both for the ordinary and extraordinary modes.
 Also we assume that the waves are elliptically polarized with the axial ratio of the polarization ellipse 7j;=cos@ for the extraordinary mode and Το=—-]/ecos8 for the ordinary mode (TETQ=-1)., Also we assume that the waves are elliptically polarized with the axial ratio of the polarization ellipse $T_{\mathrm{E}}=\cos\theta$ for the extraordinary mode and $T_{\mathrm{O}}=-1/\cos\theta$ for the ordinary mode $T_{\mathrm{E}}T_{\mathrm{O}}=-1$ ).
" The above relations follow from the magnetoionic theory when wy/w| and ωρ/ω—0 (Melrose Dulk 1991: Willes. Melrose. Robinson 1994)); however. we found that the simulation results are not very sensitive to the exact value of the axial ratio T,, provided that for the quasi-transversal propagation |7).|«| (and. accordingly. |To|>> 1)."," The above relations follow from the magnetoionic theory when $\omega_{\mathrm{B}}/\omega\to 1$ and $\omega_{\mathrm{p}}/\omega\to 0$ (Melrose Dulk \cite{mel91}; Willes, Melrose, Robinson \cite{wil94}) ); however, we found that the simulation results are not very sensitive to the exact value of the axial ratio $T_{\sigma}$ provided that for the quasi-transversal propagation $|T_{\mathrm{E}}|\ll 1$ (and, accordingly, $|T_{\mathrm{O}}|\gg 1$ )."
" As an illustration. the following parameters were chosen: magnetic field B=1430 G that corresponds to the electron cyclotron frequency of fy=+ GHz (a typical value for the radio emission of ultracool dwarfs). plasma to cyclotron frequency ratio c/«p=107 that corresponds to the electron concentration 4j=Amy2x10? em7. typical energy of the energetic electrons £j,=10 keV. and the loss-cone boundary ας=60°."," As an illustration, the following parameters were chosen: magnetic field $B=1430$ G that corresponds to the electron cyclotron frequency of $f_{\mathrm{B}}=4$ GHz (a typical value for the radio emission of ultracool dwarfs), plasma to cyclotron frequency ratio $\omega_{\mathrm{p}}/\omega_{\mathrm{B}}=10^{-3}$ that corresponds to the electron concentration $n=n_{\mathrm{b}}=2\times 10^5$ $\mathrm{cm}^{-3}$, typical energy of the energetic electrons $E_{\mathrm{b}}=10$ keV, and the loss-cone boundary $\alpha_{\mathrm{c}}=60^{\circ}$."
 In Table I.. this parameter set corresponds to the model 15.," In Table \ref{tab1}, this parameter set corresponds to the model 15."
 Figure | shows the contours of growth rates of the ordinary and extraordinary modes at the initial moment (¢= 0)., Figure \ref{incr_V} shows the contours of growth rates of the ordinary and extraordinary modes at the initial moment $t=0$ ).
 Only the positive growth rates are shown: the contour levels are evenly distributed between zero and the maximal growth rate value for a given mode which ts also shown at the figure., Only the positive growth rates are shown; the contour levels are evenly distributed between zero and the maximal growth rate value for a given mode which is also shown at the figure.
 One can see that the most effective wave amplification takes place slightly below the fundamental cyclotron frequency., One can see that the most effective wave amplification takes place slightly below the fundamental cyclotron frequency.
 Both emission modes are generated mainly in the perpendicular direction with respect to the magnetic field. with a slight asymmetry due tothe loss-cone feature.," Both emission modes are generated mainly in the perpendicular direction with respect to the magnetic field, with a slight asymmetry due tothe loss-cone feature."
 Growth rates decrease rapidly with an increasing frequency. but the wave amplification (in an oblique direction) can oecur even above the cyclotron frequency: in this region. the emission directivity patterns become essentially," Growth rates decrease rapidly with an increasing frequency, but the wave amplification (in an oblique direction) can occur even above the cyclotron frequency; in this region, the emission directivity patterns become essentially"
"Using a combination of CO and mmeasurements in the local universe, Obreschkow&Rawlings(2009) have determined the density of H52 at z=0, pg,(0), to be 1.9 — 2.8 x10"" h MoMpc?.","Using a combination of CO and measurements in the local universe, \cite{OR2009} have determined the density of 2 at $z = 0$, $\mrh2(0)$, to be 1.9 – 2.8 $\times 10^7$ h $_\sun\permpccu$."
 'The range of values corresponds to different assumptions about how metallicity affects the CO-to-H32 conversion ratio (the authors quote an error of +40% on each individual calculated value)., The range of values corresponds to different assumptions about how metallicity affects the 2 conversion ratio (the authors quote an error of $\pm 40\%$ on each individual calculated value).
" Using h=0.7 gives ρη.(0)=x10"" Mpc, or an average value of 1.65x107 (1.3—2.0)MMpc-?."," Using $h=0.7$ gives $\mrh2(0) = (1.3 - 2.0) \times 10^7$ $_\sun\permpccu$, or an average value of $1.65 \times 10^7$ $_\sun\permpccu$."
" MeThis value is within of the estimate of py,(0)=1.1x107 MoMpc? by Zwaan&Prochaska(2006) (no error quoted) using a different set of observations and should therefore be reasonably reliable.", This value is within of the estimate of $\mrh2(0) = 1.1 \times 10^7$ $_\sun\permpccu$ by \cite{zp2006} (no error quoted) using a different set of observations and should therefore be reasonably reliable.
" A number of recent studies have reported measurements of the mean ccomoving mass density in galaxies, pyr(z), from the local universe to z~6."," A number of recent studies have reported measurements of the mean comoving mass density in galaxies, $\mrhi(z)$, from the local universe to $z\sim 6$."
" At z= 0, Zwaanetal.(2005) present a study of the mmass function using the PParkes All Sky Survey (HIPASS) data."," At $z=0$ , \cite{Zwaan2005} present a study of the mass function using the Parkes All Sky Survey (HIPASS) data."
" At low redshift, Lahetal.(2007) calculate by co-adding 221-cm emission from galaxies with known positions and redshifts."," At low redshift, \cite{Lah2007} calculate by co-adding 21-cm emission from galaxies with known positions and redshifts."
" At higher redshifts, estimates of are mainly obtained from DLA studies (Raoetal.2006;Prochaska&Wolfe 2009)."," At higher redshifts, estimates of are mainly obtained from DLA studies \citep{Rao2006,
 PW2009}."
. We use these observations of tto construct an analytical form for as a function of z., We use these observations of to construct an analytical form for as a function of $z$.
 We fit a straight line to these points in log(1+z) vs space., We fit a straight line to these points in $\log(1+z)$ vs $\log(\mrhi)$ space.
 The data points and the fit are shown in Fig. log(our)2.., The data points and the fit are shown in Fig. \ref{HIfig}.
" The observations suggest very little evolution of,, which increases by only factor of 2-3 between z—0 and z—6."," The observations suggest very little evolution of, which increases by only a factor of 2-3 between $z=0$ and $z=6$."
 A more recent DLAa study (Noterdaemeetal. using SDSS DR7 finds the values at z~2 to 3 to be 2009)somewhat higher than the Prochaska&Wolfe(2009) points; this would make our linear fit in Fig., A more recent DLA study \citep{Noterdaeme2009} using SDSS DR7 finds the values at $z\sim 2$ to 3 to be somewhat higher than the \cite{PW2009} points; this would make our linear fit in Fig.
 2 correspond even better to the observations., \ref{HIfig} correspond even better to the observations.
" The observations above allow us to compare the evolution of the SFRD and dpg,/dt.", The observations above allow us to compare the evolution of the SFRD and $d\mrh2/dt$.
" We start with the simplest model possible: a closed box of H22 being turned into stars at the rate of dp,/dt.", We start with the simplest model possible: a closed box of 2 being turned into stars at the rate of $d\mrh2/dt$.
" In this closed box model, we initially assume that the MGDR is constant in time (the restricted closed box model; refRCBM)) and subsequently relax this condition (the general closed box model; refclosedbox))."," In this closed box model, we initially assume that the MGDR is constant in time (the restricted closed box model; \\ref{RCBM}) ) and subsequently relax this condition (the general closed box model; \\ref{closedbox}) )."
" The failure of both models motivates us to consider an open box model refopenbox)), where we allow the(6 densities of all four phases of the IGM - stars, H92, and -— to vary to match the observational constraints."," The failure of both models motivates us to consider an open box model \\ref{openbox}) ), where we allow the densities of all four phases of the IGM – stars, 2, and – to vary to match the observational constraints."
" We assume that the molecular gas is depleted only through star formation, and that any Ἡο dissociated or ionized by star formation is instantaneously returned to the molecular state."," We assume that the molecular gas is depleted only through star formation, and that any $\mathrm{H_2}$ dissociated or ionized by star formation is instantaneously returned to the molecular state."
" Given the short timescales for the formation of molecular gas from its atomic form, ~10° yr at the relevant densities (Hollenbach&Salpeter1971;Cazaux&Tielens 2004),, this approximation should be a good one for the purposes of this paper."," Given the short timescales for the formation of molecular gas from its atomic form, $\sim10^6$ yr at the relevant densities \citep{hs1971,ct2004}, this approximation should be a good one for the purposes of this paper."
" In any event, if we define doy,/dt to be thenet flow rate of molecular gas into stars, then there is no ambiguity."," In any event, if we define $d\rho_{H_2}/dt$ to be the flow rate of molecular gas into stars, then there is no ambiguity."
" We write the statement that star formation occurs only through the depletion of molecular gas: 'This equation was used to infer the individual MGDR for a sample of nearby galaxies with a wide of range of SFRs and H» column densities, where it was found that MGDRz0.5 Gyr-! to a remarkable constancy (Leroyetal. 2008,, see refmgdrsec))."," We write the statement that star formation occurs only through the depletion of molecular gas: This equation was used to infer the individual MGDR for a sample of nearby galaxies with a wide of range of SFRs and $H_2$ column densities, where it was found that $MGDR \approx 0.5$ $^{-1}$ to a remarkable constancy \citealt{Leroy2008}, see \\ref{mgdrsec}) )."
 We use this prescription for the SFRD for our models but we note that the form may change at higher redshift (See refSFRDvar))., We use this prescription for the SFRD for our models but we note that the form may change at higher redshift (See \\ref{SFRDvar}) ).
" It is also interesting to note that if we divide the observed global star formation rate density SFRD(z—0)~ (0.8 - 1.8) x107?M5Mpc? yr-! (Hopkins& by the observed ρη.(0)=(1.3—2.0)x107MMpc (Obreschkow&Rawlings 2009), we obtain a range of MGDR that is consistent with the values of Leroyetal.(2008)."," It is also interesting to note that if we divide the observed global star formation rate density $(z=0) \sim$ (0.8 – 1.8) $\times 10^{-2}
 M_\sun\permpccu$ $^{-1}$ \citep{HB2006, Salim} by the observed $\mrh2(0) = (1.3 - 2.0) \times 10^7 M_\sun\permpccu$ \citep{OR2009}, we obtain a range of MGDR that is consistent with the values of \cite{Leroy2008}."
". Given that stars must form from molecular gas, this result is not surprising."," Given that stars must form from molecular gas, this result is not surprising."
" Nevertheless, the agreement is reassuring because it is based on different data sets and different methods of determining the relevant quantities."," Nevertheless, the agreement is reassuring because it is based on different data sets and different methods of determining the relevant quantities."
" It also suggests, combined with the arguments in refintro that Eq."," It also suggests, combined with the arguments in \\ref{intro} that Eq."
 1 can be extrapolated to all z., \ref{sfe} can be extrapolated to all $z$.
" We use mass densities instead of mass surface densities, which are used by observers, but note that these are roughly equivalent because for the most part, the stars, aand H32 aregenerally confined to thin disks within galaxies."," We use mass densities instead of mass surface densities, which are used by observers, but note that these are roughly equivalent because for the most part, the stars, and 2 aregenerally confined to thin disks within galaxies."
" In the closed box model, we consider only stars and H32, and allow pz, to be converted into stars at the star formation rate density, SFRD:"," In the closed box model, we consider only stars and 2, and allow $\rho_{H_2}$ to be converted into stars at the star formation rate density, SFRD:"
"sources (Carilli,Gnedin.&Owen2002:FurlanettoLoeb.2002).","sources \citep{carilli02,furlanetto02}."
. But this method requires luminous sources of radio continuum at hieh redshifts. which remain largely speculative.," But this method requires luminous sources of radio continuum at high redshifts, which remain largely speculative."
 The GRB radio afterglows would be à natural candidate., The GRB radio afterglows would be a natural candidate.
 ILowever their radio fIuxes at high redshift have not been estimated in detail., However their radio fluxes at high redshift have not been estimated in detail.
 Furlanetto&Loeb(2002) suggested that GRD alterglows are (oo dim. but they did not explicitly caleulate the sell-absorption frequencies. nor did they include the reverse shock emission.," \citet{furlanetto02} suggested that GRB afterglows are too dim, but they did not explicitly calculate the self-absorption frequencies, nor did they include the reverse shock emission."
 Another (vpe of related energetic radio source which has not been investigated in this respect are hvpernovae (IINe). which max be as [recquent as GRBs.," Another type of related energetic radio source which has not been investigated in this respect are hypernovae (HNe), which may be as frequent as GRBs."
 In (his paper. we investigate ihe GRB and LN radio afterglows at high. redshifts. including also the possibly of more energetic versions of (hese objects which may be associated with a population HI of stars.," In this paper, we investigate the GRB and HN radio afterglows at high redshifts, including also the possibly of more energetic versions of these objects which may be associated with a population III of stars."
 We first estimate the maximum redshifts out to which these alterglows could be detected with the Verv Large Array Low Frequency Array and Square Kilometer Array This can be important for broadband afterglow fits aàmed at determining total energies and physical parameters (e.g..Panailescu&Kumar2002) of bursts. oul to very hieh redshifts.," We first estimate the maximum redshifts out to which these afterglows could be detected with the Very Large Array Low Frequency Array and Square Kilometer Array This can be important for broadband afterglow fits aimed at determining total energies and physical parameters \citep[e.g.,][]{panaitescu02} of bursts, out to very high redshifts."
 We also quantify the detectabilitv of the 21 em absorption line in the allerelows., We also quantify the detectability of the 21 cm absorption line in the afterglows.
 since (he afterglow model parameters. such as (he energy and duration of the GRBs as well as (he ambient density may be different at high redshifts. we examine several possibilities of ihe model parameters.," Since the afterglow model parameters, such as the energy and duration of the GRBs as well as the ambient density may be different at high redshifts, we examine several possibilities of the model parameters."
 However. even with the SIXA. which has the best sensitivity among the proposed telescopes. and wilh our improved spectral-temporal evolution moclel. we find that the chances are low lor the detection of 21 cim absorption lines from GRBs auc Ne. except in (he case of exceplionally energetic sources.," However, even with the SKA, which has the best sensitivity among the proposed telescopes, and with our improved spectral-temporal evolution model, we find that the chances are low for the detection of 21 cm absorption lines from GRBs and HNe, except in the case of exceptionally energetic sources."
 This equantilication exercise should be useful for (he assessment and planning of future observing projects wilh advanced facilities and detectors., This quantification exercise should be useful for the assessment and planning of future observing projects with advanced facilities and detectors.
 The paper is organized as follows., The paper is organized as follows.
 In 27.. we calculate the GRB and. LN radio alterglows al hieh recdshilt. based on theafterelow model summarized in ?7..," In \ref{sec:flux}, we calculate the GRB and HN radio afterglows at high redshift, based on theafterglow model summarized in \ref{sec:model}."
 By comparing the afterglow fluxes with the VLA. LOFAR and SIXA sensitivities. we discuss (he maxiniun redshift out to which the radio afterglow of these sources would be detectable.," By comparing the afterglow fluxes with the VLA, LOFAR and SKA sensitivities, we discuss the maximum redshift out to which the radio afterglow of these sources would be detectable."
 We calculate the 21 em line absorption in the IGM ancl intervening structures al 2Z Gin ?? and ad αςσω ος ??.., We calculate the 21 cm line absorption in the IGM and intervening structures at $z \simg 6$ in \ref{sec:z>6} and at $z \siml 6$ in \ref{sec:z<6}.
 Our conclusions are sununarized in ??.., Our conclusions are summarized in \ref{sec:con}. .
 Throughout the paper we use a, Throughout the paper we use a
"pair (®,;,A®,).","pair $(\Phi_{\mathrm{l},i},\Delta\Phi_\mathrm{c})$."
" Since Ab,=O(10*km?s?Mpc?h?), we choose the next step to be AD;=Ad,+6(A®) with δ(ΔΦ)=10?km?s?Mpc?I? and calculate now R and the spectral moments for the pair (Di;,AD;)."," Since $\Delta\Phi_\mathrm{c}=\mathcal{O}(10^4\ \mbox{km}^2\ \mbox{s}^{-2}\ \mbox{Mpc}^{-2}\ h^2)$, we choose the next step to be $\Delta\Phi_\mathrm{1}=\Delta\Phi_\mathrm{c}+\delta(\Delta\Phi)$ with $\delta(\Delta\Phi)=10^2\ \mbox{km}^2\ \mbox{s}^{-2}\ \mbox{Mpc}^{-2}\ h^2$ and calculate now $R$ and the spectral moments for the pair $(\Phi_{\mathrm{l},i},\Delta\Phi_1)$."
" We continue with AO;=AO,+6(A®) for a given ©); until we fulfill the following convergence criterion."," We continue with $\Delta\Phi_j=\Delta\Phi_{j-1}+\delta(\Delta\Phi)$ for a given $\Phi_{\mathrm{l},i}$ until we fulfill the following convergence criterion."
 We approximate the integral in Eq. (49)), We approximate the integral in Eq. \ref{eq:intPotFunc}) )
 bythe trapezium rule., bythe trapezium rule.
" Hence, for a specific potential value ®);, the number density n(®,) is evaluated numerically as This summation is stopped at an index N chosen as the first index for which the relative contribution to n(0;) is smaller than 107$."," Hence, for a specific potential value $\Phi_{\mathrm{l},i}$, the number density $n(\Phi_{\mathrm{l},i})$ is evaluated numerically as This summation is stopped at an index $N$ chosen as the first index for which the relative contribution to $n(\Phi_{\mathrm{l},i})$ is smaller than $10^{-6}$."
" This is a proper break condition since Π(Φι,A®) tends rapidly towards zero for AD—ov."," This is a proper break condition since $\tilde{n}(\Phi_\mathrm{l},\Delta\Phi)$ tends rapidly towards zero for $\Delta\Phi\to\infty$."
" Having evaluated n(T;) for each i, we have an appropriate approximation of the X-ray temperature function's shape in the chosen temperature interval."," Having evaluated $n(T_i)$ for each $i$, we have an appropriate approximation of the X-ray temperature function's shape in the chosen temperature interval."
" As mentioned before, we need a proper high-pass filter in order to remove disturbing large-scale potential modes since we want the potential associated with a collapsed structure to be defined with respect to the large-scale potential level in its direct vicinity, and we want the structure to haveno peculiar motion so that the constraint 7=0 is applicable."," As mentioned before, we need a proper high-pass filter in order to remove disturbing large-scale potential modes since we want the potential associated with a collapsed structure to be defined with respect to the large-scale potential level in its direct vicinity, and we want the structure to haveno peculiar motion so that the constraint $\vec{\eta}=\vec{0}$ is applicable."
 The natural filter choice is a sharp cut-off in k-space since this will effectively remove both large-scale potential modes and potential gradients., The natural filter choice is a sharp cut-off in $k$ -space since this will effectively remove both large-scale potential modes and potential gradients.
 Let again R be the top-hat filter radius of Eq. (6)), Let again $R$ be the top-hat filter radius of Eq. \ref{eq:topHat}) )
" and Ry,the filter radius related to a sharp cut-off wave number kpp in k-space by kpp=271/Rhp.", and $R_\mathrm{hp}$the filter radius related to a sharp cut-off wave number $k_\mathrm{hp}$ in $k$ -space by $k_\mathrm{hp}=2\pi/R_\mathrm{hp}$.
" The ratio of both quantities defines a, Since R and Ry, define a low- and a high-pass filter, respectively, one can expect a>1."," The ratio of both quantities defines $\alpha$, Since $R$ and $R_\mathrm{hp}$ define a low- and a high-pass filter, respectively, one can expect $\alpha>1$."
 It is now a legitimate question how α should be chosen for calculating the correct number density of objects having a particular potential depth., It is now a legitimate question how $\alpha$ should be chosen for calculating the correct number density of objects having a particular potential depth.
 We argue that a natural choice exists., We argue that a natural choice exists.
" In Fig. B],"," In Fig. \ref{fig:maxima},"
" we plot Π(Τ,A®) as a function of a."," we plot $\tilde{n}(T,\Delta\Phi)$ as a function of $\alpha$."
" For each pair (T;,A9), the number density peaks at some value ama, which is a function of both the temperature and the Laplacian of the potential."," For each pair $(T_i,\Delta\Phi_k)$, the number density peaks at some value $\alpha_\mathrm{max}$, which is a function of both the temperature and the Laplacian of the potential."
 It increases with both increasing Laplacian and increasing temperature., It increases with both increasing Laplacian and increasing temperature.
" For low temperatures, the maximum is rather sharp, but broadens as the temperature increases."," For low temperatures, the maximum is rather sharp, but broadens as the temperature increases."
 This is most pronounced for the EdS universe., This is most pronounced for the EdS universe.
 This behaviour can be understood as follows., This behaviour can be understood as follows.
" On the one hand, decreasing the radius of the high-pass filter starting from a value much larger than the typical size of the object excludes more and more modes on decreasing scales."," On the one hand, decreasing the radius of the high-pass filter starting from a value much larger than the typical size of the object excludes more and more modes on decreasing scales."
" In this way, potential areremoved which would cause a non-zero peculiar velocity and thus a deviation from the constraint 7=0."," In this way, potential areremoved which would cause a non-zero peculiar velocity and thus a deviation from the constraint $\vec{\eta}=\vec{0}$."
" Since more and more objects are put to rest, the number density of objects with vanishing potential gradient increases."," Since more and more objects are put to rest, the number density of objects with vanishing potential gradient increases."
 Figure [5] shows that the increase of the number density is less steep for large (or hot) than for small (or cool) objects., Figure \ref{fig:maxima} shows that the increase of the number density is less steep for large (or hot) than for small (or cool) objects.
" This reflects the fact that that massive, hot objects, e.g. with 7=10 keV, are more likely located at potential minima, thus removing large-scale modes has less effect on their number counts."," This reflects the fact that that massive, hot objects, e.g. with $T=10$ keV, are more likely located at potential minima, thus removing large-scale modes has less effect on their number counts."
" On the other hand, increasing the radius of the high-pass filter starting from a value much smaller than the typical size of the relevant object adds more and more modes."," On the other hand, increasing the radius of the high-pass filter starting from a value much smaller than the typical size of the relevant object adds more and more modes."
" While the window between low- and high-pass filtering is too small, modes relevant for the structures considered are filtered out and the halo number density remains approximately zero."," While the window between low- and high-pass filtering is too small, modes relevant for the structures considered are filtered out and the halo number density remains approximately zero."
" Once modes are included that compose the structures, the halo number density steeply rises as α is increased."," Once modes are included that compose the structures, the halo number density steeply rises as $\alpha$ is increased."
" At a certain Q@max, the number density of objects reaches a maximum where both effects are balanced."," At a certain $\alpha_\mathrm{max}$ , the number density of objects reaches a maximum where both effects are balanced."
" Then, all modes relevant for structures of size R=./—20;/AQ; are included, but larger modes are excluded which would create a non-vanishing potential gradient."," Then, all modes relevant for structures of size $R=\sqrt{-2\Phi_i/\Delta\Phi_j}$ are included, but larger modes are excluded which would create a non-vanishing potential gradient."
" Thus, @max can be used to define Kmin for each pair (Τι,Aj) individually by This definition of kmin has to be used when evaluating the spectral moments with Eq. (24))."," Thus, $\alpha_\mathrm{max}$ can be used to define $k_\mathrm{min}$ for each pair $(T_i,\Delta\Phi_j)$ individually by This definition of $k_\mathrm{min}$ has to be used when evaluating the spectral moments with Eq. \ref{eq:specMomPot}) )."
" In appendix [AppendixÀj, we present an alternative way to define a physically reasonable cut-off wave vector for the evaluation of the spectral moments."," In appendix \ref{ap:kmin}, we present an alternative way to define a physically reasonable cut-off wave vector for the evaluation of the spectral moments."
" It turns out, however, that the number density of objects with a lowX-ray temperature function is highly underestimated in this way."," It turns out, however, that the number density of objects with a lowX-ray temperature function is highly underestimated in this way."
" In order to compare the X-ray temperature function that we have derived from the statistics of gravitational potential perturbations to the classical Press-Schechter theory, we need a proper, albeit idealised, and consistent mass-temperature relation."," In order to compare the X-ray temperature function that we have derived from the statistics of gravitational potential perturbations to the classical Press-Schechter theory, we need a proper, albeit idealised, and consistent mass-temperature relation."
" Since we used the virial theorem, Eq. (0),"," Since we used the virial theorem, Eq. \ref{eq:virTheorem}) ),"
" to relate the temperature to the potential, we will start at the same point to relate the temperature to a mass."," to relate the temperature to the potential, we will start at the same point to relate the temperature to a mass."
" Note that this has nothing to do with an assumption on real clusters, but merely serves thepurpose of a theoretical cross-comparison between the mass- Press-Schechter approach and our direct derivation of the temperature function."," Note that this has nothing to do with an assumption on real clusters, but merely serves thepurpose of a theoretical cross-comparison between the mass-based Press-Schechter approach and our direct derivation of the temperature function."
" We saw earlier that for a spherical and homogeneous overdensity the potential depth in the centre is 69=—2zGpR?, where p=ὄνρυ is the constant density inside theperturbation."," We saw earlier that for a spherical and homogeneous overdensity the potential depth in the centre is $\Phi_0=-2\pi G \bar{\rho} R^2$, where $\bar{\rho}=\delta_\mathrm{v}\rho_\mathrm{b}$ is the constant density inside the."
". Nonetheless, we shall replace 6, by the virial overdensity A,, which is a good approximation because ὃν=A,—1 and A,= O(10))."," Nonetheless, we shall replace $\delta_\mathrm{v}$ by the virial overdensity $\Delta_\mathrm{v}$ , which is a good approximation because $\delta_\mathrm{v}=\Delta_\mathrm{v}-1$ and $\Delta_\mathrm{v}=\mathcal{O}(10^2)$ ."
" The mass of the overdensity is M= απρκ», where ρ is the total density inside the sphere."," The mass of the overdensity is $M=\frac{4}{3}\pi\rho R^3$ , where $\rho$ is the total density inside the sphere."
 It is related tothe background density by p= Aypy., It is related tothe background density by $\rho=\Delta_\mathrm{v}\rho_\mathrm{b}$ .
" According to the previous statements, we can identify p and p, thus p« p."," According to the previous statements, we can identify $\bar{\rho}$ and $\rho$ , thus $\bar{\rho}\approx\rho$ ."
 Combining the equations for, Combining the equations for
of the bulge.,of the bulge.
 This radius corresponds to the core radius ry for the unperturbed isothermal sphere., This radius corresponds to the core radius $r_0$ for the unperturbed isothermal sphere.
 To caleulate any relationship between the DII mass and the velocity dispersion away from the centre of the svstem. we must use (he DF ealeulated. from the adiabatie growth framework.," To calculate any relationship between the BH mass and the velocity dispersion away from the centre of the system, we must use the DF calculated from the adiabatic growth framework."
 Fieure T shows the results for the three svstems studied. calculated at three clilferent radii: 2=10. near the centre: 2=1. at the core radius: and 2=100. well outside of the central region.," Figure \ref{fig:mbh-sigma} shows the results for the three systems studied, calculated at three different radii: $R = 10^{-3}$, near the centre; $R = 1$, at the core radius; and $R=100$, well outside of the central region."
 The simple predictions made above (24)) are confirmed (o exist in this model asvimptotically as the mass of the hole becomes large (the solid lines in Figure 7))., The simple predictions made above \ref{eq:mbh-sigma^2}) ) are confirmed to exist in this model asymptotically as the mass of the hole becomes large (the solid lines in Figure \ref{fig:mbh-sigma}) ).
 As the ealeulations are done farther out. however. the relation exists only at larger and larger masses (the dashed and dotted lines in Figure 7)).," As the calculations are done farther out, however, the relation exists only at larger and larger masses (the dashed and dotted lines in Figure \ref{fig:mbh-sigma}) )."
 Similar calculations for the self-similar and NEW distributions vield almost identical results., Similar calculations for the self-similar and NFW distributions yield almost identical results.
 The NEW DF is notable for attaining the asvanplolic relation at smaller black hole nasses (han for the isothermal case (the sell-similar DF is intermediate). but in all cases 10 linear relation steeper than that of equation (24)) is lound to exist.," The NFW DF is notable for attaining the asymptotic relation at smaller black hole masses than for the isothermal case (the self-similar DF is intermediate), but in all cases no linear relation steeper than that of equation \ref{eq:mbh-sigma^2}) ) is found to exist."
 There is. as can be seen in the figures. a steep shoulder during the approach to the asvimptotic limit but this is ion-linear and of insignificant extent.," There is, as can be seen in the figures, a steep shoulder during the approach to the asymptotic limit but this is non-linear and of insignificant extent."
 50 1 seems that (he required relation does not arise naturally in (he adiabatic growth yamework., So it seems that the required relation does not arise naturally in the adiabatic growth framework.
", This is not suprising. of course: after all. the DII does not create a disturbance nuch further than the radius 2=my,"," This is not suprising, of course; after all, the BH does not create a disturbance much further than the radius $R = m_{bh}$."
 Thus only the largest black holes can reach ont to the (vpical radius al which observations are taken. aud (hen they establish a much flatter relation.," Thus only the largest black holes can reach out to the typical radius at which observations are taken, and then they establish a much flatter relation."
 In the next section therefore we consider an alternative to adiabatic growth wherein the central black hole and the dark matter halo form together on the dvnamical time scale., In the next section therefore we consider an alternative to adiabatic growth wherein the central black hole and the dark matter halo form together on the dynamical time scale.
 The argument in (his section is somewhat more speculative (han in the preceding sections. ancl it must ultimately be checked by extensive numerical ealeulations.," The argument in this section is somewhat more speculative than in the preceding sections, and it must ultimately be checked by extensive numerical calculations."
 However the argument is compelling on dimensional grounds ancl is consistent will well known solutions and simulations in spherical svnuuetry (Ilenribsen&Widrow1984:Bertschinger 1935).," However the argument is compelling on dimensional grounds and is consistent with well known solutions and simulations in spherical symmetry \citep{hen99, fil84, ber85}."
. We assume (hat the galaxy forms by the extended collapse of a “halo” composed of collisionless matter. and that simultaneously the central black hole is growing proportionally to the halo as matter continues (o fall in.," We assume that the galaxy forms by the extended collapse of a “halo” composed of collisionless matter, and that simultaneously the central black hole is growing proportionally to the halo as matter continues to fall in."
 We do not assume spherical svinmetry., We do not assume spherical symmetry.
one exception). while those with AL<0.3 keV are the X-ray faintest.,"one exception), while those with $kT<0.3$ keV are the X-ray faintest."
" The least eas rich ETCs are then the coolest ones, Which secmed coutrary to expectation. if low Ly ETCs loose their ISM in au outflow (e... David et al."," The least gas rich ETGs are then the coolest ones, which seemed contrary to expectation, if low $L_X$ ETGs loose their ISM in an outflow (e.g., David et al."
 1990. Ciotti ct al.," 1990, Ciotti et al."
 1991). and the hotter the eas. the stronger is the outflow (BIKE).," 1991), and the hotter the gas, the stronger is the outflow (BKF)."
 Ou average TP iucreases with o... aud most ETCs lic above a rough cstimate of the eas viral temperature (T5=pinu? kk). sugecsting the xesence of additional heating.," On average $T$ increases with $\sigma_c$, and most ETGs lie above a rough estimate of the gas virial temperature $T_{\sigma}=\mu m_p
\sigma_c^2/k$ ), suggesting the presence of additional heating."
" ETCs with a moderate o high eas content (Ly>5«10°? ere 1) follow a rend roughly parallel to that of 75: iustead. ETCs with ittle hot. eas (Ly<5«1079 ere 1) have a similar eimperature for o, ranging from 160 to 250 kins +."," ETGs with a moderate to high gas content $L_X>5\times 10^{39}$ erg $^{-1}$ ) follow a trend roughly parallel to that of $T_{\sigma}$; instead, ETGs with little hot gas $L_X< 5 \times 10^{39}$ erg $^{-1}$ ) have a similar temperature for $\sigma_c$ ranging from 160 to 250 km $^{-1}$."
 This ack ofcorrelation was attributed to a different dynamical state of the hot ISM in gas-poor with respect to gas-rich ETCGs. though a full explanation of this aspect remained o be found (BKE).," This lack of correlation was attributed to a different dynamical state of the hot ISM in gas-poor with respect to gas-rich ETGs, though a full explanation of this aspect remained to be found (BKF)."
" This work takes advantage of the new accurate ucasureineuts of the hot eas properties. and of the ""undanental relatious Lya, and T0,. derived down to galaxw masses and X-rav huninosities smaller han ever before (BIKE). to investigate the relationship tween T. the ealaxy structure. the internal eas ισαπιο inechanisnis (SNL. aud those huked to the eravitational potential). and the dynamical status of he eas flow."," This work takes advantage of the new accurate measurements of the hot gas properties, and of the fundamental relations $L_X-\sigma_c$ and $T-\sigma_c$, derived down to galaxy masses and X-ray luminosities smaller than ever before (BKF), to investigate the relationship between $T$, the galaxy structure, the internal gas heating mechanisms (SNIa's, and those linked to the gravitational potential), and the dynamical status of the gas flow."
 To his purpose. a few characteristic cluperatures are introduced. depending on the nature of the eas heating sources and the galaxw structure. and relevant for the various gas flow phases: these characteristic temperatures are then compared with the observed TZ values.," To this purpose, a few characteristic temperatures are introduced, depending on the nature of the gas heating sources and the galaxy structure, and relevant for the various gas flow phases; these characteristic temperatures are then compared with the observed $T$ values."
 Tn doing so. galaxy mass models are nult according to the most recent understanding of the ETCs structure. such as their stellar mass profile aud heir dark matter content aud distribution. as mdicated x detailed modeling of optical observations aud by the nain scaling laws (c.¢.. Cappellani et al.," In doing so, galaxy mass models are built according to the most recent understanding of the ETGs' structure, such as their stellar mass profile and their dark matter content and distribution, as indicated by detailed modeling of optical observations and by the main scaling laws (e.g., Cappellari et al."
 2006. Weijuians et al.," 2006, Weijmans et al."
 2009. Auger ct al.," 2009, Auger et al."
 2010. Napolitano ct al.," 2010, Napolitano et al."
 2010.mn Shen Cebhardt 2010).," 2010, Shen Gebhardt 2010)."
 The ais are to address the ollowiug «questious: cau the eas heating sources above account for the observed F's?, The aims are to address the following questions: can the gas heating sources above account for the observed $T$ 's?
 how are the various iuput ΟΠΟΙΟΥ sources for the gas used iu the different flow yhases?, how are the various input energy sources for the gas used in the different flow phases?
 is there auv relation between f aud the flow oiase?, is there any relation between $T$ and the flow phase?
 We present in Sect., We present in Sect.
 ??. the sources of mass aud heating or the hot ISM. in Sect.," \ref{heat} the sources of mass and heating for the hot ISM, in Sect."
 ?? the conditions for the gas o escape from the galaxy. iun Sect.," \ref{esc} the conditions for the gas to escape from the galaxy, in Sect."
 ?? the galaxy imass uodoels. in Sect.," \ref{mass} the galaxy mass models, in Sect."
 77. the comparison between observed and predicted temperatures. iu Sect.," \ref{disc} the comparison between observed and predicted temperatures, in Sect."
 77. the relation )etween gas temperature aud flow status. aud in Sect.," \ref{temp} the relation between gas temperature and flow status, and in Sect."
 the conclusions., \ref{concl} the conclusions.
 Iu ETCs the hot gas comes frou stellar mass losses produced by evolved stars. mainly during the red eiaut. asvinptotic eiut branch. and planetary uchula phases. aud by SNlIas. that are the only ones observed iu an old stellar population (e.g.. Cappellaro ot al.," In ETGs the hot gas comes from stellar mass losses produced by evolved stars, mainly during the red giant, asymptotic giant branch, and planetary nebula phases, and by SNIa's, that are the only ones observed in an old stellar population (e.g., Cappellaro et al."
 1999)., 1999).
 The fist. more quiescent. type of losses originates ejecta that initially have the velocity of the parcut star. then individually iuteract with the mass lost from other stars or with the hot ISAL aud mix with it (Mathews 1990. Parriott Breeman 2008).," The first, more quiescent, type of losses originates ejecta that initially have the velocity of the parent star, then individually interact with the mass lost from other stars or with the hot ISM, and mix with it (Mathews 1990, Parriott Bregman 2008)."
 For a galaxy of total stellar mass M... the evolution of the stellar mass loss rate AL(t) can be caleulated using sinele burst stellar population svuthesis models (Maraston 2005). for a Salpeter aud for a. IWkroupa Tnitial Mass Function (AIF). assmuius for example solar abundance.," For a galaxy of total stellar mass $M_*$, the evolution of the stellar mass loss rate $\dot M_*(t)$ can be calculated using single burst stellar population synthesis models (Maraston 2005), for a Salpeter and for a Kroupa Initial Mass Function (IMF), assuming for example solar abundance."
 So doiug. at an age of 12 Cis. a rate is recovered of AL.= DB «10ἩLpiLp..)Mixx. +. where Lg ds the galactic D-baud hDwuninositv at an age of 12 Gyr. and Όλι or B=1.9 for the Salpeter or Iroupa IMF (see also Pellegrini 2011).," So doing, at an age of 12 Gyrs, a rate is recovered of $\dot M_* $ = B $\times 10^{-11} \, L_B(L_{B,\odot})\,\,
M_{\odot}$ $^{-1}$, where $L_B$ is the galactic B-band luminosity at an age of 12 Gyr, and B=1.8 or B=1.9 for the Salpeter or Kroupa IMF (see also Pellegrini 2011)."
 This value is in reasonable agrecnmient with the average derived for mine local ETCs from FSO data (Athey et al., This value is in reasonable agreement with the average derived for nine local ETGs from $ISO$ data (Athey et al.
 2002) of A—T78.102Lp(Lp.) Moy | an estimate based onu individual observed values that vary by a factor of ~10. though. which was attributed to differeut ages aud inetallicities.," 2002) of $\dot M_*=7.8\times
10^{-12}\,\, L_B(L_{B,\odot})$ $_{\odot}$ $^{-1}$, an estimate based on individual observed values that vary by a factor of $\sim 10$, though, which was attributed to different ages and metallicities."
 The total mass loss rate of a stellar population AJ is eiven by the sum AY=AL.|Max. where Max is the rate of mnass loss via SNIa eveuts for the whole galaxy., The total mass loss rate of a stellar population $\dot M$ is given by the sum $\dot M=\dot M_*+ \dot M_{\rm SN}$ where $ \dot M_{\rm SN}$ is the rate of mass loss via SNIa events for the whole galaxy.
 Max is even by Max=ΔίονRex. where May=1.LAL... is the ejected mass by one event. and Rex is the explosion rate.," $\dot
M_{\rm SN}$ is given by $\dot M_{\rm SN} =M_{SN}\, R_{\rm SN}$, where $M_{SN}=1.4M_{\odot}$ is the ejected mass by one event, and $R_{\rm
SN}$ is the explosion rate."
 Πως has heen determined for local ETCs to be Rex= 1l. where Fy is the Iubble coustant iu units of Jan + Mpe| (Cappellaro et al.," $R_{\rm SN}$ has been determined for local ETGs to be $R_{\rm SN}=0.16 (H_0/70)^2 \times 10^{-12} \,L_B
(L_{B,\odot}) \, {\rm yr}^{-1} $ , where $H_0$ is the Hubble constant in units of km $^{-1}$ $^{-1}$ (Cappellaro et al."
 1999)., 1999).
 More receut iieasureimenuts of the observed rates of supernovae in the local Universe (Li et al., More recent measurements of the observed rates of supernovae in the local Universe (Li et al.
 2010) eive a SNias rate in ETCs consisteut with that of Cappellaro et al. (, 2010) give a SNIa's rate in ETGs consistent with that of Cappellaro et al. (
1999).,1999).
 For Wy=70 lau + Mpe|l one obtains Max —22.410Πρίν.) Mi 1l. that is ~NO fines simaller than AL. derived above for an age of 12 Cr: therefore. the main source of mass for the hot eas is provided by AZ...," For $H_0=70$ km $^{-1}$ $^{-1}$, one obtains $\dot M_{\rm SN}$ $ \times 10^{-13}
L_B(L_{B,\odot}) $ $_{\odot}$ $^{-1}$, that is $\sim 80$ times smaller than $\dot M_* $ derived above for an age of 12 Gyr; therefore, the main source of mass for the hot gas is provided by $\dot M_*$ ."
 A reasonable assuniptiou is that the eas is shed by stars with a radial depeudeuce that follows that of the stellar distribution. so that the deusity profile of the injected gas is οκ)Mpal). where patr) is the stellar density profile.," A reasonable assumption is that the gas is shed by stars with a radial dependence that follows that of the stellar distribution, so that the density profile of the injected gas is $\rho_{gas}(r)\propto \rho_*(r)$, where $\rho_*(r)$ is the stellar density profile."
 This asswuptiou is adopted hereafter. aud the characteristic temperatures presented below apply to a eas distribution following Ομ)xpatr) (out see also Sect. ο).," This assumption is adopted hereafter, and the characteristic temperatures presented below apply to a gas distribution following $\rho_{gas}(r)\propto \rho_*(r)$ (but see also Sect. \ref{disc}) )."
 The material lost by stars is ejected at a velocity of few tens of kins + and at a temperature of <10! Is (Parriott Bregman 2008). auc is subsequently heated to high. X-rav cutting temperatures by the thermalization of the stellar velocity dispersion. as it collides with the mass ost from other stars. or with the ambicut hot gas. aud is shocked.," The material lost by stars is ejected at a velocity of few tens of km $^{-1}$ and at a temperature of $\lsim 10^4$ K (Parriott Bregman 2008), and is subsequently heated to high, X-ray emitting temperatures by the thermalization of the stellar velocity dispersion, as it collides with the mass lost from other stars, or with the ambient hot gas, and is shocked."
 Another source of heating for the stellar mass osses is provided by the thermalization of the kinetic cherey of SNIa’s events., Another source of heating for the stellar mass losses is provided by the thermalization of the kinetic energy of SNIa's events.
" The internal cucrey eiven bv hese heating processes to the wit mass of injected gas is BAT), (with & the Doltziuiumn constant. the xotou mass. 4/2),(01, the macau particle mass. with jii,—0.62 or solar abuudauce]: 7;,; 18 determiued by the heating due to thermalization of the motions of the eas-losing stars μες]. aud of the velocity of the SNIa's ejecta (Poy).and is written as (e.g.BS.Casler 1976. White"," The internal energy given by these heating processes to the unit mass of injected gas is $3kT_{inj}/2\mu m_p$ (with $k$ the Boltzmann constant, $m_p$ the proton mass, $\mu m_p$ the mean particle mass, with $\mu=0.62$ for solar abundance); $T_{inj}$ is determined by the heating due to thermalization of the motions of the gas-losing stars $T_{star}$ ), and of the velocity of the SNIa's ejecta $T_{SN}$ ),and is written as (e.g.,Gisler 1976, White"
 Tideutity teu hos stars of transiting exoplanet svstenis in which either the host star is more slowly rotating relativo to the simple model than any of the stars in the control sauple of SPOCS stays or in which he star is viewed pole-on aud the SYsten d8 spiu-orbit misaligned.,I identify ten host stars of transiting exoplanet systems in which either the host star is more slowly rotating relative to the simple model than any of the stars in the control sample of SPOCS stars or in which the star is viewed pole-on and the system is spin-orbit misaligned.
 This degeneracy cau IC broken for at cast I&kepler-5 and Iepler-7. as high xecisiou heper photometry should reveal their rotatio1] lods.," This degeneracy can be broken for at least Kepler-5 and Kepler-7, as high precision Kepler photometry should reveal their rotation periods."
 If the seriods are indeed short. then the smal| observed csiu? indicates spin-orbit misalienineut along the line of sieht.," If the periods are indeed short, then the small observed $v\sin{i}$ indicates spin-orbit misalignment along the line of sight."
 Alternatively. if the observed rotation xxiods. are long. hen EKepler-5 aud Iepler-7 are svstelis in which the jost star has somehow lost auguar monientuni nnch nore quickly than expecte by models of Sun-like stellar spin-down.," Alternatively, if the observed rotation periods are long, then Kepler-5 and Kepler-7 are systems in which the host star has somehow lost angular momentum much more quickly than expected by models of Sun-like stellar spin-down."
 High precision photomerv from other sources weht enable similar nieasuremoeuts for the other xvsteLis ideutified as :nalously sow projected rotators., High precision photometry from other sources might enable similar measurements for the other systems identified as anomalously slow projected rotators.
 All ten of the candidate uisaligeued systems have host stars inore lassive than the Sun., All ten of the candidate misaligned systems have host stars more massive than the Sun.
 Te» cletermine the sienificance of fus observation. I performed a Monte Carlo. simulation in whic1 1 randomly selected with replacement teu stars from the 175 listed in Table 1..," To determine the significance of this observation, I performed a Monte Carlo simulation in which I randomly selected with replacement ten stars from the 75 listed in Table \ref{tbl-1}."
 I repeated this srocess 109 times., I repeated this process $^6$ times.
 Less than of he Moute Carlo trials produced a sample in which all en stars had 1.2Af.zM.E1.5AL...," Less than of the Monte Carlo trials produced a sample in which all ten stars had $1.2~M_{\odot}
\lesssim M_{\ast} \lesssim 1.5~M_{\odot}$."
 Therefore. the xobaην that all ten caudidate iuisaligued svstenis woul be identified around massive stars bv chance is oss. fiui one iu teu thousand.," Therefore, the probability that all ten candidate misaligned systems would be identified around massive stars by chance is less than one in ten thousand."
 Siwilarly. the median uass of the planets in the candidate inajened svstenmis is 24.0.6 Jupiter-1nasses. while the lutiu mass. of he whole sample of 75 svstcus is DEGG Jupite-uasses.," Similarly, the median mass of the planets in the candidate misaligned systems is $2
\pm 0.6$ Jupiter-masses, while the median mass of the whole sample of 75 systems is $1 \pm 0.1$ Jupiter-masses."
 Tuterestinely. there may be a liut of the sale apparent overrepresentation oD qnassivo planets im spin-orbit misaieued svsteimis from Bossiter-MeLaughliliu neasureclents (Torresetal.2010).," Interestingly, there may be a hint of the same apparent overrepresentation of massive planets in spin-orbit misaligned systems from Rossiter-McLaughlin measurements \citep{tor10}."
. Ecceutricity also seems to be relaed to nisaliguimoenut xobabilitv., Eccentricity also seems to be related to misalignment probability.
 Thcced. of the mine transiting svstems with FGK stellar hosts ancl eccentricity ¢0.1 CID 80606. TWD 17156. ILAT-P-2. WASDP-5. XO-3. ILAT-P-]1. WASP-17. CoRoT-9. aid ITAT-P-1D. there is now evidence frou either measureneuts of the RM effect or this analysis that six of those svstoenis are significantly misaligned (ID 50606. ΠΟ 17156. WASDP-5. NO-3. WASP-17. aud ILAT-P-11).," Indeed, of the nine transiting systems with FGK stellar hosts and eccentricity $e > 0.1$ (HD 80606, HD 17156, HAT-P-2, WASP-8, XO-3, HAT-P-11, WASP-17, CoRoT-9, and HAT-P-14), there is now evidence from either measurements of the RM effect or this analysis that six of those systems are significantly misaligned (HD 80606, HD 17156, WASP-8, XO-3, WASP-17, and HAT-P-14)."
 The other lice systeus are unlikely to be observed to be spiu-orüt miusaliered for a iunber of reasons: ITAT-P-2 is likely au iutriusicallv. fast-rotator and therefore uulikelv to be ideifified as misaligned along the line of sight bw the nethod described im this paper even if it does have sienificaut spin-orlit nusaliguiuent along the ine of sight. ILAT-P-11 may have too small a radius for οο...its of the RAL effecti and CoRoT-9 has the largest periccuter distance by far of the known transiting planets and almost certainly has had a different evotionary history than the other known trausitiie planets.," The other three systems are unlikely to be observed to be spin-orbit misaligned for a number of reasons: HAT-P-2 is likely an intrinsically fast-rotator and therefore unlikely to be identified as misaligned along the line of sight by the method described in this paper even if it does have significant spin-orbit misalignment along the line of sight, HAT-P-11 may have too small a radius for measurements of the RM effect, and CoRoT-9 has the largest pericenter distance by far of the known transiting planets and almost certainly has had a different evolutionary history than the other known transiting planets."
 To determine the significance of the apparewt correlation between eccentricity and spin-orbit nisaligeumoenut. I performed a Monte Carlo simulation to determine the expected umber of systems with eccentricity e20.1 from a sample of 16 transiting systems — fhe total number of significantly iuisaligued svstenis from RAL measurements ando this analysis expected if spin-orbit musaliguiment is uurelated to eccentricity.," To determine the significance of the apparent correlation between eccentricity and spin-orbit misalignment, I performed a Monte Carlo simulation to determine the expected number of systems with eccentricity $e > 0.1$ from a sample of 16 transiting systems – the total number of significantly misaligned systems from RM measurements and this analysis -- expected if spin-orbit misalignment is unrelated to eccentricity."
 I x:udonlv selected with replacement the eccentricities of 16 planets from the 75 systems listed iu Table 1.. and [repeated this process 109 times.," I randomly selected with replacement the eccentricities of 16 planets from the 75 systems listed in Table \ref{tbl-1}, and I repeated this process $^6$ times."
 Efiud that the nean number of svsteius with «>OL in a sample of 16 svstenus 1uder the assumption that ecceutricity is unrelated to spin-orbit nüsalieument is 2+2., I find that the mean number of systems with $e > 0.1$ in a sample of 16 systems under the assumption that eccentricity is unrelated to spin-orbit misalignment is $2 \pm 2$.
 The observation of six svstems with ο20.1 from a random saluple of 16 svsteis occurs im less than oue in 150 trials. indicating that ecceutricity and spin-orbit misaligeuimoeut are related.," The observation of six systems with $e > 0.1$ from a random sample of 16 systems occurs in less than one in 150 trials, indicating that eccentricity and spin-orbit misalignment are related."
 Tn an attempt to derive some constraüut on the nuderlving spin-orbit misaliguiment distribution iu the 75 systems in Table 1.. I performed two completeness calculations for two differeut ἐς distributions resulti18o frou wo different c distributions: (1) correspondiug to plaue-plauet scattering (e.g.Chatterjeeetal.2008).. a Case 1i which the iuclination of the host star is completely indepeneut of the aliguiieut of the orbit of its planet le;™arecos(1UU). where (~nif(0.1): and (2) COLLCSPOxding to Ixozai eveles aud tidal friction (e.g.Fal-rvckv&Tremaine2007 ).. a case iu which the distribution of ἐς is derived. from the c distribution as elven iu Figure 10 of Fabrveky&Tremaine(2007) axl the A distribution as given iu the RM compilation in Table 5Γ of Tui et al. (," In an attempt to derive some constraint on the underlying spin-orbit misalignment distribution in the 75 systems in Table \ref{tbl-1}, I performed two completeness calculations for two different $i_s$ distributions resulting from two different $\psi$ distributions: (1) corresponding to planet-planet scattering \citep[e.g.][]{cha08}, a case in which the inclination of the host star is completely independent of the alignment of the orbit of its planet $i_{s} \sim \arccos(1-U)$, where $U \sim Unif\left(0,1\right)$; and (2) corresponding to Kozai cycles and tidal friction \citep[e.g.][]{fab07}, a case in which the distribution of $i_s$ is derived from the $\psi$ distribution as given in Figure 10 of \citet{fab07} and the $\lambda$ distribution as given in the RM compilation in Table 5 of Triaud et al. ("
2010. submitted).,"2010, submitted)."
 Asstuine a thiresrold cletectale iisalienuieut that corresponds to the suxlest degree oπαρα Ποιοι ideutified plszHO) πι το systems. I find that the expected Πιο of detections iu case (1) is 17+| and in case (2) it is FE ," Assuming a threshold detectable misalignment that corresponds to the smallest degree of misalignment identified $|i_p - i_s| \gtrsim 50^{\circ}$ ) in 75 systems, I find that the expected number of detections in case (1) is $17 \pm 4$ and in case (2) it is $27 \pm 4$."
Even though e:veh mmdiidual svsteimi duo thus veuls is securely identified as an anolously slow xojected rotator alc therefore possibly sou-orbit κααςred. voth completeness calculations are stronely depeudent on the smallest deeree of müsaligunieut this technique can ddentifv.," Even though each individual system in this analysis is securely identified as an anomalously slow projected rotator and therefore possibly spin-orbit misaligned, both completeness calculations are strongly dependent on the smallest degree of misalignment this technique can identify."
 TIlat quantity is uncertanm. and as such nore data is nec‘essary before any couclusive statements about distriution of spiu-orbi aligninents is made.," That quantity is uncertain, and as such more data is necessary before any conclusive statements about distribution of spin-orbit alignments is made."
 The ll s:uuple of Ixepler trausitiis planet detections aud lost star properties available at the eud of its four vear Mission nieht xsolve this issue., The full sample of Kepler transiting planet detections and host star properties available at the end of its four year mission might resolve this issue.
 Collectivev. this analvsis 1didicates that predictious [spurorbit iuisalieunieut for a population of close-in dlanets eutielv produced by planet-plauet scattering or Ilozai cycles with tidal fictioi overpredict the umber of lisajened systems.," Collectively, this analysis indicates that predictions of spin-orbit misalignment for a population of close-in planets entirely produced by planet-planet scattering or Kozai cycles with tidal friction overpredict the number of misaligned systems."
 As a result. there seeus to IO two populations of close-in planets: a popilation hat is spin-orbit aieued aud a population that is apparently spir-orbit nusaligued.," As a result, there seems to be two populations of close-in planets: a population that is spin-orbit aligned and a population that is apparently spin-orbit misaligned."
" The processes that lead 0o πα, or to the survival of misaligned svstenmis seen to operates nios efficiently in svstems with massive vost stars ale planets.", The processes that lead to misalignment or to the survival of misaligned systems seem to operates most efficiently in systems with massive host stars and planets.
" Indeed. the striking appearance of cancidate wisaligjd systems at AL,21.2M. in Figure | indicates that the xocesses that lead to spin-orbit nisalieuucut aid survival are threshold processes.4. simular to the ra pidiicrease In exoplauet incidence with host stellar metallicity (c.g.Saiosetal.2004:Fischer&Valenti 2005)."," Indeed, the striking appearance of candidate misaligned systems at $M_{\ast} \gtrsim 1.2~M_{\odot}$ in Figure \ref{fig04} indicates that the processes that lead to spin-orbit misalignment and survival are threshold processes, similar to the rapid increase in exoplanet incidence with host stellar metallicity \citep[e.g.][]{san04,fis05}."
. Th tjs case. the transition to frequent apparent spin-orbit παοπλο! occurs at tie sale stellar mass at which Sun-lise stars with near solar metallicitv develop radiative envelopes.," In this case, the transition to frequent apparent spin-orbit misalignment occurs at the same stellar mass at which Sun-like stars with near solar metallicity develop radiative envelopes."
 This dramatic change in stellar structure will liselv stronely effect the poorlv-uuderstood tidal processes that play a major role in the formation. evolution. aix lous-terii survival of spin-orbit nüsaligued systems.," This dramatic change in stellar structure will likely strongly effect the poorly-understood tidal processes that play a major role in the formation, evolution, and long-term survival of spin-orbit misaligned systems."
eas in the nebula to a cdilfuse IGM. but we will model all σας as either part of the nebula or part ofthe LGAL.,"gas in the nebula to a diffuse IGM, but we will model all gas as either part of the nebula or part of the IGM."
 Since the nebula and IGM are neutral in the absence of stellar radiation. both have the potential to play an important role in reprocessing ionizing photons from a star.," Since the nebula and IGM are neutral in the absence of stellar radiation, both have the potential to play an important role in reprocessing ionizing photons from a star."
 We examine two limiting cases for the importance of each of these phases: In the first case. sstars are enshrouded in dense nebulae. and all of the reprocessing of ionizing radiation takes place in the halo (the IGM still plays a role in scattering Lya photons).," We examine two limiting cases for the importance of each of these phases: In the first case, stars are enshrouded in dense nebulae, and all of the reprocessing of ionizing radiation takes place in the halo (the IGM still plays a role in scattering $\lya$ photons)."
 In this case. the escape fraction of ionizing photons from the nebula. fo... is zero.," In this case, the escape fraction of ionizing photons from the nebula, $f_{esc}$, is zero."
 In the second case. the nebula plays no role and fos.=1: all reprocessing occurs in the IGM.," In the second case, the nebula plays no role and $f_{esc}=1$; all reprocessing occurs in the IGM."
 In the rest of this section. we discuss the reprocessecl spectrum of sstellar racliation for each of these cases.," In the rest of this section, we discuss the reprocessed spectrum of stellar radiation for each of these cases."
 umerical simulations suggest that when a sstar forms. the nebula consists of a higher density phase. with ngc10em and a lower density. phase (e.g... Bronun ct al.," Numerical simulations suggest that when a star forms, the nebula consists of a higher density phase, with $n_{\mathrm
H}\simeq10^4~\cmt$, and a lower density phase (e.g., Bromm et al."
 1999)., 1999).
 We make the simplifving assumption hat half of the nebulas mass is contained in a homogeneous ohase with density ny=101em* that completely covers he star(s). and ignore the lower density σας.," We make the simplifying assumption that half of the nebula's mass is contained in a homogeneous phase with density $n_{\mathrm H}=10^4~\cmt$ that completely covers the star(s), and ignore the lower density gas."
 We take the mass fractions of hvdrogen and helium as Y=0.75 and )=0.25. respectively.," We take the mass fractions of hydrogen and helium as $X=0.75$ and $Y=0.25$, respectively."
 lonizing raciation from the star(s) creates an rregion in the dense nebula., Ionizing radiation from the star(s) creates an region in the dense nebula.
 Because of the hardness of our input spectrum. in the inner part of the rreeion helium is doubly ionized (the rregion).," Because of the hardness of our input spectrum, in the inner part of the region helium is doubly ionized (the region)."
 In the outer part it is singly ionized (the rregion: the spectrum is hard enough. that. there. is. no rregion)., In the outer part it is singly ionized (the region; the spectrum is hard enough that there is no region).
 The majority of photons above the iionization threshold ionize rrather thanIli., The majority of photons above the ionization threshold ionize rather than.
 We iteratively solve for the sizes and temperatures of these regions. using the thermocwnanic equations from Cen (1992).," We iteratively solve for the sizes and temperatures of these regions, using the thermodynamic equations from Cen \shortcite{cen92}."
. Phe rregion. comprising 0.4 of the volume of the rregion. has a temperature of 3.60!Is: the rregion is cooler. at 2107Ek.," The region, comprising $0.4$ of the volume of the region, has a temperature of $3.6\times10^4~\kelvin$; the region is cooler, at $2.7\times10^4~\kelvin$."
 In both regions& the 1primary cooling mechanism is rrecombination. but free-[ree cnission anc cooling via collisional excitation of aare also important.," In both regions the primary cooling mechanism is recombination, but free-free emission and cooling via collisional excitation of are also important."
 Ehe total nebular emission is not. very sensitive to the relative sizes of the aanel rregions., The total nebular emission is not very sensitive to the relative sizes of the and regions.
 Given these properties of an ionized region. we can determine fos.’ The volume and mass of the rregion are where Quo is the stellar. emission. rate of photons energetic enough to ionizeHel. αμο.7) is the Case 3 recombination coellicient. forHel1L. here AM. is the mass of the ionizing sstar(s) in the nebula and yp=1.2 is the mean molecular weight.," Given these properties of an ionized region, we can determine $f_{\mathrm esc}$ The volume and mass of the region are where $Q_{\mathrm HeII}$ is the stellar emission rate of photons energetic enough to ionize, $\alpha_{\mathrm B}(\heii,T)$ is the Case B recombination coefficient for, here $M_*$ is the mass of the ionizing star(s) in the nebula and $\mu=1.2$ is the mean molecular weight."
 Because in the rregion recombinations to pprovide enough photons to keep the hydrogen ionized. the volume and mass of the rregion are where Quy is the rate of Lli-tonizing photons. and api1) is the hydrogen recombination cocllicient.," Because in the region recombinations to provide enough photons to keep the hydrogen ionized, the volume and mass of the region are where $Q_{\mathrm HI}$ is the rate of -ionizing photons, and $\alpha_{\mathrm B}(\hi,T)$ is the hydrogen recombination coefficient."
 The total mass of ionized gas is then 1.7 times the mass of the ionizing star(s)., The total mass of ionized gas is then $1.7$ times the mass of the ionizing star(s).
" In the single-burst. star-formation model. the nebula mass in. ng=101em;"" gas. Mya. ds. For η=0.4. Mj=0.75M. which implies fis.=0 would be dillieult to achieve for our model nebula."," In the single-burst star-formation model, the nebula mass in $n_{\mathrm H}=10^4~\cmt$ gas, $M_{\mathrm neb}$ , is For $\eta=0.4$, $M_{\mathrm neb}=0.75\,M_*$, which implies $f_{\mathrm
esc}=0$ would be difficult to achieve for our model nebula."
 In the ongoing model. though. Adie.(19)M;/(24)) in general. ancl especially at. lower redshift.," In the ongoing model, though, $M_{\mathrm neb}>(1-\eta)M_*/(2\eta)$ in general, and especially at lower redshift."
 This is due to the short lifetime of sstaws: star formation in a halo subsequent to the first episode usually takes place after. previous &enerations of stars in the halo have stopped radiating., This is due to the short lifetime of stars: star formation in a halo subsequent to the first episode usually takes place after previous generations of stars in the halo have stopped radiating.
 Thus fi.=0 may be possible in these haloes., Thus $f_{\mathrm esc}=0$ may be possible in these haloes.
 Approximately 1/2 of the energy raciated by a sstarin a nebula is ultimately reraciated by recombinations., Approximately $1/2$ of the energy radiated by a starin a nebula is ultimately reradiated by recombinations.
 The mean energy of a free electron just before it recombines with a proton may be estimated from where Sp(lli.2) ds the. recombination emission coefficient.," The mean energy of a free electron just before it recombines with a proton may be estimated from where $\beta_{\mathrm B}(\hi,T)$ is the recombination emission coefficient."
Using 7=31U IK. we lind (£3=L4eV ,"Using $T=3\times10^4~\kelvin$ , we find $\left<E\right>=1.4~\ev$ "
from an astrophysical aid observational point of view is the most relevant quantity. depends on the temperature through the thermalized (me scale.,"from an astrophysical and observational point of view is the most relevant quantity, depends on the temperature through the thermalized time scale."
The number densitv of pairs can be accurately evaluated by using the Selwinger formalism and by taking into account the inhomogeneities in (he electric field distribution.,The number density of pairs can be accurately evaluated by using the Schwinger formalism and by taking into account the inhomogeneities in the electric field distribution.
 However. the boundary effects also induce large quantitative and qualitative deviations of the particle production rate Irom what one deduces with the Schwinger formula and its generalization lor the inhomogeneous electric fielcl of the electrosphere.," However, the boundary effects also induce large quantitative and qualitative deviations of the particle production rate from what one deduces with the Schwinger formula and its generalization for the inhomogeneous electric field of the electrosphere."
 Due to all these ellects. we estimate that al hish temperatures the energy flux due to €=e pars production could be lower than in the initial proposal of Usov(1998a.b).," Due to all these effects, we estimate that at high temperatures the energy flux due to $e^{-}-e^{+}$ pairs production could be lower than in the initial proposal of \citet{Us98a,Us98b}."
. llowever. this flux could still be the main observational signature of a quark star.," However, this flux could still be the main observational signature of a quark star."
 On the other haud. the presence of a strong magnetic field at (he quark star surface may significantly enhance (he electron-positron Εις.," On the other hand, the presence of a strong magnetic field at the quark star surface may significantly enhance the electron-positron flux."
 The possible astrophysical and observational implications of Che direct pair production effect will be considered in a future publication., The possible astrophysical and observational implications of the direct pair production effect will be considered in a future publication.
 The authors would like to thank (he anonymous releree for verv helpful comments ancl suggestions., The authors would like to thank the anonymous referee for very helpful comments and suggestions.
 This work is supported by a RGC grant of the government of the Hong Ixong SAR., This work is supported by a RGC grant of the government of the Hong Kong SAR.
general. we only consider stars in narrow spectral type bins (or effective temperature) to avoid contamination from the intrinsic variation in magnetic activity with stellar mass.,"general, we only consider stars in narrow spectral type bins (or effective temperature) to avoid contamination from the intrinsic variation in magnetic activity with stellar mass."
 The age range «0.07 Gyr is covered by stars in the 1026002. [C2391]. and α PPersei clusters.," The age range $<$ 0.07 Gyr is covered by stars in the IC2602, IC2391, and $\alpha$ Persei clusters."
 We should point out that these clusters are younger than 0.07 Gyr. (0.03. 0.03. and 0.05 Gyr. respectively). but since saturation is present until 0.07 Gyr or longer in the case of K and M stars(22).. considering data with different ages does not affect the value of logLx.," We should point out that these clusters are younger than 0.07 Gyr, (0.03, 0.03, and 0.05 Gyr, respectively), but since saturation is present until 0.07 Gyr or longer in the case of K and M stars, considering data with different ages does not affect the value of $\log L\rm_X$."
 For ages above 1 Gyr cluster or moving group membership is not a useful age determination method., For ages above 1 Gyr cluster or moving group membership is not a useful age determination method.
 Other age indicators. such as the use of rotation period. age-activity relations. asteroseismology. or theoretical isochrones. are more useful in this age domain.," Other age indicators, such as the use of rotation period, age-activity relations, asteroseismology, or theoretical isochrones, are more useful in this age domain."
 Some of them have been used to obtain the ages of a few GKM stars older than «1 Gyr., Some of them have been used to obtain the ages of a few GKM stars older than $\sim$ 1 Gyr.
 Only a handful of field stars have reliable ages in this domain., Only a handful of field stars have reliable ages in this domain.
 These stars are a Cen B and Proxima Cen. with ages determined from the isochrone and asteroseismologic age of their close companion a Cen A(?).. HR7703 shows space motions that are typical of a thick disk star. so we can very roughly assume an age of 10 Gyr.," These stars are $\alpha$ Cen B and Proxima Cen, with ages determined from the isochrone and asteroseismologic age of their close companion $\alpha$ Cen A. HR7703 shows space motions that are typical of a thick disk star, so we can very roughly assume an age of 10 Gyr."
 The X-ray luminosities of these stars were determined from the ROSAT database following the calibration in 2..For stars older than ~6 Gyr long-term changes in high-energy emissions are. in general. difficult to distinguish from short-term stellar activity variations.," The X-ray luminosities of these stars were determined from the ROSAT database following the calibration in .For stars older than $\sim$ 6 Gyr long-term changes in high-energy emissions are, in general, difficult to distinguish from short-term stellar activity variations."
 We have put together all the compilled data. and the evolution of logLx with age for three spectral type intervals (G0-5. K0-5. and Μ0-5) is illustrated in Fig. 1..," We have put together all the compilled data, and the evolution of $\log L\rm_X$ with age for three spectral type intervals (G0-5, K0-5, and M0-5) is illustrated in Fig. \ref{lx_age}."
 In the case of G type stars. the plotted values are quite reliable as they come from a thorough analysis of the sample.," In the case of G type stars, the plotted values are quite reliable as they come from a thorough analysis of the sample."
 For K- and M-type stars. however. the plotted values are very crude results corresponding to the few stars described above that just have rough age estimates.," For K- and M-type stars, however, the plotted values are very crude results corresponding to the few stars described above that just have rough age estimates."
 It is likely that the uncertainty of each point is at least of a few tenths of a dex., It is likely that the uncertainty of each point is at least of a few tenths of a dex.
 The figure shows. as expected. that M-type stars stay at saturated activity levels for a longer period of time than G-type stars.," The figure shows, as expected, that M-type stars stay at saturated activity levels for a longer period of time than G-type stars."
 According to these results. solar-like GO-5 stars are at saturated emission levels until ages of - 100 Myr. and their X-ray luminosity decreases rapidly.," According to these results, solar-like G0-5 stars are at saturated emission levels until ages of $\sim$ 100 Myr, and their X-ray luminosity decreases rapidly."
 K-type stars have saturated emission levels for a little longer (~200 Myr) and then also decrease rapidly., K-type stars have saturated emission levels for a little longer $\sim$ 200 Myr) and then also decrease rapidly.
 Finally. MO-MS stars seem to have saturated emission levels up to 0.5 Gyr or more and then decrease in an analogous way to G- and K-type stars.," Finally, M0-M5 stars seem to have saturated emission levels up to 0.5 Gyr or more and then decrease in an analogous way to G- and K-type stars."
 The interval corresponding to ages youger than ~0.7 Gyr is covered well by the used cluster and moving group stars (1C2602. IC2391. Pleiades. a PPersei. Hyades clusters. and the Ursa Majoris moving group). but it is obvious from our preliminary analysis that a more complete sample of older stars is needed to define reliable age-activity relationships.," The interval corresponding to ages youger than $\sim$ 0.7 Gyr is covered well by the used cluster and moving group stars (IC2602, IC2391, Pleiades, $\alpha$ Persei, Hyades clusters, and the Ursa Majoris moving group), but it is obvious from our preliminary analysis that a more complete sample of older stars is needed to define reliable age-activity relationships."
 Furthermore. it is interesting to note that an age of 0.7 Gyr is a key point for our Sun. since it represents the time at which life is supposed to have appeared on the Earth’s surface.," Furthermore, it is interesting to note that an age of 0.7 Gyr is a key point for our Sun, since it represents the time at which life is supposed to have appeared on the Earth's surface."
 Thus. data for older stars of different types will be very important for modelling planetary atmospheres in a regime that can be relevant to potential life on their surface. as happened to our Earth.," Thus, data for older stars of different types will be very important for modelling planetary atmospheres in a regime that can be relevant to potential life on their surface, as happened to our Earth."
 We have developed an age determination method based on the use of wide binaries where one of the components. a WD. is used as a chronometer.," We have developed an age determination method based on the use of wide binaries where one of the components, a WD, is used as a chronometer."
 The members of a wide binary are assumed to have been born simultaneously and with the same chemical composition., The members of a wide binary are assumed to have been born simultaneously and with the same chemical composition.
 Since they are well separated (100-1000 AU). we can assume that no interaction has occurred between them in the past and they have evolved as single stars(2).," Since they are well separated (100-1000 AU), we can assume that no interaction has occurred between them in the past and they have evolved as single stars."
. We are interested in wide binaries composed by a WD and a star with GKM spectral type., We are interested in wide binaries composed by a WD and a star with GKM spectral type.
 The evolution of a WD can be described as a cooling process. which is relatively well understood at present(?).," The evolution of a WD can be described as a cooling process, which is relatively well understood at present."
. The total age of the WD can be expressed as the sum of its cooling time plus the pre-WD lifetime of its progenitor., The total age of the WD can be expressed as the sum of its cooling time plus the pre-WD lifetime of its progenitor.
 Thus. ages can be obtained from an initial-final mass relationship and stellar tracks to account for the pre-WD lifetime.," Thus, ages can be obtained from an initial-final mass relationship and stellar tracks to account for the pre-WD lifetime."
 This procedure ts analogous to that described by?., This procedure is analogous to that described by.
. It is sound to assume that the age of the WD is the same as the that of the low-mass companion. since both members of a wide binary were born simultaneously.," It is sound to assume that the age of the WD is the same as the that of the low-mass companion, since both members of a wide binary were born simultaneously."
 We selected a sample of 30 wide binaries containing a WD and a GKM star (see Table 1))., We selected a sample of 30 wide binaries containing a WD and a GKM star (see Table \ref{sample}) ).
 In the sample we favour WD components classified as a DA. r.e.. with the unique presence of Balmer lines in their spectra.," In the sample we favour WD components classified as a DA, i.e., with the unique presence of Balmer lines in their spectra."
 As we demonstrate in section 6. the fits to the spectral features of these WDs yield realistic values for the atmospheric. parameters (effective temperature. Tar and surface gravity. log e).," As we demonstrate in section 6, the fits to the spectral features of these WDs yield realistic values for the atmospheric parameters (effective temperature, $_{\rm eff}$ and surface gravity, $\log g$ )."
 We have collected a large amount of observational data (photometry and spectroscopy) on the WD components to implement the proposed approach., We have collected a large amount of observational data (photometry and spectroscopy) on the WD components to implement the proposed approach.
 We aim at determining total system ages (1e.. cooling ages plus progenitor lifetime) with precisions of10-2066.," We aim at determining total system ages (i.e., cooling ages plus progenitor lifetime) with precisions of."
.. This is sufficient 1n. our context since magnetic activity is an intrinsically variable phenomenon and any relationship will, This is sufficient in our context since magnetic activity is an intrinsically variable phenomenon and any relationship will
"assemble the lightest SMDIIS — gas accretion triggered by uajor mergers, with a Dovlau-I&olchiu dvnamical friction nerecr timescale — the mereer rate drops by a factor of wo.","assemble the lightest SMBHs – gas accretion triggered by major mergers, with a Boylan-Kolchin dynamical friction merger timescale – the merger rate drops by a factor of two."
 Over a 32 vear LISA observation. we should be able o detect at least TO SMDII inergers with a signal-to-noise ratio greater than 30 that are involved in assembling the ight end of the SMDIT mass spectrum in the Universe.," Over a 3 year LISA observation, we should be able to detect at least 70 SMBH mergers with a signal-to-noise ratio greater than 30 that are involved in assembling the light end of the SMBH mass spectrum in the Universe."
 If we consider longer LISA observation windows. the iuuber of observed sources ducreascs. naturallv. aud for a 10 vear observation. LISA should detect nearly 500 uereers from the assembly of the lightest SMIBUs alone.," If we consider longer LISA observation windows, the number of observed sources increases, naturally, and for a 10 year observation, LISA should detect nearly 500 mergers from the assembly of the lightest SMBHs alone."
 By concentrating on a small cosmological volume. we aave heen able to model the eravitational wave sources iat result from the assembly of a~10°10*AL. black role in a Milkv Way iass halo.," By concentrating on a small cosmological volume, we have been able to model the gravitational wave sources that result from the assembly of a $\sim 10^6 - 10^7 \,{\rm M}_\odot$ black hole in a Milky Way mass halo."
 We have calculated ιο gravitational wave strain for each of the black hole nerecrs In our volume. and have deteriunued which of rose Will be detectable with LISA.," We have calculated the gravitational wave strain for each of the black hole mergers in our volume, and have determined which of those will be detectable with LISA."
 Of the 1500 mergers oei our vole. we uucovered approximately 0001200 uereers detectable with a signal-to-noise ratio ercater ran 5 over a LISA observation span of 3 years.," Of the 1500 mergers in our volume, we uncovered approximately $300-1200$ mergers detectable with a signal-to-noise ratio greater than 5 over a LISA observation span of 3 years."
 We found that the most common class of observable jack hole merger m our vole is between a SMDII and IMDBII (of 200-2000 AL. ) at z0.05., We found that the most common class of observable black hole merger in our volume is between a SMBH and IMBH (of 200-2000 $_\odot$ ) at $<0.05$.
 These IMDBIIS originally resided in small dark matter halos that mereecd with the massive primary halo at high redshift aud aad very Ίος dyuaudcal friction timescales., These IMBHs originally resided in small dark matter halos that merged with the massive primary halo at high redshift and had very long dynamical friction timescales.
 Before the uereecr occurs. the incoming IMDBIT may be observed with he next ecneration of X-ray telescopes as a ULX source with a rate of about  3-7 vr| for 1 € z € 5.," Before the merger occurs, the incoming IMBH may be observed with the next generation of X-ray telescopes as a ULX source with a rate of about $\sim$ 3 - 7 $^{-1}$ for 1 $\leq$ z $\leq$ 5."
 Because of their potential tie to observable ULNs. aud. because lis class of source has a different waveform character han other well kuown gravitational wave sources. such as equal mass niergers. intermediate mass ratio mspirals (INIBIS). or extreme mass ratio iuspirals (EMBIS). we ie nonunally dubbed this class of source as an Ultra Large Inspirals (ULIs).," Because of their potential tie to observable ULXs, and because this class of source has a different waveform character than other well known gravitational wave sources, such as equal mass mergers, intermediate mass ratio inspirals (IMRIs), or extreme mass ratio inspirals (EMRIs), we have nominally dubbed this class of source as an Ultra Large Inspirals (ULIs)."
 The other class was IMDBIT-IMDII nerecrs at z=3.δ., The other class was IMBH-IMBH mergers at $=2-8$.
 Note that we could not resolve the growth of the most nassive dark matter halos with our technique. aud iiss he mergers that arise from the assembly of the most uassive SAIBUs.," Note that we could not resolve the growth of the most massive dark matter halos with our technique, and miss the mergers that arise from the assembly of the most massive SMBHs."
 Therefore. we consider our approach o be complementary to studies like 7? we believe hat the high mass ratio mergers identified here would add to the rates found iu previous studies that focus ou black hole growth iu present-day halos larger thin Q0HAE.," Therefore, we consider our approach to be complementary to studies like \citet{Sesana:05, Sesana:07} – we believe that the high mass ratio mergers identified here would add to the rates found in previous studies that focus on black hole growth in present-day halos larger than $\sim 10^{11} \, {\rm M}_\odot$."
 However. the high mass-ratio channel that we have identified may well dominate the LISA black hole merger events.," However, the high mass-ratio channel that we have identified may well dominate the LISA black hole merger events."
 Tn fact. even our most pessimistic estimate (~70 eveuts) vields a comparable nuuber of LISA sources as Ὁ (—90 events).," In fact, even our most pessimistic estimate $\sim 70$ events) yields a comparable number of LISA sources as \citet{Sesana:05} $\sim 90$ events)."
 The sigual-to-wise ratio liste in Table 1 ASSTLLUCS simple data anavais techuiques., The signal-to-noise ratio listed in Table \ref{tab:snrdist} assumes simple data analysis techniques.
 If we were to ciplov simular sophisticated tools as are |seine developed for the extreme nis raio iuspirals. such as matched filtering. it is possible tha this class of source may probe slightly larger distances fhan EMBIs.," If we were to employ similar sophisticated tools as are being developed for the extreme mass ratio inspirals, such as matched filtering, it is possible that this class of source may probe slightly larger distances than EMRIs."
 Iun any eveut. current data analysis tecliniqtes to extract SMDII siguals from LISA data streams all assunie the binaries have mass ratios close to unitv.," In any event, current data analysis techniques to extract SMBH signals from LISA data streams all assume the binaries have mass ratios close to unity."
 As we have shown here. the more complex waveform struct)re of these ULIS may require a cliffercut data analysis sti:iteey.," As we have shown here, the more complex waveform structure of these ULIs may require a different data analysis strategy."
 Tn our volume. the SMDIT was in place at nearly its final mass at redshift 5. aud it was built by merecrs of Inmdveds of O(102)AE. black holes at z»5 /— these nerecrs were too weak to be detectable by LISA.," In our volume, the SMBH was in place at nearly its final mass at redshift 5, and it was built by mergers of hundreds of $O(10^2) \, {\rm M}_\odot$ black holes at $>5$ – these mergers were too weak to be detectable by LISA."
 Iu act. the first trace of the erowing SMBIT occurs at about redshift 7 for our most aseressivo gas accretion wescription.," In fact, the first trace of the growing SMBH occurs at about redshift 7 for our most aggressive gas accretion prescription."
 This may lave implications for how well LISA observations cau coustrain the carly erowth of hese lightest SMDITs., This may have implications for how well LISA observations can constrain the early growth of these lightest SMBHs.
 Tn this paper. we have neglected gravitational wave recoil. a potentially important mechanisin that iav inhibit black hole erowth.," In this paper, we have neglected gravitational wave recoil, a potentially important mechanism that may inhibit black hole growth."
 Binary black holes stronely radiate linear momentum in the form of eravitational waves dunug the plunec phase of the iuspiral resulting ina “kick” to the new black hole., Binary black holes strongly radiate linear momentum in the form of gravitational waves during the plunge phase of the inspiral – resulting in a “kick” to the new black hole.
 This. iu itself. has long been predicted as a cousequeuce of an asviunetryv in the binary orbit or spin configuration.," This, in itself, has long been predicted as a consequence of an asymmetry in the binary orbit or spin configuration."
 Previous kick velocity estimates. though. were either highly uncertain or sugeested that the resulting eravitational wave recoil velocity was relatively simall. astroplivsically speaking.," Previous kick velocity estimates, though, were either highly uncertain or suggested that the resulting gravitational wave recoil velocity was relatively small, astrophysically speaking."
" Now. recent results indicate the recoil cau drive a eravitational wave kick velocity as fast as ~1000 liisd (e.c.TTTTTTT,."," Now, recent results indicate the recoil can drive a gravitational wave kick velocity as fast as $\sim 4000 \,$ $\kms$ \citep[e.g.][]{Herrmann07a,
Gonzalez:2006md, gonzalez:07b, Baker:2006vn, Koppitz:07kick, Campanelli:2007cg, schnittman:07}."
 Iu realitv. much siuadler values han this maxi may be expected in eas-rich galaxies due to the alignment of the orbital angular moment and spins of both black holes (7).," In reality, much smaller values than this maximum may be expected in gas-rich galaxies due to the alignment of the orbital angular momentum and spins of both black holes \citep{tamara:07spin}."
 Towever. even vpieal kick velocities (~200laus1) are interestingly aree when compared to the escape velocity of typical astrononicalsvstenis — low mass galaxies. as au example. lave an escape velocity of ~200kins5 (eg.2).," However, even typical kick velocities $\sim 200~\kms$ ) are interestingly large when compared to the escape velocity of typical astronomical systems – low mass galaxies, as an example, have an escape velocity of $\sim 200~\kms$ \citep[e.g.][]{KHB:2007recoil}."
 The effect of large kicks combined with low escape velocity roni the ceuters of small dark matter halos at high redshift plavs a major role in suppressing the erowtl of black hole seeds iuto SAIBIT., The effect of large kicks combined with low escape velocity from the centers of small dark matter halos at high redshift plays a major role in suppressing the growth of black hole seeds into SMBH.
 Even the most massive dark matter halo at zz11 can uot retain a black hole that receives > 150 dans| kick (??)..," Even the most massive dark matter halo at $\geq$ 11 can not retain a black hole that receives $\geq$ 150 ${\rm km \,s^{-1}}$ kick \citep{Merritt:2004xa, micic:2006}."
 We have submitted a conrpanion paper that incorporates the effect of recoil velocity on the expected merger rates and the erowtl of a Milkv Wiauass SMDIT. aud fud that if there is no spin alieument mechamisin. then a Pop III seed black hole cin τοσο] 109NE. oulv 20% of the time through merecr- gas accretion.," We have submitted a companion paper that incorporates the effect of recoil velocity on the expected merger rates and the growth of a Milky Way-mass SMBH, and find that if there is no spin alignment mechanism, then a Pop III seed black hole can reach $10^6 \, {\rm M}_\odot$ only $\%$ of the time through merger-driven gas accretion."
 We are exploring the effect of recoil ou the gravitational wave signal in a forthcoming paper., We are exploring the effect of recoil on the gravitational wave signal in a forthcoming paper.
route is more relevant here: see. e.g.. Preston&Sneden 2000.. Mathieu&Geller2009.. Knigge.Leigh&Sills 2009)).,"route is more relevant here; see, e.g., \citealp{preston00}, , \citealp{mathieu09}, \citealp{knigge09}) )."
 Choosing stars in the broad color cut broad cuts —-0.5<g—r«O0 and 0.8«u—g<1.6 gives a sample that is 73:1 blue stragglers toBHB stars for the full &=19.5 sample of Xueetal.(2008).. and blue stragglers for the g=18 subsample of interest in this paper.," Choosing stars in the broad color cut broad cuts $-0.5<g-r<0$ and $0.8<u-g<1.6$ gives a sample that is $\ga$ 3:1 blue stragglers toBHB stars for the full $g \la 19.5$ sample of \citet{xue08}, and blue stragglers for the $g \la 18$ subsample of interest in this paper."
 This ratio is based entirely on the sample of stars with SDSS spectroscopy as reported in Xueetal.(2008)., This ratio is based entirely on the sample of stars with SDSS spectroscopy as reported in \citet{xue08}.
. The algorithm used to select stars for SDSS spectroscopy changed a number of times during the survey. and does not select uniformly in color-color space: these contamination fractions are affected by this selection (we present alternative estimates of contamination fractions derived in other ways later).," The algorithm used to select stars for SDSS spectroscopy changed a number of times during the survey, and does not select uniformly in color-color space; these contamination fractions are affected by this selection (we present alternative estimates of contamination fractions derived in other ways later)."
 Yet. it is clear that such a level of contamination is a significant challenge for quantitative analysis of the BHB star population using photometry alone.," Yet, it is clear that such a level of contamination is a significant challenge for quantitative analysis of the BHB star population using photometry alone."
 Our goal here is to combine the extensive SDSS spectroscopy with SDSS photometry to devise color selection criteria that are significantly improved in their trade-off between BHB purity andsample size (similar in spirit to the “stringent” color cut of Sirkoetal.20041: see also Clewleyetal.2005 and Kinmanetal.2007 for selection of BHB stars from ugrphotometry)., Our goal here is to combine the extensive SDSS spectroscopy with SDSS photometry to devise color selection criteria that are significantly improved in their trade-off between BHB purity andsample size (similar in spirit to the 'stringent' color cut of \citealp{sirko04}; see also \citealp{clewley05} and \citealp{kinman07} for selection of BHB stars from $ugr$photometry).
 Using a calibration sample of BHB star candidates with spectra (Xueetal. 2008.. following the methodology of Sirkoetal. 2004)). we have determined the probability that à BHB candidate with given 4—$ and g—r color is a spectroscopic BHB star: the rest of the candidates are other blue stars. and are presumably mostly blue stragglers.," Using a calibration sample of BHB star candidates with spectra \citealp{xue08}, , following the methodology of \citealp{sirko04}) ), we have determined the probability that a BHB candidate with given $u-g$ and $g-r$ color is a spectroscopic BHB star; the rest of the candidates are other blue stars, and are presumably mostly blue stragglers."
 reffig:selection illustrates this process. showing stars with g<18 (appropriate for BHB stars in the Heliocentric distance range considered here) and with spectra from the SDSS (Xue 2008).," \\ref{fig:selection} illustrates this process, showing stars with $g < 18$ (appropriate for BHB stars in the Heliocentric distance range considered here) and with spectra from the SDSS \citep{xue08}."
. Blue data points show stars that are very likely to be BHB stars on the basis of their spectra (à contamination. of much less than has been argued by Xueetal.2008 and Sirkoetal.2004. at these e<18 limits)., Blue data points show stars that are very likely to be BHB stars on the basis of their spectra (a contamination of much less than has been argued by \citealp{xue08} and \citealp{sirko04} at these $g<18$ limits).
 The thick contour outlines the region of color-color space where the fraction of BHB candidates that are spectroscopically-classified BHB stars is 250% and there were more than 16 stars in a bin of «0.04 mag (to ensure a high enough number of stars to measure the BHB star fraction with some fidelity)., The thick contour outlines the region of color-color space where the fraction of BHB candidates that are spectroscopically-classified BHB stars is $>50$ and there were more than 16 stars in a bin of $\times$ 0.04 mag (to ensure a high enough number of stars to measure the BHB star fraction with some fidelity).
 This region is well-approximated bythe selection region shown in red in reffig:selection:: 0.98ος 1.28. 0.2<e—r«—0.06 and excluding the region with ές—0.98]/0.215)Εςνε (formulated similarly to the ‘stringent’ color cut of Sirkoetal. 2004.. shown in green).," This region is well-approximated bythe selection region shown in red in \\ref{fig:selection}:: $0.98<u-g<1.28$ , $-0.2<g-r<-0.06$ and excluding the region with $([u-g-0.98]/0.215)^2 + ([g-r+0.06]/0.17)^2 < 1$ (formulated similarly to the 'stringent' color cut of \citealp{sirko04}, , shown in green)."
we areuc below. size. can be used to constrain the deceleration parameter. qu.,"we argue below, size, can be used to constrain the deceleration parameter, $q_0$."
 This is a difficult task. due to uucertaiuties in galaxy formation aud evolution and in the dust content of voung ealaxies.," This is a difficult task, due to uncertainties in galaxy formation and evolution and in the dust content of young galaxies."
 At redshifts :l. though. the difficulty is nütisated bv the order-ofanagnitude seusitivitv of the ditfereutial conioviue volume. dV/d:. to the deceleration parameter.1.," At redshifts $z\ga 1$, though, the difficulty is mitigated by the order-of-magnitude sensitivity of the differential comoving volume, $dV/dz$, to the deceleration parameter,."
. The attempt to detect primeval galaxies by means of strong οwission lines. particuarlyLya.. has failed.," The attempt to detect primeval galaxies by means of strong emission lines, particularly, has failed."
 Blind surveys for cnussion-line objects have covered more than 109Mpe? with no detections! We now uudoerstauk that even a minuscule amount of dust can strongly attenuate the pphotous that are resonantly trapped in the cluitting region., Blind surveys for emission-line objects have covered more than $10^6 \mpc^3$ with no \cite{thompson95} We now understand that even a minuscule amount of dust can strongly attenuate the photons that are resonantly trapped in the emitting region.
 Ilieh redshift galaxies are occasionally discovered sereudipitoisly. for cxzunple when they are leused by a foreground cluster? More recently. a possible vich cluster of galaxies was identified at :2.56 in a field in which there are two confirmed quasars of that redshift separated Dy 1/55. aAC a nearby nulcrowawve vackeround decrenent that is plausibly explained as «lue to the Suivaev-Zoe‘dovich effect? The cluster candidates. which await sCTLOSCOPIC confirmatio1. were identified |»* neans of excess Clussion at the redsüfted wwavelenet1 detected by medii-baud imaglug.," High redshift galaxies are occasionally discovered serendipitously, for example when they are lensed by a foreground \cite{frye98} More recently, a possible rich cluster of galaxies was identified at $z=2.56$ in a field in which there are two confirmed quasars of that redshift separated by 5, and a nearby microwave background decrement that is plausibly explained as due to the Sunyaev-Zel'dovich \cite{campos98} The cluster candidates, which await spectroscopic confirmation, were identified by means of excess emission at the redshifted wavelength detected by medium-band imaging."
 Wile such discoveries are of interest even a sinele confirmice ich cluster at lieh recshift could significauth constrain sone cosmological models real uudoersauding conies only with systematic surveys based on a successful detection scheme., While such discoveries are of interest — even a single confirmed rich cluster at high redshift could significantly constrain some cosmological models — real understanding comes only with systematic surveys based on a successful detection scheme.
 Such a mechuisu) has become available ouly in the last CW Vears. wlh the realizaion fat absorption at the Lyman luit aud in the Lwo--forest reeion lead to broad-hand color features which provide reliable photometric redshift estimates The nec for photometricue redshifts becaue acute as the limiting imagine uaenitude was pushed to anter hnmmats tha rare accessible to spectroscopy. urtieulaulv the faint ealaxies of the Inble Deep Field (UDF).," Such a mechanism has become available only in the last few years, with the realization that absorption at the Lyman limit and in the -forest region lead to broad-band color features which provide reliable photometric redshift \cite{steidel92}
 The need for photometric redshifts became acute as the limiting imaging magnitude was pushed to fainter limits than are accessible to spectroscopy, particularly the faint galaxies of the Hubble Deep Field (HDF)."
 The methods or photouxtry and redshift estimates vary 80newhat between ditfereut eroups. nit there is now aereecnieu ou the resuItaut redshifts. aud they are in good agreement with spectroscopic redshifts serere available? Tn fact. the spectroscopy Is Lot easv ¢ther. aud under careful scruΠιν nore spectroscopic redshifts were ound to be in error than photometric Les? The first and forcimost fact to emerge from the imagine of the TDF was the sanall aneular size of the bulk of the galaxies.," The methods for photometry and redshift estimates vary somewhat between different groups, but there is now agreement on the resultant redshifts, and they are in good agreement with spectroscopic redshifts where \cite{hogg98} In fact, the spectroscopy is not easy either, and under careful scrutiny more spectroscopic redshifts were found to be in error than photometric \cite{lanzetta98a}
 The first and foremost fact to emerge from the imaging of the HDF was the small angular size of the bulk of the galaxies."
 The right panel of slows, The right panel of shows
available.,available.
 Nevertheless. their color temperature may be derived from the fux density ratio (ITeuniug et al.," Nevertheless, their color temperature may be derived from the flux density ratio (Henning et al."
 1990)., 1990).
 Taking iuto account the uucertaiutv iu the computed color temiperature. I select those distributions g(T) that produce the observed flux deusity ratio aud whose peak temperature is comparable to the color temperature.," Taking into account the uncertainty in the computed color temperature, I select those distributions $g(T)$ that produce the observed flux density ratio and whose peak temperature is comparable to the color temperature."
 The Rrected family of ο} funcious obviously satisfies two conditions which are not independer (beiug both related to the flux ratio)., The selected family of $g(T)$ functions obviously satisfies two conditions which are not independent (being both related to the flux ratio).
 This fact could in principle affect the reliability of the results., This fact could in principle affect the reliability of the results.
 It turus out. however. hat pairs of id 5 which satisfy the same constraints eive the salle dius lass Within the fiux uncertaiuty.," It turns out, however, that pairs of $\beta$ and $\gamma$ which satisfy the same constraints give the same dust mass within the flux uncertainty."
 Iun order to check the method a simple numerical sinulation has been performed., In order to check the method a simple numerical simulation has been performed.
 I cousidered au artificial ealaxy which is not resolve by IRAS both at 60 and 100 jan with a eiven dust mass aud. then. T evaluated the dust coutel following the preseut method.," I considered an artificial galaxy which is not resolved by IRAS both at 60 and 100 $\mu$ m with a given dust mass and, then, I evaluated the dust content following the present method."
" The galaxy dust mass AM, may be roughly exnated by using the equatiou where s ds the dust gram radius. p is the dust grain density and Ny is the nunber of dust eramus."," The galaxy dust mass $M_d$ may be roughly estimated by using the equation where $r$ is the dust grain radius, $\rho$ is the dust grain density and $N_d$ is the number of dust grains."
 Eq., Eq.
 3 nuples the following approximations., 3 implies the following approximations.
 The dust eraiu radius p characterizes the dust erams which contribute to the thermal enüssion a FIR wavelengths., The dust grain radius $r$ characterizes the dust grains which contribute to the thermal emission at FIR wavelengths.
 Actually. the nucertaimtics in the total dust nass lutroduced by using an average radius mstead of a size distribution. are much simaller than those arising from uncertainties in the dust emissivitv spectral trend aud in the dust temperature distribution (Iwan & Xie 1992).," Actually, the uncertainties in the total dust mass introduced by using an average radius instead of a size distribution, are much smaller than those arising from uncertainties in the dust emissivity spectral trend and in the dust temperature distribution (Kwan $\&$ Xie 1992)."
 Ny is the total nuuber of the dust erains within the observing beam., $N_d$ is the total number of the dust grains within the observing beam.
 Since the ealaxy is not resolved by IRAS. as it is often the case. Eq.," Since the galaxy is not resolved by IRAS, as it is often the case, Eq."
 3 gives the total dust mass of the point source., 3 gives the total dust mass of the point source.
" A dust radius of 0.1 jan aud a dust grain deusitv of 2 @ ein""5m are currently adopted (Hildebraud 1985) aud also used in the preseut article.", A dust radius of 0.1 $\mu$ m and a dust grain density of 3 g $^{-3}$ are currently adopted (Hildebrand 1983) and also used in the present article.
 Concerning the dust eimissivity. the power-law nuuatiou e(A)=eg(Àg/A) is used in the computations. with a=1 aud ey29.38:10.! (IHildebraud 1983).," Concerning the dust emissivity, the power-law mation $\epsilon(\lambda) = \epsilon_0(\lambda_0/\lambda)^\alpha$ is used in the computations, with $\alpha = 1$ and $\epsilon_0 = 9.38\cdot
10^{-4}$ (Hildebrand 1983)."
 By assuming a value for Ny aud the teniperature distribution οἱ). aud accounting for the spectra responses of the TRAS detectors at 60 aud LOO san. the huumositv enütted by the Ny dust eraius. as observe within the filter bandpass. can be estimated.," By assuming a value for $N_d$ and the temperature distribution $g(T)$, and accounting for the spectral responses of the IRAS detectors at 60 and 100 $\mu$ m, the luminosity emitted by the $N_d$ dust grains, as observed within the filter bandpass, can be estimated."
 The fiux ratio and the color temperature of the galaxy are then colpited. while the dust mass of the source is derives fro Vy by Eq.," The flux ratio and the color temperature of the galaxy are then computed, while the dust mass of the source is derived from $N_d$ by Eq."
 ee, 3.
"3 A galaxy with Ίο M, turns out to coutain about Ag21.6:1072 dust exiius."," A galaxy with $M_d$ $10^5$ $_\odot$ turns out to contain about $N_d$ $\cdot 
10^{52}$ dust grains."
 The galaxy distance is assuiue to be LO Mpe., The galaxy distance is assumed to be 10 Mpc.
 By assuiiiug a g(P) with §—5 aud 5=20 the flux ratio turus out to be about 0.5., By assuming a $g(T)$ with $\beta=5$ and $\gamma=20$ the flux ratio turns out to be about 0.5.
 I use the flux ratio as a constraint to select the pairs ο producing a gCD) whose peak temperature is conrparable to the galaxy color temperature of 36 IX. Among these pairs also the pair j—jB5 and >=20 is fouud (i.c. exactly the values adopted or the input g(£})., I use the flux ratio as a constraint to select the pairs $\beta$ $\gamma$ producing a $g(T)$ whose peak temperature is comparable to the galaxy color temperature of 36 K. Among these pairs also the pair $\beta=5$ and $\gamma=20$ is found (i.e. exactly the values adopted for the input $g(T)$ ).
 For the selected pairs of paraiucters I compute the dust mass which ranges between 69.10! AL..., For the selected pairs of parameters I compute the dust mass which ranges between $6-9\cdot10^4$ $_\odot$.
 Therefore. the derived values are in agreement with he dust mass of the galaxy within the fiux uncertainties. lus confirming the reliability of the method.," Therefore, the derived values are in agreement with the dust mass of the galaxy within the flux uncertainties, thus confirming the reliability of the method."
 Futhermore. wousine the sinele temperature model aud taking iuto account the differcut foxiuulae available (see sect.," Futhermore, by using the single temperature model and taking into account the different formulae available (see sect."
" 2) a dust mass of 1:10! ME, is derived.", 2) a dust mass of $4\cdot10^4$ $_\odot$ is derived.
 This result shows that he single temperature model underestimates the dust content., This result shows that the single temperature model underestimates the dust content.
" The same test was performed for different ""artificial ealaxies"" with different tempperratturre distributions obtaining always cousisteut results.", The same test was performed for different “artificial galaxies” with different re distributions obtaining always consistent results.
 The 21 elliptical galaxies Listed in Table 1 have been extracted from the sample by 695. who evaluated their FIR huuinositv and dust mass from FIR data.," The 21 elliptical galaxies listed in Table 1 have been extracted from the sample by GJ95, who evaluated their FIR luminosity and dust mass from FIR data."
 All the ealaxies are classified as E both in RSA aud iu de, All the galaxies are classified as E both in RSA and in de
error.,error.
 For the chosen CFL uuuber. the hich frequency modes have the largest phase errors but they are highly damped.," For the chosen CFL number, the high frequency modes have the largest phase errors but they are highly damped."
 Some of the modes having lagging phase errors are not lighly damped., Some of the modes having lagging phase errors are not highly damped.
 We will subsequently see how this becomes important., We will subsequently see how this becomes important.
 A rigourous test of the 1-D Las-Weudroff scheme and other flux assigunieut schemes we will discuss is he linear advection of a square wave., A rigourous test of the 1-D Lax-Wendroff scheme and other flux assignment schemes we will discuss is the linear advection of a square wave.
 The challenge is ο accurately advect this discontinuous function where he edges müunüc Riemann shock frouts., The challenge is to accurately advect this discontinuous function where the edges mimic Riemann shock fronts.
 In Figure (2)) we show how the Lax-Woendroff scheme docs at advecting the square wave once (dashed line) aud teu ines (dotted line) through a periodic box of LOO erid cells at speede= Land A=0.9., In Figure \ref{fig:lax}) ) we show how the Lax-Wendroff scheme does at advecting the square wave once (dashed line) and ten times (dotted line) through a periodic box of 100 grid cells at speed$v=1$ and $\lambda=0.9$.
 Note that this scheme oxoduces nuuerical oscillations., Note that this scheme produces numerical oscillations.
 Recall that a square wave cal be represented by a sun of Fourier or sine waves., Recall that a square wave can be represented by a sum of Fourier or sine waves.
 These waves will be damped aud disperse when acvectec sing the Lax-Woudroff scheme., These waves will be damped and disperse when advected using the Lax-Wendroff scheme.
 Figure (1)) shows that the modes having lageine phase errors are not damped away., Figure \ref{fig:lwdispersion}) ) shows that the modes having lagging phase errors are not damped away.
 Tence. the Lax-Wendroff scheme is liehly ispersive and the oscillations im Figure (2)) are due to cispersion.," Hence, the Lax-Wendroff scheme is highly dispersive and the oscillations in Figure \ref{fig:lax}) ) are due to dispersion."
 We leave it as an exercise for the reader to advect a sine wave using the Lax-Wendroff scheme., We leave it as an exercise for the reader to advect a sine wave using the Lax-Wendroff scheme.
 Since there is ouly oue frequency mode in this case. there will be no spurious oscillations due to dispersion. but a phase error will be preseut.," Since there is only one frequency mode in this case, there will be no spurious oscillations due to dispersion, but a phase error will be present."
 For a comprehensive discussion on the family of Lax-Wendroff schemes and other centered schemes. see Ilrsch(1990). ancl Laney(1998).," For a comprehensive discussion on the family of Lax-Wendroff schemes and other centered schemes, see \citet{hir90} and \citet{lan98}."
 Upwiud methods take tuto account the plivsical nature of the flow when assigning fiuxes for the discrete solution., Upwind methods take into account the physical nature of the flow when assigning fluxes for the discrete solution.
 This class offlux assigunmieut schemes. whose origin dates back to the work of Courant.Isaason.&Reeves (1952).. has been shown to be excellent at capturing shocks aud also being highly stable.," This class offlux assignment schemes, whose origin dates back to the work of \citet{cou52}, has been shown to be excellent at capturing shocks and also being highly stable."
 We start with a simple first-order upwiud scheme o solve the linear advection equation., We start with a simple first-order upwind scheme to solve the linear advection equation.
 Consider the case where the advection velocity is positive and flow is to the right., Consider the case where the advection velocity is positive and flow is to the right.
" The fux of the plysical quantity « hrough the cell boundary vy,)y/o will originate fro cell o.", The flux of the physical quantity $u$ through the cell boundary $x_{n+1/2}$ will originate from cell $n$.
 The upwind scheme proposes that. to first-order. the fluxes Frπο at cell boundaries be taken rou the cell-ceutered fluxes Fi=ene which is iu the upwind direction.," The upwind scheme proposes that, to first-order, the fluxes $F_{n+1/2}^t$ at cell boundaries be taken from the cell-centered fluxes $F_n^t=vu_n^t$, which is in the upwind direction."
" If the advection velocity 1s negative and flow is to the left. the boundary fluxcs F’112 are taken from the eell-centered fluxes £7,4=eii,yy."," If the advection velocity is negative and flow is to the left, the boundary fluxes $F_{n+1/2}^t$ are taken from the cell-centered fluxes $F_{n+1}^t=vu_{n+1}^t$."
 The first-order upyiud fiux assiguiuent scheme cau be stuunarized as follows: Uulike ceutral difference schemes; wpwincl schemes are explicitly asviunietric.," The first-order upwind flux assignment scheme can be summarized as follows: Unlike central difference schemes, upwind schemes are explicitly asymmetric."
 The CFL condition for the first-order upwiud scheme can be determined from the vou Nema analysis., The CFL condition for the first-order upwind scheme can be determined from the von Neumann analysis.
 We cousider the case of a positive advection velocity., We consider the case of a positive advection velocity.
 After i time steps. the Fourier modes evolve according to where A=ολΑι aud o=2rhAv/N.," After $m$ time steps, the Fourier modes evolve according to where $\lambda\equiv v\Dt/\Dx$ and $\phi=2\pi k\Dx/N$ ."
 The dispersion relation is given by The CFL condition for solving the linear advection equation with this scheme is to have Ax 1. identical," The dispersion relation is given by The CFL condition for solving the linear advection equation with this scheme is to have $\lambda\leq1$ , identical"
15.5pl We examine (he correlation between compact radio quasars (redshilts in the range 2 —0.3—2.2) and the arrival direction of ultrahigh energy cosmic ravs forming clusters.,15.5pt We examine the correlation between compact radio quasars (redshifts in the range $z = 0.3 - 2.2$ ) and the arrival direction of ultrahigh energy cosmic rays forming clusters.
 Our Monte Carlo simulation reveals a statistically significant correlation on the AGASA saniple: the chance probability of this effect being less than 154., Our Monte Carlo simulation reveals a statistically significant correlation on the AGASA sample: the chance probability of this effect being less than $1\%$.
 The implications of {his result on the origin of ultrahigh energy cosmic ravs are discussed., The implications of this result on the origin of ultrahigh energy cosmic rays are discussed.
adio selection remains one of the most efficient wars of finding hieh-redshift ACN.,Radio selection remains one of the most efficient ways of finding high-redshift AGN.
 This approach has the further advantage of being less prone to selection effects than optical selection. since radio enmüsson is unaffected bv either intrinsic or extrinsic absorption due to dust.," This approach has the further advantage of being less prone to selection effects than optical selection, since radio emission is unaffected by either intrinsic or extrinsic absorption due to dust."
 The specific aiu of this work was to find optically bright. radio-selected high-redshift quasars.," The specific aim of this work was to find optically bright, radio-selected high-redshift quasars."
 These cau be used for mubiased studies of damped Για alpha svstenis at ligh redshift and other follow-up studies such as searclies for associated high-redshitt galaxy clusters., These can be used for unbiased studies of damped Lyman alpha systems at high redshift and other follow-up studies such as searches for associated high-redshift galaxy clusters.
 We therefore οσα to cary out a large. svsteiatic survey aimed specifically at 2>L QSOs.," We therefore began to carry out a large, systematic survey aimed specifically at $z>4$ QSOs."
 Our ucthod involves the optical identification of flat-xpectriun radio sources and the spectroscopic follow-up of the red stellar identifications., Our method involves the optical identification of flat-spectrum radio sources and the spectroscopic follow-up of the red stellar identifications.
 This approach exploits the fact that quasars at high redshift have redder optical colours than their low-redshift counterparts due to absorption by iuterveniug IIT (see figure Lin Wook 1995). aud las proved successful at finding high-redshift quasars in the past (Hook et al 1995. 1996. 1998).," This approach exploits the fact that quasars at high redshift have redder optical colours than their low-redshift counterparts due to absorption by intervening HI (see figure 1 in Hook 1995), and has proved successful at finding high-redshift quasars in the past (Hook et al 1995, 1996, 1998)."
 Previous work using we-clefined quasar samples las shown that 21 radio-loud quasars are likely to be rare objects. both because the quasar population as a whole appears to decline at redshifts above ~23 (Ixenncfick et al.," Previous work using well-defined quasar samples has shown that $z>4$ radio-loud quasars are likely to be rare objects, both because the quasar population as a whole appears to decline at redshifts above $\sim 2-3$ (Kennefick et al."
 1996: Tawkins Verou 1996. Schiuidt. Sclineider Cunu 1995: Warren. Howett Osimer 1995) aud because racio-loud quasars represeu oulv about of the full quasar population.," 1996; Hawkins Veron 1996, Schmidt, Schneider Gunn 1995; Warren, Hewett Osmer 1995) and because radio-loud quasars represent only about of the full quasar population."
 Specific studies of the radio-loud quasar populatio- have shown that these objects are indeed rare at 2>, Specific studies of the radio-loud quasar population have shown that these objects are indeed rare at $z>4$.
" Dunlop Peacock (1990) presented strong evidence for a drop in the space density of radio-loud quasars between >=2 and +3 based on radio-solected samples reaching Soscan,=LOOmdy.", Dunlop Peacock (1990) presented strong evidence for a drop in the space density of radio-loud quasars between $z=2$ and $z\approx 3$ based on radio-selected samples reaching $\rm S_{2.7GHz} = 100mJy$.
 More receuth. significant progress has been made towards understanding the evolution of the radio-loud quasar population out to 2~lI aud the potential effects of absorption bv dust. bv the study of a completely ideutified. large area. flat-xpectrum radio sample with Sorccπμν (Shaver et al.," More recently, significant progress has been made towards understanding the evolution of the radio-loud quasar population out to $z\sim 4$ and the potential effects of absorption by dust, by the study of a completely identified, large area, flat-spectrum radio sample with $\rm
S_{2.7GHz}\ge250mJy$ (Shaver et al."
 1996: Wall et al.," 1996; Wall et al.,"
 iu preparation)., in preparation).
 The low numbers of lieh-+veshift quasars found iu these studies demonstrates that there is a distiuct drop-off in the space deusitv of quasars at z23., The low numbers of high-reshift quasars found in these studies demonstrates that there is a distinct drop-off in the space density of quasars at $z>3$.
 Thus for our new survey to be successful it must reach faiuter radio flux density limits than the above surveys (to salple further down the Iuniuositv function). aud cover a senificaut fraction of the sk.," Thus for our new survey to be successful it must reach fainter radio flux density limits than the above surveys (to sample further down the luminosity function), and cover a significant fraction of the sky."
 Tere we preset he firs results of this new survey for high-redshift radio-loud quasars., Here we present the first results of this new survey for high-redshift radio-loud quasars.
 As will be seen in section 2. the survev uses deeper radio aud optical data than previous radio-loud quasar surveys. and covers a very huge area in the Southern sky (~7.500 sq dee. colmparable to that of the plauned 10.000. sq dee of the Sloan survev).," As will be seen in section 2, the survey uses deeper radio and optical data than previous radio-loud quasar surveys, and covers a very large area in the Southern sky $\sim 7,500$ sq deg, comparable to that of the planned 10,000 sq deg of the Sloan survey)."
 Our survey has produced seven 2>I flat-spectrum quasars. one of which was previously known.," Our survey has produced seven $z>4$ flat-spectrum quasars, one of which was previously known."
 The survey complements the survey of Sucllen et al (2001) which coutains four fat-spectrmm +> bin the Northern sky. selected using aa simular method.," The survey complements the survey of Snellen et al (2001) which contains four flat-spectrum $z>4$ in the Northern sky, selected using a a similar method."
 The parent radio sample used iu this study is based ou the Parkes-\OT-NRAO radio survey (PAIN. απ et al 1995 aud references therin). selected at 5GIIz.," The parent radio sample used in this study is based on the Parkes-MIT-NRAO radio survey (PMN, Griffith et al 1995 and references therin), selected at 5GHz."
 The data cover the southern sky with à<107 to a fux density uit of 20-7210Jx depending on declination., The data cover the southern sky with $\delta < 10^\circ$ to a flux density limit of 20-72mJy depending on declination.
 No additional fiux density lamit was applied when carrving out the survey. but note that when cousideriug the statistics of the final quasar sanple. we consider sources with S2 T21uJw.," No additional flux density limit was applied when carrying out the survey, but note that when considering the statistics of the final quasar sample, we consider sources with $\rm S\ge 72mJy$ ,"
by cooling.,by cooling.
 Only the rim of the cloud becomes diluted and so susceptible to evaporation., Only the rim of the cloud becomes diluted and so susceptible to evaporation.
 Thus. the mass loss (2%)) is reduced by more than a factor of 15 with respect to the analytical approach (~2305€ )).," Thus, the mass loss ) is reduced by more than a factor of 15 with respect to the analytical approach $\sim 30$ )."
 In this case the fate of the cloud depends only on small-scale. not on large-scale dynamics.," In this case the fate of the cloud depends only on small-scale, not on large-scale dynamics."
 These results cannot be extrapolated to homogeneous clouds (model K). where already a small additional heat input is sufficient to unbind the whole cloud.," These results cannot be extrapolated to homogeneous clouds (model K), where already a small additional heat input is sufficient to unbind the whole cloud."
 Although KH instability is also suppressed in this case. the cloud becomes elongated by the Bernoulli effect. which ts very efficient because of the flat gravitational potential.," Although KH instability is also suppressed in this case, the cloud becomes elongated by the Bernoulli effect, which is very efficient because of the flat gravitational potential."
 It is this large-scale phenomenon that peels off a large fragment of the cloud., It is this large-scale phenomenon that peels off a large fragment of the cloud.
 Evaporation due to heat conduction is negligible in comparison to the mass loss due to the large-scale effects., Evaporation due to heat conduction is negligible in comparison to the mass loss due to the large-scale effects.
 Heat conduction acts instead as an agent to extend the survival time of the cloud by the suppression of KH instability and therefore the reduction of the mass-loss 1 comparison to the non-conductive case., Heat conduction acts instead as an agent to extend the survival time of the cloud by the suppression of KH instability and therefore the reduction of the mass-loss in comparison to the non-conductive case.
" While the analytical approach of model K leads to cloud destruction by evaporatio within almost 6 τι, (Fig.", While the analytical approach of model K leads to cloud destruction by evaporation within almost 6 $\tau_{\mbox{\tiny dyn}}$ (Fig.
" 9aa). 1n the numerical simulatio the dynamies stabilize the cloud and heat conduction extends the destruction time to more than I] 7,,."," \ref{f13}a a), in the numerical simulation the dynamics stabilize the cloud and heat conduction extends the destruction time to more than 11 $\tau_{\mbox{\tiny dyn}}$."
 After the loss of ¢£5 large fragment with of the cloud mass. the total disruptio of the cloud is most likely a consequence of the dynamics of the streaming ISM.," After the loss of a large fragment with of the cloud mass, the total disruption of the cloud is most likely a consequence of the dynamics of the streaming ISM."
 It is therefore not surprising that the analytical evaporation rates of CM77 are incompatible with our calculations., It is therefore not surprising that the analytical evaporation rates of CM77 are incompatible with our calculations.
 Comparing the CM77 evaporation rates with computed model;uw for the static case (Paper [)) reveal that condensation may occur for large clouds in temperature regimes where CM77 also predict evaporation., Comparing the \cite{cm77} evaporation rates with computed models for the static case \cite{vh05}) ) reveal that condensation may occur for large clouds in temperature regimes where \cite{cm77} also predict evaporation.
 The three presented models can only give a first insight into the fundamental importance of dynamies and. additional heat conduction in the evolution of the ISM phases., The three presented models can only give a first insight into the fundamental importance of dynamics and additional heat conduction in the evolution of the ISM phases.
 Further investigations are necessary to determine the dependence of evaporation and condensation or the physical state of the phases., Further investigations are necessary to determine the dependence of evaporation and condensation on the physical state of the phases.
 This is necessary for the simulation of galaxy evolution (see e.g. Samland et al., This is necessary for the simulation of galaxy evolution (see e.g. Samland et al.
 1997. Hensler 1999).," 1997, Hensler 1999)."
 The effect of magnetic fields or heat conduction must be discussed. especially when it is neglected. as 1n our models.," The effect of magnetic fields on heat conduction must be discussed, especially when it is neglected, as in our models."
 When electrons of the HIM enter an interstellar cloud. they transfer their energy to the clouc by collisions with. mainly neutral HI atoms of density ny and collisional cross section qui., When electrons of the HIM enter an interstellar cloud they transfer their energy to the cloud by collisions with mainly neutral HI atoms of density $_{\rm HI}$ and collisional cross section $_{\rm HI}$.
 Their collisional mean free path Αι jji/em = (nir: quu) is about 1075 em.," Their collisional mean free path $\lambda_{c,{\rm HI}}$ /cm = $_{\rm HI} \cdot$ $_{\rm HI})^{-1}$ is about $^{16}$ cm."
" On the other hand. magnetic fields force charged particles with mass m, moving with velocity à,jc) perpendicular to the 7 field vectors due to the Lorentz force to gyrate with the so-called (Larmor) radius ay."," On the other hand, magnetic fields force charged particles with mass $_e$ moving with velocity $u_{e,ICM}$ perpendicular to the $B$ field vectors due to the Lorentz force to gyrate with the so-called (Larmor) radius $_{L,e}$."
" For electrons this reads as Since we ""Wdassumeo in our modelling that at the clouds’ surface. both gas phases achieve pressure equilibrium we can also set the magnetic pressure P,,,, = D?/sz in equilibrium."," For electrons this reads as Since we assume in our modelling that at the clouds' surface, both gas phases achieve pressure equilibrium we can also set the magnetic pressure $_{mag}$ = $B^2/8\pi$ in equilibrium."
" Replacing 1,ULaea by the ICM temperature Tye:a;. it cancels with that of the pressure and one gets 107n12 ILJIC'M- "," Replacing $m_e\cdot u_{e,ICM}^2$ by the ICM temperature $_{ICM}$, it cancels with that of the pressure and one gets $_{L,e} = 4.46\cdot 10^5 {\rm n}^{-1/2}_{{\rm H},ICM}$ ."
The ratio of mean free path to electron Larmor radius then is For weaker B fields than in σας pressure equilibrium the ratio can become much larger. 1.9. the free electron motion is less hampered.," The ratio of mean free path to electron Larmor radius then is For weaker B fields than in gas pressure equilibrium the ratio can become much larger, i.e. the free electron motion is less hampered."
 But this only holds for the fraction of electrons moving perpendicular to the B vector. while electrons moving parallel to the B field are unaffected. Malyshkin&Kulsrud(2001).," But this only holds for the fraction of electrons moving perpendicular to the B vector, while electrons moving parallel to the B field are unaffected. \cite{mk01},"
. however. showed that random magnetic fields also reduce the diffusivity of electrons traveling along the B field lines.," however, showed that random magnetic fields also reduce the diffusivity of electrons traveling along the B field lines."
" This reduction of free electron motion and thus heat conduction by magnetic fields can also be parametrized by 9, as in eq.5..", This reduction of free electron motion and thus heat conduction by magnetic fields can also be parametrized by $\Phi_s$ as in \ref{qsat}. .
algorithm to compute AL=ΕΣ)—Z](qu.po) using equation ,"algorithm to compute $\Delta I \equiv
\left[\left(F^*_t \mathcal{I}\right) - \mathcal{I}\right]\left(q_0,
p_0 \right)$ using equation ."
(40).. AL is a two form. and in a two-dimensional phase space it can be writtenAJ=f(/)dqydpy. where f is a scalar.," $\Delta I$ is a two form, and in a two-dimensional phase space it can be written$\Delta I = f(t) dq_0 \wedge dp_0$, where $f$ is a scalar."
 The value of / is plotted versus / for the three algorithms in Figure 5.., The value of $f$ is plotted versus $t$ for the three algorithms in Figure \ref{OneDSymplecticityError}.
 It is clear that the ordinary. adaptive stepsize integrator does not conserve the sviplectie form while the block-power-ol-Gwo ancl constant stepsize integrators do. as expected from Section 4.," It is clear that the ordinary adaptive stepsize integrator does not conserve the symplectic form while the block-power-of-two and constant stepsize integrators do, as expected from Section \ref{AdaptiveTimesteps}."
 As explained in Section 4.. though the energy error performance of the three algorithuns is comparable. only the constant-timestep and block-power-ol-two schemes are sviplectic. ancl only the block-power-ol-two scheme is svimplectie allows [or adaptive timesteps.," As explained in Section \ref{AdaptiveTimesteps}, though the energy error performance of the three algorithms is comparable, only the constant-timestep and block-power-of-two schemes are symplectic, and only the block-power-of-two scheme is symplectic allows for adaptive timesteps."
 In (his subsection we report on (he application of the individual ancl adaptive tüinmestep integration algorithm described in this paper to simulations of manuy-bocly svstems., In this subsection we report on the application of the individual and adaptive timestep integration algorithm described in this paper to simulations of many-body systems.
" We choose to use the so-called ""standard units: units in whieh (he total svstem mass A=1. G=1. andl the total energv £=—1/4 (?).."," We choose to use the so-called “standard units”: units in which the total system mass $M = 1$, $G
= 1$, and the total energy $E = - 1/4$ \citep{Heggie1986}."
 In stancarel units. the virial radius of (he system is 1.," In standard units, the virial radius of the system is 1."
 Our initial conditions are a randomly samplecl Plummer model shifted into a coordinate svstem where the center of mass is at the origin and the (otal linear momentum is zero., Our initial conditions are a randomly sampled Plummer model shifted into a coordinate system where the center of mass is at the origin and the total linear momentum is zero.
 Because our code makes no special provision for close encounters. we used a softened eravitational wilh e=4/N. where N is the number of bodies in a particular simulation.," Because our code makes no special provision for close encounters, we used a softened gravitational with $\epsilon = 4/N$, where $N$ is the number of bodies in a particular simulation."
 This is a standard technique in codes which do not carefully regularize (he singular two-body potential (?).., This is a standard technique in codes which do not carefully regularize the singular two-body potential \citep{Aarseth2001}.
 An N-body system with NV bodies has a 6[N-dimensional phase space., An $N$ -body system with $N$ bodies has a $6N$ -dimensional phase space.
 Given an integration mapping. with Jacobian matrix J. conservation of the Poincaré integral invariant by the mapping implies that where 5 is (he “svimplectic unit”:," Given an integration mapping, with Jacobian matrix$J$ , conservation of the Poincaré integral invariant by the mapping implies that where $S$ is the “symplectic unit”:"
Asstuning this is (rue for other stars. i dgpcd; after a flare. it is also true before a flare.,"Assuming this is true for other stars, if $\delta_{SP} \sim d_{i}$ after a flare, it is also true before a flare."
 The data analyzed in this paper pertain to [lares in sun-like stars. but the underlving cdenamices of reconnection is general.," The data analyzed in this paper pertain to flares in sun-like stars, but the underlying dynamics of reconnection is general."
 Our model applies equally well to mucro- and nano-fares in the quiet corona., Our model applies equally well to micro- and nano-flares in the quiet corona.
 Using values for the quiet sun of T& 1 MIX. ~10? em7. BDe 5 G. and LLOM em. we lind 9sp~770 cm and d;~720 em. in agreement with the model.," Using values for the quiet sun of $T \sim$ 1 MK, $n \sim 10^{9}$ ${\rm
cm}^{-3}$ , $B \sim $ 5 G, and $L \sim 10^{10}$ cm, we find $\delta_{SP} \sim 770$ cm and $d_{i} \sim 720$ cm, in agreement with the model."
 The present result may have important implications for sell-organized criticality (SOC) models of the solar corona., The present result may have important implications for self-organized criticality (SOC) models of the solar corona.
 SOC occurs in driven. dissipative svstems when the svstem is driven (o a eritical state where it undergoes a major reconfiguration (Bakοἱal.1937)..," SOC occurs in driven, dissipative systems when the system is driven to a critical state where it undergoes a major reconfiguration \citep{Bak87}. ."
 SOC leads to power law statistics. which encouraged Lu&ILaanilton(1991) to propose the corona undergoes SOC.," SOC leads to power law statistics, which encouraged \citet{Lu91} to propose the corona undergoes SOC."
 Subsequent studies of SOC in the corona exist (Luetal.1993:Vlahos1995:Longcope&Noonan2000:Islikeretal. 2001).. but a firm physical foundation of the mechanism for sel[-driving aud (he physical condition setting the critical state is often traded for the ease of performing cellular automaton simulations [see Charbonneauetal.(2001) for a review].," Subsequent studies of SOC in the corona exist \citep{Lu93,Vlahos95,Longcope00,Isliker01}, but a firm physical foundation of the mechanism for self-driving and the physical condition setting the critical state is often traded for the ease of performing cellular automaton simulations [see \citet{Charbonneau01}
 for a review]."
 The present result provides a physical mechanism for sell-driving (embedded Sweet- reconnection) and the critical state (marginal collisionalitv). which may provide an avenue lor developing «quantitative predictions of SOC to compare wilh coronal observations.," The present result provides a physical mechanism for self-driving (embedded Sweet-Parker reconnection) and the critical state (marginal collisionality), which may provide an avenue for developing quantitative predictions of SOC to compare with coronal observations."
 An alternate mechanism (Uzdenskyv.2006.2007a.b) for heating the solar corona uses a change in densitv to achieve sell-regulation.," An alternate mechanism \citep{Uzdensky06,Uzdensky07,Uzdensky07b} for heating the solar corona uses a change in density to achieve self-regulation."
 After an eruption. chromospheric evaporation increases (he coronal density. decreasing the ion gvroradius [eq. (2))]," After an eruption, chromospheric evaporation increases the coronal density, decreasing the ion gyroradius [eq. \ref{didef}) )]"
 and making subsequent eruplions more difficult., and making subsequent eruptions more difficult.
 The extent to which Uzdenskvs and our mechanisms regulate coronal heating is an open question., The extent to which Uzdensky's and our mechanisms regulate coronal heating is an open question.
 The present model assumes that Sweel-Parker scaling is appropriate for (hin current sheets of large extent., The present model assumes that Sweet-Parker scaling is appropriate for thin current sheets of large extent.
 Long current sheets are known to [ragment due to secondary instabilities. but the effect of this on the reconnection rate is unknown.," Long current sheets are known to fragment due to secondary instabilities, but the effect of this on the reconnection rate is unknown."
 Verification of the present model would entail testing whether Sweet-Parker reconnection in extended current sheets remains much slower than Lall reconnection. [, Verification of the present model would entail testing whether Sweet-Parker reconnection in extended current sheets remains much slower than Hall reconnection. [
See Uzdensky(2007b) for further discussion of this point as well as other future research directions.|,See \citet{Uzdensky07b} for further discussion of this point as well as other future research directions.]
 The authors thank J. F. Drake. A. IxIimas. E. Ott. ο. Owocki. P. So and D. Uzdensky for helpfulconversations.," The authors thank J. F. Drake, A. Klimas, E. Ott, S. Owocki, P. So and D. Uzdensky for helpfulconversations."
 This work was supportedin part bv the Delaware Space Grant., This work was supportedin part by the Delaware Space Grant.
the motion of charges in a plasma under solar conditions in order to investigate dynamic screening.,the motion of charges in a plasma under solar conditions in order to investigate dynamic screening.
 The advantage of the molecular-dynamics method is that it does not assume a mean field., The advantage of the molecular-dynamics method is that it does not assume a mean field.
" Nor does it assume a long-time average potential for the scattering of any two charges, which 1s necessary in the statistical way to solve Poisson's equation to obtain the mean potential in a plasma."," Nor does it assume a long-time average potential for the scattering of any two charges, which is necessary in the statistical way to solve Poisson's equation to obtain the mean potential in a plasma."
 Shaviv and Shaviv attribute the differences between their simulations and Salpeter's theory to dynamic effects., Shaviv and Shaviv attribute the differences between their simulations and Salpeter's theory to dynamic effects.
" Since their claims have been met with skepticism, we have conducted independent simulations to confirm the existence of dynamic effects."," Since their claims have been met with skepticism, we have conducted independent molecular-dynamics simulations to confirm the existence of dynamic effects."
 Our numerical simulation consisted of a 3-dimensional box with 1000 particles (half protons and half electrons) interacting via the Coulomb potential., Our numerical simulation consisted of a 3-dimensional box with 1000 particles (half protons and half electrons) interacting via the Coulomb potential.
 The temperature and density of the solar core (T=1.6x10’K. p= 1.6x10°kg/m*) provided the velocity distributions and inter-particle spacing.," The temperature and density of the solar core $T=1.6 x 10^7 
\rm{K}$, $\rho = 1.6 x 10^5 \rm{kg/m^3}$ ) provided the velocity distributions and inter-particle spacing."
 We applied periodic boundary conditions and the minimum-image convention., We applied periodic boundary conditions and the minimum-image convention.
" In order to deal with the long-range nature of the Coulomb potential, we implemented a cut-off radius."," In order to deal with the long-range nature of the Coulomb potential, we implemented a cut-off radius."
" With this method, particles separated by a greater distance than the cut-off radius were not included in the potential sums as explained in section 2.1.."," With this method, particles separated by a greater distance than the cut-off radius were not included in the potential sums as explained in section \ref{sect:long-range}."
 Quantum effects were included through the use of effective potentials as desribed in section 2.2.., Quantum effects were included through the use of effective potentials as desribed in section \ref{sect:QM_effects}.
" Detailed analysis of the use of a cutoff radius and effective quantum potential as well as tests of our simulations can be found elsewhere al.2006,2007;Mussack 2007).."," Detailed analysis of the use of a cutoff radius and effective quantum potential as well as tests of our simulations can be found elsewhere \citep{Mao_2004,Mao_thesis,
Mussack_2006, Mussack_2007, Mussack_thesis}."
 The long-range nature of the Coulomb potential introduces additional challenges to a simulation., The long-range nature of the Coulomb potential introduces additional challenges to a molecular-dynamics simulation.
" For many of the materials commonly modeled using molecular dynamics, short-range potentials such as the Lennard-Jones potential of a simple fluid"," For many of the materials commonly modeled using molecular dynamics, short-range potentials such as the Lennard-Jones potential of a simple fluid"
halo area to LOO times that of the PSF (Baeanoffotal. 2003).,halo area to 100 times that of the PSF \citep{bag03}.
. The best-fit spectral paraiueters and derived huninosities for the two detected sources are listed iu Table 3.. and the spectra are displaved in Figure L.," The best-fit spectral parameters and derived luminosities for the two detected sources are listed in Table \ref{tab:spec}, and the spectra are displayed in Figure \ref{fig:xspec}."
 Finally. massive stars. particularly those mo binaries. are often variable X-ray sources.," Finally, massive stars, particularly those in binaries, are often variable X-ray sources."
 Therefore. we examined whether N-vav liebt curves for II2 aud wwere consistent with a constaut nean flux using IXoleomiorv-Siirnov (K-S) aud 4? tests.," Therefore, we examined whether X-ray light curves for H2 and were consistent with a constant mean flux using Kolgomorv-Smirnov (K-S) and $\chi^2$ tests."
 The pliotou flix from increased by a factor of 3 from (3.1£0.2)&109 bbetween 1999 aud 2003 to (7840.7)«10© ddiwing 2001 July, The photon flux from increased by a factor of 3 from $(3.1\pm0.2)\times10^{-6}$ between 1999 and 2003 to $(7.8\pm0.7)\times10^{-6}$ during 2004 July.
 The probability that this increase in flux resulted from a coustaut flux was <2«410* under a 47 test., The probability that this increase in flux resulted from a constant flux was $<2\times10^{-7}$ under a $\chi^2$ test.
 We fiud no evidence that the larducss of the spectrum varied caving the outburst., We find no evidence that the hardness of the spectrum varied during the outburst.
 We defined a harducss ratio (ios)/(hh|s) using s as the πο of counts in the 0.53.0 keV. band. aud 7 as the counts in the 3.08.0 keV baud.," We defined a hardness ratio $(h-s)/(h+s)$ using $s$ as the number of counts in the 0.5–3.0 keV band, and $h$ as the counts in the 3.0–8.0 keV band."
" The average hardness from 1999 through 2003 was 0.02+0.06. while that durius 2001 July was — 0O,012:0.09. so there is no evidence for spectral variations coincident with the merease iu fiux."," The average hardness from 1999 through 2003 was $\pm$ 0.06, while that during 2004 July was $-$ $\pm$ 0.09, so there is no evidence for spectral variations coincident with the increase in flux."
 Ou shorter time scales. there appears to have been a flare near or before the start of the observation ou 2001 July 6. which decaved with a time scale of z1 lr.," On shorter time scales, there appears to have been a flare near or before the start of the observation on 2004 July 6, which decayed with a time scale of $\approx$ 1 hr."
 During that observation. the chance that the data were produced by a coustaut flux was <0.3% according to a I&-8 test.," During that observation, the chance that the data were produced by a constant flux was $<$ according to a K-S test."
 The X-rav lieht curve from I2 has a z:304. chance of being produced by a coustaut mean flux uuder both tests., The X-ray light curve from H2 has a $\approx$ chance of being produced by a constant mean flux under both tests.
 ILowever. because the count rate from II2 is lower than that frou1516.1.. we cannot exclude the liypothesis that a flux variation with a factor of <3 amplitude occurred between 1999 and 2001.," However, because the count rate from H2 is lower than that from, we cannot exclude the hypothesis that a flux variation with a factor of $\la$ 3 amplitude occurred between 1999 and 2004."
 After realizing that M2 had been previously identified ax a voung star. we returued to previous searches for such stars to determine if others were also strong X-ray sources.," After realizing that H2 had been previously identified as a young star, we returned to previous searches for such stars to determine if others were also strong X-ray sources."
 For the purposes of this paper. we are interested in relatively isolated so we searched for N-ray counterparts to the four enission-liue stars (besides ΤΠ) iu Coteraetal.(1999).," For the purposes of this paper, we are interested in relatively isolated so we searched for X-ray counterparts to the four emission-line stars (besides H2) in \citet{cot99}."
.. We first compared the locations and JN magnitudes ofthe Coteraetal.(1999) stars to the 2ATASS catalog in order to refine their locations., We first compared the locations and $JHK$ magnitudes of the \citet{cot99} stars to the 2MASS catalog in order to refine their locations.
 We found that the published positions of the stars disaereed, We found that the published positions of the stars disagreed
nieasure energies of neutrinos above 10!* eV by IceCube.,measure energies of neutrinos above $10^{17}$ eV by IceCube.
 Figure 4 shows the detectable nunmber of spectra of neutrinos in this case., Figure 4 shows the detectable number of spectra of neutrinos in this case.
" The vertical axis e,:N4(e,) roughly. corresponds to the detectable number in each energy range.", The vertical axis $\epsilon_\nu N_{\rm d}(\epsilon_\nu)$ roughly corresponds to the detectable number in each energy range.
 Di this case. the expectation value of neutrinos from kaons is Q.1-1 by a Lkin? detector.," In this case, the expectation value of neutrinos from kaons is 0.1-1 by a $1 {\rm km}^2$ detector."
 However. for very high enerey neutrinos. it may be possible to build detectors with effective volume orders of magnitude larger than Lanz. such as the Extreme Universe Space Observatory (10kn? detector). because the Earth is thick for such neutrinos.," However, for very high energy neutrinos, it may be possible to build detectors with effective volume orders of magnitude larger than $1 {\rm km}^2$, such as the Extreme Universe Space Observatory $10^5 {\rm km}^2$ detector), because the Earth is thick for such neutrinos."
 From the step-Dfunction-like features in (he spectra. we can easily distinguis[un origins of neutrmos.," From the step-function-like features in the spectra, we can easily distinguish origins of neutrinos."
 As shown in Figures 3 and 4. ATj-decay neutrinos are dominant above 107 eV. althoug[un the fIux is too dim to detect on the Earth.," As shown in Figures 3 and 4, $K^0_{\rm L}$ -decay neutrinos are dominant above $10^{18}$ eV, although the flux is too dim to detect on the Earth."
 On the other hand. the contribution of AT-decay in the energy band below LO! eV is not so prominent.," On the other hand, the contribution of $K^0_{\rm L}$ -decay in the energy band below $10^{18}$ eV is not so prominent."
 since we have considered many decaving modes. the production ratio of hieh-enerey," Since we have considered many decaying modes, the production ratio of high-energy"
neutron star X-ray binaries have been found. to have an »proximately. common radio luminosity (Fender Lendry )00).,neutron star X-ray binaries have been found to have an approximately common radio luminosity (Fender Hendry 2000).
 In Fige 1 we summarise our understandingὃν of the relation ob radio emission to transient BIC systems., In Fig 1 we summarise our understanding of the relation of radio emission to transient BHC systems.
 From quiescence (phase 1 in Fig 1). in which systems generally. spend of their time. they are observed. to evolve rapidly within a dav or so to a much brighter state. presumably indicative of an increased accretion rate (e.g. Chen et al.," From quiescence (phase 1 in Fig 1), in which systems generally spend of their time, they are observed to evolve rapidly within a day or so to a much brighter state, presumably indicative of an increased accretion rate (e.g. Chen et al."
 1997)., 1997).
 Subsequently. the systems remain bright ancl only slowly [ade away. often taking months to return to their. pre-outburst levels.," Subsequently the systems remain bright and only slowly fade away, often taking months to return to their pre-outburst levels."
 The X-ray spectrum in this stage is generally either disc- or power-law dominated. corresponding to the canonical high/soft or low/hared X-rav. states respectively.," The X-ray spectrum in this stage is generally either disc- or power-law dominated, corresponding to the canonical high/soft or low/hard X-ray states respectively."
 These two states have dramatically different radio properties (Lender 2000a.0: Fender et al.," These two states have dramatically different radio properties (Fender 2000a,c; Fender et al."
 1999b) from which we infer dilfering acerctionoutflow couplings. as indicated by phases 3a and 3b in Fig 1 respectively.," 1999b) from which we infer differing accretion/outflow couplings, as indicated by phases 3a and 3b in Fig 1 respectively."
 However this is not directly relevant to this study: here we focus on the discrete ejection events which appear to be associated. with the transition hase (stage 2 in Fig 1) at the start of the outburst., However this is not directly relevant to this study; here we focus on the discrete ejection events which appear to be associated with the transition phase (stage 2 in Fig 1) at the start of the outburst.
 For he NS systems. which are all (probably) essentially similar (ic.," For the NS systems, which are all (probably) essentially similar (ie."
 low magnetic. Lick ‘atoll’ sources) there is no clear distinction in. post-peak X-ray. behaviour. and as with the BlICs the optically thin radio emission. observed: during ransicnt outbursts is probably associated with the rapi ejection of matter during the rise to peak of the outburs (stage 2 in Fig 1).," low magnetic field `atoll' sources) there is no clear distinction in post-peak X-ray behaviour, and as with the BHCs the optically thin radio emission observed during transient outbursts is probably associated with the rapid ejection of matter during the rise to peak of the outburst (stage 2 in Fig 1)."
 Whilst in the early vears of X-ray astronomy few racio counterparts to transient systems were reported. in the pas decade many more such svstems have been found. in par due to a greater awareness of the ubiquity and significance of the radio emission.," Whilst in the early years of X-ray astronomy few radio counterparts to transient systems were reported, in the past decade many more such systems have been found, in part due to a greater awareness of the ubiquity and significance of the radio emission."
 In this paper we gather together all the data for BEIC and NS transients for which there are reportec X-ray and radio observations. with the aim of establishing whether there is a clear relation between X-ray. and. racio emission from neutron star and black hole candidate X-ray transicnts both within each class and across the classes.," In this paper we gather together all the data for BHC and NS transients for which there are reported X-ray and radio observations, with the aim of establishing whether there is a clear relation between X-ray and radio emission from neutron star and black hole candidate X-ray transients both within each class and across the classes."
 ‘Table 1 lists the available information for the 22 X-ray transients for which we have some information on the peak X-ray and radio lux. subdivided into 6 neutron-star systems (NSs) and. 16 black-hole candidates CDIICS).," Table 1 lists the available information for the 22 X-ray transients for which we have some information on the peak X-ray and radio flux, subdivided into 6 neutron-star systems (NSs) and 16 black-hole candidates (BHCs)."
 In. addition we tabulate the same information for the two recently cliscoverecl ‘Last transients’ (EVs). NPE J0421|560/€1 Cam and SAN .1819.3-2525/V4641 Ser (the latter of which is likely to be a DIIC Orosz et al.," In addition we tabulate the same information for the two recently discovered `fast transients' (FTs), XTE J0421+560/CI Cam and SAX J1819.3-2525/V4641 Sgr (the latter of which is likely to be a BHC – Orosz et al."
 2000a. b). ancl several unusual sources of recurrent X-ray and radio outbursts. namely (νο X-3. GRS 19151105 and Cir N-1.," 2000a, b), and several unusual sources of recurrent X-ray and radio outbursts, namely Cyg X-3, GRS 1915+105 and Cir X-1."
 As well as information on the source name. outburst date. peak N-rav and radio [luxes and references. we indicate whether the radio emission from the source has been directly. resolved wv radio interferometry.," As well as information on the source name, outburst date, peak X-ray and radio fluxes and references, we indicate whether the radio emission from the source has been directly resolved by radio interferometry."
 This reveals that none of the 7 NS ransient events. three (maybe four) ofthe 16 BLIC events. and all of the FTPs and unusual recurrent svstenis have been directly resolved into jets.," This reveals that none of the 7 NS transient events, three (maybe four) of the 16 BHC events, and all of the FTs and unusual recurrent systems have been directly resolved into jets."
 The data sets for the NS and DEC xus EE svstems are plotted separately in Figs 2(a) and 1(hb) respectively., The data sets for the NS and BHC plus FT systems are plotted separately in Figs 2(a) and 1(b) respectively.
 We briefly cliseuss below how the peak X-ray. aid raclio luxes presented in Table 1 were estimated and compiled., We briefly discuss below how the peak X-ray and radio fluxes presented in Table 1 were estimated and compiled.
 ltadio data are based. on peak observed (lux density at a [requenev of 5 CGllz., Radio data are based on peak observed flux density at a frequency of 5 GHz.
" Where measurements at 5 Cllz are not available. we assume a spectral index of a=Aloes,/ANlogyv=0.5 inorder to estimate the (lux density at 5 Gllz based on other observations."," Where measurements at 5 GHz are not available, we assume a spectral index of $\alpha = \Delta \log S_{\nu} / \Delta \log \nu =
-0.5$ in order to estimate the flux density at 5 GHz based on other observations."
 While most transient outbursts are accompanied by optically thin events for which this will be à good approximation. some svstems produce," While most transient outbursts are accompanied by optically thin events for which this will be a good approximation, some systems produce"
observation which result in a rather wide point spread function at their locations.,observation which result in a rather wide point spread function at their locations.
 Therefore. al least a [fraction of the N-ravs from T2 can possibly be contributed by T1.," Therefore, at least a fraction of the X-rays from T2 can possibly be contributed by T1."
 Furthermore. the tidal radius of 47 Tue is 43° and it is possible that a good fraction of millisecond pulsars are located outside the half-mass radius but within the tidal radius.," Furthermore, the tidal radius of 47 Tuc is 43' and it is possible that a good fraction of millisecond pulsars are located outside the half-mass radius but within the tidal radius."
 Consequently the center of this extended [aint X-ray source T2 may not coincide with the hall-mass radius., Consequently the center of this extended faint X-ray source T2 may not coincide with the half-mass radius.
 With this consideration. the flux measured [roi T2 should be considered as an upper limit.," With this consideration, the flux measured from T2 should be considered as an upper limit."
 The largest model predicted. X-rav energy flix in 3. radius resulling from optical photons is ~(3/17)3.2x10Pergem7s1+.," The largest model predicted X-ray energy flux in 3' radius resulting from optical photons is $\sim
(3'/1^{\circ})^2 3.2 \times 10^{-13}\rm erg~cm^{-2}s^{-1} \sim 10^{-15}\rm erg
~cm^{-2}s^{-1}$."
 However. it is very important to note that the actual emission region of gamma-ray can be much smaller than 1° as this estimate is limited by the angular resolution of LAT.," However, it is very important to note that the actual emission region of gamma-ray can be much smaller than $1^{\circ}$ as this estimate is limited by the angular resolution of LAT."
 Therefore. a dedicated X-ray observation with T2 on-axis can provide important constraint lor the model parameters.," Therefore, a dedicated X-ray observation with T2 on-axis can provide important constraint for the model parameters."
 According Eq. (, According Eq. (
27). the energy fixes at 1 Gllz are 9x10.Litereem7s+ for optical photons. 5x10.Mereem7s+ for IR photons and 5x10Pereem7s! [or relic photons in 1. which correspond to 9 Jv. 5 Jv. and 0.5 Jv respectively.,"27), the energy fluxes at 1 GHz are $9\times 10^{-14} \rm erg ~cm^{-2} s^{-1}$ for optical photons, $5\times 10^{-14} \rm erg ~cm^{-2} s^{-1}$ for IR photons and $5\times 10^{-15} \rm erg ~cm^{-2} s^{-1}$ for relic photons in $1^{\circ}$, which correspond to 9 Jy, 5 Jy, and 0.5 Jy respectively."
 At 400 MIIz the corresponding fhixes will be equal to 18 Jv. 8 Jv and 0.7 Jv.," At 400 MHz the corresponding fluxes will be equal to 18 Jy, 8 Jy and 0.7 Jy."
 The radio (ix from region with diameter 1° is 19 Jv at 408 MIIZ (Haslam et al. (, The radio flux from region with diameter $^\circ$ is 19 Jy at 408 MHz (Haslam et al. (
1982)) and 27 Jv at 1420 MIIZz (Reich et al. (,1982)) and 27 Jy at 1420 MHz (Reich et al. (
2001)).,2001)).
 Lowever. in view of the poor resolution οἱ the instrument. there may have contamination by other sources.," However, in view of the poor resolution of the instrument, there may have contamination by other sources."
 Therefore. the (rue radio fhixes due to the pulsar wind al these frequencies should be lower (han the aforementioned values.," Therefore, the true radio fluxes due to the pulsar wind at these frequencies should be lower than the aforementioned values."
 Ii view of (his. these observed values should only be considered as the upper limits.," In view of this, these observed values should only be considered as the upper limits."
 since the theoretical estimate at 400 MIIz lor the background optical photons (i.e. 18 Jv) is comparable with (he observational limit reported by Laslam et al. (, Since the theoretical estimate at 400 MHz for the background optical photons (i.e. 18 Jy) is comparable with the observational limit reported by Haslam et al. (
1982). there is a hieh probability that the IC: model with the optical photons as the soft photon field may over-predict the radio flux.,"1982), there is a high probability that the IC model with the optical photons as the soft photon field may over-predict the radio flux."
 In section 5.4. we have pointed out that the injected energy {τι Ior optical photons is less than the model predicted value by a factor of ~10 if the," In section 5.4, we have pointed out that the injected energy $E_{inj}$ for optical photons is less than the model predicted value by a factor of $\sim 10$ if the"
profile would be more likely to have higher ionization states. something that can be checked in the existing data.,"profile would be more likely to have higher ionization states, something that can be checked in the existing data."
 lt is also interesting to investigate the distribution of halo masses giving rise to DLAS., It is also interesting to investigate the distribution of halo masses giving rise to DLAS.
 Fig., Fig.
 12. shows the distribution. of circular velocities of the halos containing clises that eive rise to DLAS., \ref{fvc} shows the distribution of circular velocities of the halos containing discs that give rise to DLAS.
 Also shown is the average cross section for DLAS as a function of circular velocity. which agrees fairly well with the results of (1997) (slope =2.94) and Llachnelt. (1999) (slope = 2.5). butnot those of (1999) who incs a much shallower slope of 0.9. ," Also shown is the average cross section for DLAS as a function of circular velocity, which agrees fairly well with the results of \nocite{gard:97}{ (1997) (slope =2.94) and \nocite{hsr:99}{, (1999) (slope = 2.5), butnot those of \nocite{gard:99}{ (1999) who finds a much shallower slope of 0.9. \nocite{hsr:99}{"
determine weir average cross section bw fitting to the observed: Av listribution so we expect. that the relationship between 10 circular velocity of the halo and the Ae of the DLAS mt arise in it must be the same in our modeling and the simulations of ., determine their average cross section by fitting to the observed $\delv$ distribution so we expect that the relationship between the circular velocity of the halo and the $\delv$ of the DLAS that arise in it must be the same in our modeling and the simulations of \nocite{hsr:99}{.
 Phis is in fact the case as can oe seen. by comparing from Fig., This is in fact the case as can be seen by comparing from Fig.
 13. and Fig., \ref{fdvvc} and Fig.
 1 in (1999)., 1 in \nocite{hsr:99}{ (1999).
 This seems to suggest that the very. dilferent approaches of νάνο simulations ancl S.XMs are converging on à common picture for the nature of the DLAS., This seems to suggest that the very different approaches of hydro simulations and SAMs are converging on a common picture for the nature of the DLAS.
 We have explored. the properties of DLAS in semi-analytic models of galaxy formation., We have explored the properties of DLAS in semi-analytic models of galaxy formation.
 These models produce. good agreement with many optical properties of galaxies. at low anc high redshift. and the total mass of cold gas at redshift ~3 is also in reasonable agreement with observations.," These models produce good agreement with many optical properties of galaxies at low and high redshift, and the total mass of cold gas at redshift $\sim 3$ is also in reasonable agreement with observations."
 Lt is therefore interesting to ask whether the kinematic properties. metallicities. and column densities of DLAS in these models are in agreement with observations.," It is therefore interesting to ask whether the kinematic properties, metallicities, and column densities of DLAS in these models are in agreement with observations."
 We investigated the dependences of these properties. on cosmology. the distribution of satellite orbits. and. gaseous disc scale. height. anc found that our results were not sensitive to these assumptions.," We investigated the dependences of these properties on cosmology, the distribution of satellite orbits, and gaseous disc scale height, and found that our results were not sensitive to these assumptions."
 Our results are to our assumptions about the radial distribution of cold gas within galactic discs., Our results are to our assumptions about the radial distribution of cold gas within galactic discs.
 Given that one believes the other components of our model. one can then perhaps learn about the distribution of cold neutral gas at high redshi," Given that one believes the other components of our model, one can then perhaps learn about the distribution of cold neutral gas at high redshift."
 Currently. popular theories of cise formation posit that the racial size of a galactic disc is determined by the initial specific angular momentum of the dark. matter halo in which it forms. and that the cold gas traces the stellar component.," Currently popular theories of disc formation posit that the radial size of a galactic disc is determined by the initial specific angular momentum of the dark matter halo in which it forms, and that the cold gas traces the stellar component."
 Often. the profile of the disc is assumed to have an exponential form.," Often, the profile of the disc is assumed to have an exponential form."
 We investigate several variants of such models. based on ideas in the literature such as (1980). (1998) and (1996).," We investigate several variants of such models, based on ideas in the literature such as \nocite{fe:80}{ (1980), \nocite{mmw:98}{ (1998) and \nocite{kauf:96}{ (1996)."
 We lind that the kinematies of DLAS arising in such mocdoels are in strong conllict with the observations of (1997h. 1998).," We find that the kinematics of DLAS arising in such models are in strong conflict with the observations of \nocite{pw:97,pw:98}{ (1997b, 1998)."
 This is consistent with the previous work of Wolfe.. in which it was shown that if the," This is consistent with the previous work of \nocite{pw:97}{ , in which it was shown that if the"
wawas observed with the VIMOS Integral Field Unit (Ub) as part of a project to map the properties of a sample of 24 ealaxies selected: randomly. from the Sloan Digital Sky Survey (SDSS: York et al.,was observed with the VIMOS Integral Field Unit (IFU) as part of a project to map the properties of a sample of 24 galaxies selected randomly from the Sloan Digital Sky Survey (SDSS; York et al.
 2000)., 2000).
 ALL data were obtained using the medium resolution setup (wavelength range: 5000 - 9000Α.. dispersion. 2.5 AX//pix) covering a field-of-view of 27x27 aresee (0.67 arcsec/spaxel).," All data were obtained using the medium resolution setup (wavelength range: 5000 - 9000, dispersion, 2.5 /pix) covering a field-of-view of 27x27 arcsec (0.67 arcsec/spaxel)."
 We obtained two 30 minute exposures on this galaxy during service mode observations in January 2007 (in seeing conditions of about 1.5 arcsec)., We obtained two 30 minute exposures on this galaxy during service mode observations in January 2007 (in seeing conditions of about 1.5 arcsec).
 X detailed: description of the data reduction will be given in a forthcoming paper (Gerssen et al. 2008).," A detailed description of the data reduction will be given in a forthcoming paper (Gerssen et al, 2008)."
 Briellv.. we usec the ESO VIMOS pipeline to perform the basic reduction steps up to spectrum extraction and wavelength calibration.," Briefly, we used the ESO VIMOS pipeline to perform the basic reduction steps up to spectrum extraction and wavelength calibration."
 The post-processing steps (e.g. throughput correction. flux calibration. anc exposure combination) to create the final data cube Gry A) were performed using custom written IDL scripts.," The post-processing steps (e.g. throughput correction, flux calibration, and exposure combination) to create the final data cube $x$ $y$ $\lambda$ ) were performed using custom written IDL scripts."
 To analyse the emission line data. we independently fi the Ho |NH] group. the OLI] doublet. the SH] double and the LL emission line.," To analyse the emission line data, we independently fit the $\alpha$ +[NII] group, the [OIII] doublet, the [SII] doublet and the $\beta$ emission line."
 Each Line is fit with a single Gaussian. and for cach set of lines the relative position anc widths of cach line are fixed to each other as they trace the same kinematics.," Each line is fit with a single Gaussian, and for each set of lines the relative position and widths of each line are fixed to each other as they trace the same kinematics."
 For example. in a three componen fit to the Ho |NH] emission lines we tie the eentroids aix line widths to the Hla line.," For example, in a three component fit to the $\alpha$ +[NII] emission lines we tie the centroids and line widths to the $\alpha$ line."
 In this case there are six [ree parameters: the amplitudes of the three emission lines. the line centroid and line width. anc à constant continuum level.," In this case there are six free parameters: the amplitudes of the three emission lines, the line centroid and line width, and a constant continuum level."
 We do not include an additional broad. Lla component in the emission line analysis., We do not include an additional broad $\alpha$ component in the emission line analysis.
 Phe Broad Line Region (DLIU in iis only detectable in the Hao line and then only in spectra close to the nucleus. where it is so broad as to have no inlluence on the fit.," The Broad Line Region (BLR) in is only detectable in the $\alpha$ line and then only in spectra close to the nucleus, where it is so broad as to have no influence on the fit."
 The galaxy iis a Sevfert 1.9 at a of 293 Ape., The galaxy is a Seyfert 1.9 at a of 293 Mpc.
 dts. basic properties are listed in Table 1., Its basic properties are listed in Table 1.
 It stands out in our sample because the radial dependence of its strongest emission lines (llo. Ll. NHJ6584. OLLI]J5007) indicates a high ionization state out to large radii. see Figure 1..," It stands out in our sample because the radial dependence of its strongest emission lines $\alpha$, $\beta$, [NII]6584 [OIII]5007) indicates a high ionization state out to large radii, see Figure \ref{f:bpt}."
" In this so-called. DI""T diagram (Baldwin et al.", In this so-called BPT diagram (Baldwin et al.
 1981) we plot results derived: using a synthetic annulus-aperture (2 aresec width) of increasing racius., 1981) we plot results derived using a synthetic annulus-aperture (2 arcsec width) of increasing radius.
" Remarkably. iis located on the AGN ""wing of the DIE diagram out to a radius of at least 9 kpc."," Remarkably, is located on the AGN `wing' of the BPT diagram out to a radius of at least 9 kpc."
 This implies a role for strong ionizing radiation. probably associated with the AGN. on ealaxy-wicle scales.," This implies a role for strong ionizing radiation, probably associated with the AGN, on galaxy-wide scales."
 displays a complex morphological structure that is strongly wavelength: dependent., displays a complex morphological structure that is strongly wavelength dependent.
 A composite colour image derived from our cata cube is shown in panel (a) of Figure 2.., A composite colour image derived from our data cube is shown in panel (a) of Figure \ref{f:mosaic}.
 The two brightest knots. labelled A and D. coincide respectively with the nucleus of the host and that of a nearby. galaxy.," The two brightest knots, labelled A and B, coincide respectively with the nucleus of the host and that of a nearby galaxy."
 Their projected: separation is 11.6 kpe., Their projected separation is 11.6 kpc.
 Hence. it is likely that the two svstems are interacting.," Hence, it is likely that the two systems are interacting."
 A clear manifestation of interaction are the faint. knots visible NE of the host nucleus., A clear manifestation of interaction are the faint knots visible NE of the host nucleus.
 The knot labelled € coincides with the peak of a resolved ultraviolet source (GALEN database)., The knot labelled C coincides with the peak of a resolved ultraviolet source (GALEX database).
 The lla and OLLI] line flux maps (panels d and ο) respectively) do show prominent features in region € as well., The $\alpha$ and [OIII] line flux maps (panels d and g) respectively) do show prominent features in region C as well.
 Interestingly. the Lla peak intensity is somewhat stronger for the oll-centre peak than on the nucleus κο.," Interestingly, the $\alpha$ peak intensity is somewhat stronger for the off-centre peak than on the nucleus itself."
 Phe average line ratios over region € (shown in Figure 1) are consistent with ionization by voung stars., The average line ratios over region C (shown in Figure 1) are consistent with ionization by young stars.
 This area is likely associated with oll-centre star formation (at 10.5 projected kpe from the nucleus) trigecred by interaction with the companion galaxy., This area is likely associated with off-centre star formation (at $\sim 10.8$ projected kpc from the nucleus) triggered by interaction with the companion galaxy.
 As the companion galaxy shows no emission lines we establish its nature using a near infrared. L-band image obtained from the UINIICE Infrared. Deep Sky Survey (UIXIDSS: Lawrence ct al., As the companion galaxy shows no emission lines we establish its nature using a near infrared H-band image obtained from the UKIRT Infrared Deep Sky Survey (UKIDSS; Lawrence et al.
 2007)., 2007).
 The surface brightness profile (derived with ELLIPSE in LRA) is consistent with the light profile of an elliptical galaxy with an elfective radius of Ray~2.5 kpe., The surface brightness profile (derived with ELLIPSE in IRAF) is consistent with the light profile of an elliptical galaxy with an effective radius of $R_{\rm eff} \sim 2.5$ kpc.
 ‘To constrain the systemic velocities of the host. ancl the companion we sum the spectra in our data cube over regions A and D. We use the pixel-fitting method of (Cappellari Emscllen 2004) to fit the summed. specta. with a set of stellar. template spectra observed. with EMMLE on. the NT (convolved to the VIMOS instrumental resolution of 7.8 A))., To constrain the systemic velocities of the host and the companion we sum the spectra in our data cube over regions A and B. We use the pixel-fitting method of (Cappellari Emsellem 2004) to fit the summed specta with a set of stellar template spectra observed with EMMI on the NTT (convolved to the VIMOS instrumental resolution of 7.3 ).
 Phe comparison. shown in panel (b). between the stellar absorption line spectra extracted at locations A aud 3 demonstrates that the companion galaxy is also close in velocity space (Ary.~600 Km ly ," The comparison, shown in panel (b), between the stellar absorption line spectra extracted at locations A and B demonstrates that the companion galaxy is also close in velocity space $\Delta
v_{\rm los} \sim 600$ km $^{-1}$ )."
To examine the result shown in Figure 1 in more detail we can use our data cube to derive the line ratios in each spatial clement individually., To examine the result shown in Figure \ref{f:bpt} in more detail we can use our data cube to derive the line ratios in each spatial element individually.
 The full 2-D line ratio map of OLLJA5007 /LL7 is shown in panel (c) of Figure 2.., The full 2-D line ratio map of $\lambda5007$ $\beta$ is shown in panel (c) of Figure \ref{f:mosaic}.
 Consistent with our azimuthallv averaged result. (Fig., Consistent with our azimuthally averaged result (Fig.
 1). the map shows an extended. region of highly ionizecl gas.," 1), the map shows an extended region of highly ionized gas."
 The observed line ratios require a very strong ionization field. (hard UV. spectrum) ancl are typical for the Narrow Line Region (NLR) of a Sevfert galaxy., The observed line ratios require a very strong ionization field (hard UV spectrum) and are typical for the Narrow Line Region (NLR) of a Seyfert galaxy.
 Such highly ionized, Such highly ionized
"result is obtained for fMerit, which increases by roughly 0.2 dex from no UV to FG UV and changes negligibly when X- are added.","result is obtained for $fM_{\text{crit}}$, which increases by roughly 0.2 dex from no UV to FG UV and changes negligibly when X-rays are added."
" It therefore appears that the effects of the ionizing background and of SN feedback amplify one another in a superlinear fashion, as suggested by and(2010)."," It therefore appears that the effects of the ionizing background and of SN feedback amplify one another in a superlinear fashion, as suggested by and."
". This is also indicated by Figure 2, where the dashed lines correspond to the no-feedback models: in the lowest-mass bin, adding feedback causes a greater relative reduction in baryon-conversion efficiency when an ionizing background is present."," This is also indicated by Figure 2, where the dashed lines correspond to the no-feedback models: in the lowest-mass bin, adding feedback causes a greater relative reduction in baryon-conversion efficiency when an ionizing background is present."
" To make a detailed comparison with(2010), we calculate the baryon-conversions efficiencies for a 0.2-dex-wide bin around their fiducial halo mass of 7x10°Mo."," To make a detailed comparison with, we calculate the baryon-conversions efficiencies for a 0.2-dex-wide bin around their fiducial halo mass of $7\times 10^8 M_{\odot}$."
" For their runs with UV background only, SN feedback only, and both, give efficiencies of 0.77, 0.05, and 0.03, respectively."," For their runs with UV background only, SN feedback only, and both, give efficiencies of 0.77, 0.05, and 0.03, respectively."
" Here, however, the respective models FG UV-F, No UV and FG UV give median values of, respectively, 0.16, 0.47, and 0.04."," Here, however, the respective models FG UV-F, No UV and FG UV give median values of, respectively, 0.16, 0.47, and 0.04."
" So for(2010), the feedback is by far the most important contributor, whereas here the background is somewhat more important."," So for, the feedback is by far the most important contributor, whereas here the background is somewhat more important."
" The stronger effect of the background in this work seems likely to be the effect of timing: turn their background on at z—6, by which time their galaxy has already created 20—30% of its final stellar mass; in our simulations, however, galaxies of this final size have no significant star formation until z~3 (see the next subsection)."," The stronger effect of the background in this work seems likely to be the effect of timing: turn their background on at $z=6$, by which time their galaxy has already created $20-30\%$ of its final stellar mass; in our simulations, however, galaxies of this final size have no significant star formation until $z\sim 3$ (see the next subsection)."
" Therefore our galaxies show no significant change in baryon conversion between Old UV, which turns on at z—6, and FG UV, which turns on at z=10, but we speculate that if the runs were to be repeated with FG UV, the background would have a more significant effect."," Therefore our galaxies show no significant change in baryon conversion between Old UV, which turns on at $z=6$, and FG UV, which turns on at $z=10$, but we speculate that if the runs were to be repeated with FG UV, the background would have a more significant effect."
" The relatively weak effect of feedback for our models, on the other hand, is likely the effect of environment: the relatively dense environment and corresponding IGM pressure on our small halos means that they do not have enough supernova energy to unbind their gas."," The relatively weak effect of feedback for our models, on the other hand, is likely the effect of environment: the relatively dense environment and corresponding IGM pressure on our small halos means that they do not have enough supernova energy to unbind their gas."
" The physical mechanisms governing the interaction between feedback and the background radiation are beyond the scope of this paper, but our results are consistent with a picture where feedback moves gas from the centers of halos to the outer regions, where the background then provides enough energy to unbind the gas altogether."," The physical mechanisms governing the interaction between feedback and the background radiation are beyond the scope of this paper, but our results are consistent with a picture where feedback moves gas from the centers of halos to the outer regions, where the background then provides enough energy to unbind the gas altogether."
" Thus there is no change in the fMerit when X-rays are added in the absence of feedback because the dense central gas has too short a cooling time to be affected (although the smallest halos with the lowest central densities can still be suppressed, increasing o), whereas if it were first pushed out somewhat by feedback the density would be lower, the cooling time longer and X-rays could have an effect."," Thus there is no change in the $fM_{\text{crit}}$ when X-rays are added in the absence of feedback because the dense central gas has too short a cooling time to be affected (although the smallest halos with the lowest central densities can still be suppressed, increasing $\alpha$ ), whereas if it were first pushed out somewhat by feedback the density would be lower, the cooling time longer and X-rays could have an effect."
" Conversely, with feedback and no background (the No UV case), the gas gets pushed outward, but not enough to unbind, so it falls back in and forms more stars."," Conversely, with feedback and no background (the No UV case), the gas gets pushed outward, but not enough to unbind, so it falls back in and forms more stars."
" 'To investigate the origins of the discrepancies in low-mass systems, we examine the simulations at higher redshifts."," To investigate the origins of the discrepancies in low-mass systems, we examine the simulations at higher redshifts."
 Figure 3 shows the stellar mass spectrum of the FG UV simulations from z—4 to the present.," Figure \ref{fig:nhist}
 shows the stellar mass spectrum of the FG UV simulations from $z=4$ to the present."
 The number of low-mass systems peaks at ze:2—3 and then declines as small systems merge into larger ones., The number of low-mass systems peaks at $z\approx 2-3$ and then declines as small systems merge into larger ones.
" This peak is somewhat later than the overall star-formation peak, found in to be at z74, which is precisely the “downsizing” effect first reported byresults)."," This peak is somewhat later than the overall star-formation peak, found in to be at $z\approx 4$, which is precisely the “downsizing” effect first reported by."
". This effect explains why our three UV models show such similar results here: at z=4 the Old UV and FG UV models have not finished Hell reionization, and Old UV has ~3 times lower ionization rate than the other two models, while for z<3 the three models are nearly identical in ionization rate and have all fully reionizedI)."," This effect explains why our three UV models show such similar results here: at $z=4$ the Old UV and FG UV models have not finished HeII reionization, and Old UV has $\sim 3$ times lower ionization rate than the other two models, while for $z\leq3$ the three models are nearly identical in ionization rate and have all fully reionized."
". Thus at the epoch when the smallest halos are forming stars, the gas temperatures in the three models are much more similar than when stars in more dense regions formedI)."," Thus at the epoch when the smallest halos are forming stars, the gas temperatures in the three models are much more similar than when stars in more dense regions formed."
". On the other hand, the models with X- have higher gas temperatures than the UV-only models for all z<4, which is reflected in the increased suppression of small systems."," On the other hand, the models with X-rays have higher gas temperatures than the UV-only models for all $z<4$, which is reflected in the increased suppression of small systems."
" Table 2. compares the total number of identified (stellar) systems in the whole box for the No UV and FG UV cases, as well as the baryon conversion efficiency €x,e, defined in Equation 9, for two bins in halo mass (1055<(Muaio/Mo)«10°° and 101°?<(Mmao/Mz))."," Table \ref{tab:hist} compares the total number of identified (stellar) systems in the whole box for the No UV and FG UV cases, as well as the baryon conversion efficiency $\epsilon_{\star,\text{eff}}$, defined in Equation 9, for two bins in halo mass $10^{8.5}<(M_\text{Halo}/M_{\odot})< 10^{9.5}$ and $10^{10.5}<(M_\text{Halo}/M_{\odot})$ )."
" Although No UV has ~3 times the number of systems as FG UV, they both show the same time dependence: the number of systems reaches a peak at 2S3, followed by mergers reducing the number."," Although No UV has $\sim 3$ times the number of systems as FG UV, they both show the same time dependence: the number of systems reaches a peak at $2\la z \la 3$, followed by mergers reducing the number."
" The stellar mass fraction for high-mass halos is nearly the same for the two models, while for low-mass halos the difference is a factor of 2—6."," The stellar mass fraction for high-mass halos is nearly the same for the two models, while for low-mass halos the difference is a factor of $2-6$."
" Interestingly, in the FGUV case the low-mass stellar mass fraction declines somewhat from z=2 to 0, presumably because closer satellites of the central galaxy, which had been able to form more stars than the average because of their denser environment, are finally accreted and lost (while with No UV all satellites can form many stars)."," Interestingly, in the FGUV case the low-mass stellar mass fraction declines somewhat from $z=2$ to 0, presumably because closer satellites of the central galaxy, which had been able to form more stars than the average because of their denser environment, are finally accreted and lost (while with No UV all satellites can form many stars)."
"Let us assume that AL and AL are tbe WIXB amplitudes far on the left and right hand side [rom (he point τῇ, respectively.","Let us assume that $A^L_\pm $ and $A^R_\pm $ are the WKB amplitudes far on the left and right hand side from the point $\tau^\ast$, respectively."
 It is well-known (Olver1974) that the conservation of the Wronskian leads to the following invariant: which physically corresponds to the conservation of the wave action (see.e.g.. and references therein).," It is well-known \citep{o74} that the conservation of the Wronskian leads to the following invariant: which physically corresponds to the conservation of the wave action (see, \citet{g04} and references therein)."
 If initially A’=0. the refraction and transmission coellicients may be clefined in the usual wav: while from Eq. (," If initially $A_-^L \equiv 0$, the refraction and transmission coefficients may be defined in the usual way: while from Eq. ("
56) we obviouslyhave: The dependence of the transmission coefficient DUR) on A. obtained through (he numerical solution of Eq. (,"56) we obviouslyhave: The dependence of the transmission coefficient $D(R)$ on $R$, obtained through the numerical solution of Eq. ("
52) for different values of the Drunt-Vàiisallà Irequency. N=O(solidline). l(dashedline\andN= 2(dash—doltedline). ispresentedonE ig.,"52) for different values of the Brunt-Väiisällä frequency, $N = 0$ (solid line), $N = 1$ (dashed line) and $N = 2$ (dash-dotted line), is presented on Fig."
 lA, 1.
 heinitialconditionsforlhenumericals handsideoftheresonantarea(Ix.(0) s>1. RH).," The initial conditions for the numerical solution are chosen from the WKB solutions (55) far on the left-hand side of the resonant area $K_z(0) \gg 1,R$ )."
 The analvsis leads to the conclusion that the amplification of the energy density of the perturbations is always finite. but it can become arbitrarily large wilh a proper increase of (he shearing rate.," The analysis leads to the conclusion that the amplification of the energy density of the perturbations is always finite, but it can become arbitrarily large with a proper increase of the shearing rate."
 It can be seen that the velocity shear can cause a strong amplification of the GAW. even for moderate values of the normalized shearing rates.," It can be seen that the velocity shear can cause a strong amplification of the GAW, even for moderate values of the normalized shearing rates."
 According to the munerical study of Eq. (, According to the numerical study of Eq. (
52). the amplitude of the reflected wave exceeds (he amplitude of the incident wave il: This phenomenon. originally discovered by Miles lor acoustic waves (Miles1957)... is called.,"52), the amplitude of the reflected wave exceeds the amplitude of the incident wave if: This phenomenon, originally discovered by Miles for acoustic waves \citep{m57}, is called."
over-reflection.. Bearing in mixl that the enerev of the mode is proportional to the square of its amplitude. we can easily surmise (hat in the course of its non-acliabatic evolution the perturbation energv is increasing.," Bearing in mind that the energy of the mode is proportional to the square of its amplitude, we can easily surmise that in the course of its non-adiabatic evolution the perturbation energy is increasing."
 According to Eqs. (, According to Eqs. (
56) and (53) the ratio of the total 'post-reflection! energy. {ο the initial one equals Re+D>1.,56) and (58) the ratio of the total `post-reflection' energy to the initial one equals $Re+D>1$.
 The energy. increases at the expense of the mean (shear) flow enerev., The energy increases at the expense of the mean (shear) flow energy.
 Consequently. over-rellection represents a cuite efficient mechanism lor translerring (he mean flow energy to perturbations.," Consequently, over-reflection represents a quite efficient mechanism for transferring the mean flow energy to perturbations."
 In the limit R— 0. the reflection coellicient becomes exponentially small with respect to the parameter —1/HR. i.e.. Reexp(—1/R) (Gogoberidzeetal. 2004).," In the limit $R \rightarrow 0 $ , the reflection coefficient becomes exponentially small with respect to the parameter $-1/R$ , i.e., $Re
\sim \exp(-1/R)$ \citep{g04}. ."
In the fourth case (region B). no hysteresis ancl associated orientation could be observed. but the spectrum did harden as the flux increased during the rise and decay. of flux.,"In the fourth case (region B), no hysteresis and associated orientation could be observed, but the spectrum did harden as the flux increased during the rise and decay of flux."
 In the past. two other high-peaked BL Lac AGN that are known TeV emitters (Mrk 421 and PINS 2155-304) have exhibited similar clockwise loops in their hysteresis plots (Takahashiet1996;Kataokaοἱal. 2000).," In the past, two other high-peaked BL Lac AGN that are known TeV emitters (Mrk 421 and PKS 2155-304) have exhibited similar clockwise loops in their hysteresis plots \citep{tak96, kat00}."
. As mentioned previously. Sembayetal.(2002) did not observe livsteresis wilh NMM observations of Mrk 421. however these observations did nol sample the rise and [all of the fux. variability so it is reasonable to expect a null result in that case.," As mentioned previously, \citet{sem02} did not observe hysteresis with XMM observations of Mrk 421, however these observations did not sample the rise and fall of the flux variability so it is reasonable to expect a null result in that case."
 The one flare that was completely observed by Sembayetal.(2002). did display the characteristic spectral softening with decreasing flux. but no loop could be discerned [rom this low amplitude variation.," The one flare that was completely observed by \citet{sem02} did display the characteristic spectral softening with decreasing flux, but no loop could be discerned from this low amplitude variation."
 Another observation of Mrk 421 using BeppoSAX showed evidence of spectral hysteresis with (he orientation in the counter-clockwise direction., Another observation of Mrk 421 using BeppoSAX \citep{zha02} showed evidence of spectral hysteresis with the orientation in the counter-clockwise direction.
 So. for at least one object. the hysteresis characteristics can change from flare to flare.," So, for at least one object, the hysteresis characteristics can change from flare to flare."
 Since the lLysteresis probes (he timescale of the cooling processes as a ΠΟΙΟ of energv. il appears as though some [lares have a cooling time that acts faster al higher energies. consistent with svnchrotron cooling. while other flares exhibit cooling times that are longer.," Since the hysteresis probes the timescale of the cooling processes as a function of energy, it appears as though some flares have a cooling time that acts faster at higher energies, consistent with synchrotron cooling, while other flares exhibit cooling times that are longer."
 For the 1114264423 observations presented here. (he generally clockwise hyvsteresis curves are consistent with the scenario in which the X-ray flux is dominated by svnchrotron emission since the svnchrotron cooling timescale is shorter al higher energies.," For the H1426+428 observations presented here, the generally clockwise hysteresis curves are consistent with the scenario in which the X-ray flux is dominated by synchrotron emission since the synchrotron cooling timescale is shorter at higher energies."
 In contrast. if the hvsteresis loops were oriented in the opposite direction. one would expect (he source to be in a regime in which the cooling ancl acceleration (times were approximately equal. according to models of irk.Rieger.andMastichiadis(1993).," In contrast, if the hysteresis loops were oriented in the opposite direction, one would expect the source to be in a regime in which the cooling and acceleration times were approximately equal, according to models of \citet{kir98}."
. By performing a study of the X-ray hvsteresis. one can also explore the potential οἱ various acceleration mechanisms to produce the TeV emission.," By performing a study of the X-ray hysteresis, one can also explore the potential of various acceleration mechanisms to produce the TeV emission."
 Typically. clockwise hysteresis patterns (using (he axis orientation shown in preceding ligures) are interpreted. as a solt lime lag in N-ray. enussion. while counter-clockwise orientation is interpreted. as a hard lag (IXirk.Rieger.and.Mastichiadis1998:Iataoka2000).," Typically, clockwise hysteresis patterns (using the axis orientation shown in preceding figures) are interpreted as a soft time lag in X-ray emission, while counter-clockwise orientation is interpreted as a hard lag \citep{kir98,kat00b}."
.. A soft lag at. N-ray energies. which is interpreted. as the energy dependence of the svnchrotron cooling (ime. is characteristic ol svnehrotron self Compton models (Takahashietal.1996).," A soft lag at X-ray energies, which is interpreted as the energy dependence of the synchrotron cooling time, is characteristic of synchrotron self Compton models \citep{tak96}."
.. A hard lag is not consistent wilh some standard one-zone svnchrotron self Compton models. while models such as (hat of Kirk.Rieger.andMastichiacdis(1998).. which incorporates a time-dependent. propagating shock that accelerates electrons in the region of the shock front ancl allows [or the case where TesotingvTucerte COL be consistent with a hard lag.," A hard lag is not consistent with some standard one-zone synchrotron self Compton models, while models such as that of \citet{kir98}, which incorporates a time-dependent propagating shock that accelerates electrons in the region of the shock front and allows for the case where $\tau_{cooling}\sim\tau_{accel}$, can be consistent with a hard lag."
 The observations presented here are generally in agreement with standard models that predict a soft lag. but the lack of hysteresis [rom the flare in region B does not favor simple one-component models during that particular Gime period.," The observations presented here are generally in agreement with standard models that predict a soft lag, but the lack of hysteresis from the flare in region B does not favor simple one-component models during that particular time period."
 However. the model of LiaudIxusunose.(2000) can reproduce hysteresis loops that. are not clockwise (ancl presumably intermediate situations with no obvious loop orientation) bv increasing the injection energy to the point where svnchrotron," However, the model of \citet{li00} can reproduce hysteresis loops that are not clockwise (and presumably intermediate situations with no obvious loop orientation) by increasing the injection energy to the point where synchrotron"
complex dvnamics in the cisk/magnetosphere interface.,complex dynamics in the disk/magnetosphere interface.
 Such phenomena should. be time-dependent. ancl may. be monitored in the future using Astro-I2., Such phenomena should be time-dependent and may be monitored in the future using Astro-E2.
 Our observations have revealed an Fe NNVI line at ~7 keV. detected. during low-state.," Our observations have revealed an Fe XXVI line at $\sim7$ keV, detected during low-state."
 The line is not detected. during the main-on when the X-ray spectrum. is dominated: by the strong continuum. emission., The line is not detected during the main-on when the X-ray spectrum is dominated by the strong continuum emission.
 TFhis feature has not been detected by Leahy et al. (, This feature has not been detected by Leahy et al. (
2001) using data.,2001) using data.
 Llowever. has a relatively low spectral resolution around. these energies (LS percent at 6 keV).," However, has a relatively low spectral resolution around these energies (18 percent at 6 keV)."
 observed the source with better spectral resolution (2 percent at 5.9 keV. using the SIS detectors) during an early main on state. and based on these data Indo et al. (," observed the source with better spectral resolution (2 percent at 5.9 keV using the SIS detectors) during an early main on state, and based on these data Endo et al. ("
2000) reported the first resolution of the broad feature. into two components. at 6.41. keV and 6.73 keV. However. this double structure has not been detected in other observations taken curing the main-on nor is it present in data at similar phases.,"2000) reported the first resolution of the broad feature into two components, at $6.41$ keV and $6.73$ keV. However, this double structure has not been detected in other observations taken during the main-on nor is it present in data at similar phases."
 The EPLC detectors on board. have a spectral resolution similar to at 6 keV. However. has a much greater cllective area allowing lines to be detected with a greater confidence. in particular during the states of low intensity of the binary system.," The EPIC detectors on board have a spectral resolution similar to at 6 keV. However, has a much greater effective area allowing lines to be detected with a greater confidence, in particular during the states of low intensity of the binary system."
 The feature at ~7 keV has been observed. by for the first time over several 35 day. phases., The feature at $\sim 7$ keV has been observed by for the first time over several 35 day phases.
 Also. it has been confirmed by a WEPGS observation of the source (the only one mace curing the low state) taken at Qus=O44O46 (Jimenez-Garate et al.," Also, it has been confirmed by a HETGS observation of the source (the only one made during the low state) taken at $\Phi_{35}= 0.44-0.46$ (Jimenez-Garate et al."
 2003)., 2003).
 The feature cannot be produced. by [üorescence. and. it is more Likely to be a Fe XXVI line originating in widely extended photo-ionized. plasma.," The feature cannot be produced by fluorescence, and it is more likely to be a Fe XXVI line originating in widely extended photo-ionized plasma."
 On the other hand. RGS data taken during the low and short-on states also show the presence of photo-ionized eas (Jimenez-Garate ct al.," On the other hand, RGS data taken during the low and short-on states also show the presence of photo-ionized gas (Jimenez-Garate et al."
 2002)., 2002).
 Grating spectra exhibit several narrow recombination emission lines. the most prominent being € VIN VI. N VIL. O VIL ο VILE and Ne IX.," Grating spectra exhibit several narrow recombination emission lines, the most prominent being C VI, N VI, N VII, O VII, O VIII and Ne IX."
 The line ratio @=(f| Dr. as computed for all the helium-like ion complexes. is Co4. which indicates that photoionization is the dominant Moreover. RGS spectra shows two radiative recombination continua of O ΥΠ and N VIL consistent with a low temperature of the emitting plasma (30000 Ix. «7 60000 Ix. Jimenez-Garate et al.," The line ratio $G=(f+i)/r$ , as computed for all the helium-like ion complexes, is $G \simeq 4$, which indicates that photoionization is the dominant Moreover, RGS spectra shows two radiative recombination continua of O VII and N VII, consistent with a low temperature of the emitting plasma (30000 K $<T<$ 60000 K, Jimenez-Garate et al."
 2002). which furtherly validate the photoionization moclel.," 2002), which furtherly validate the photoionization model."
 None of these features is detected during the main-on state (where instead there might be evidence for a broad O VIL iline. and perhaps O VILL Ly alpha).," None of these features is detected during the main-on state (where instead there might be evidence for a broad O VII i line, and perhaps O VIII Ly alpha)."
 The recombination X-ray line emission are not likely to originate in LIZ Lor. clue to the absence of UV induced photoexcitation signatures in the Lle-like triplets (observed with LUEPGS. Jimenez-Carate private conim.)," The recombination X-ray line emission are not likely to originate in HZ Her, due to the absence of UV induced photoexcitation signatures in the He-like triplets (observed with HETGS, Jimenez-Garate private comm.)."
 1nstead. as suggested by Jimenez-Garate et al. (," Instead, as suggested by Jimenez-Garate et al. ("
2002). an extended. photoionized accretion disk. atmosphere and corona may be responsable for such features.,"2002), an extended, photoionized accretion disk atmosphere and corona may be responsable for such features."
 Jimenez-Garate ct al... (, Jimenez-Garate et al. (
2001). computed the spectra of the atmospheric lavers of a Shakura-Sunvacy accretion disk. illuminatec by a central X-ray continuum.,"2001) computed the spectra of the atmospheric layers of a Shakura-Sunyaev accretion disk, illuminated by a central X-ray continuum."
 They found that. under these conditions. the disk develops both an extende corona which is kept hot at the Compton. temperature. and a more compact. colder. X-ray. recombination-emitting atmospheric laver. (see. Figure 9 in. Jimenez-Garate e al.," They found that, under these conditions, the disk develops both an extended corona which is kept hot at the Compton temperature, and a more compact, colder, X-ray recombination-emitting atmospheric layer (see Figure 9 in Jimenez-Garate et al."
 2002)., 2002).
 The parameters of the compact. atmospheric laver are consistent with the constraints to the narrow line emitting region. as derived from spectroscopic analysis ane modeling of the RGS features.," The parameters of the compact atmospheric layer are consistent with the constraints to the narrow line emitting region, as derived from spectroscopic analysis and modeling of the RGS features."
 Interestingly. the Fe XXVI line detected. by may be a signature of theholes! external lavers of the disk corona. which are located above the recombination-emitting lavers.," Interestingly, the Fe XXVI line detected by may be a signature of the external layers of the disk corona, which are located above the recombination-emitting layers."
 In order to detect a line at ~7 keV. an ionization parameter log£z3.3 is required. (Ixallman anc AleCrayv 1982. their Fig.," In order to detect a line at $\sim 7$ keV, an ionization parameter $\log \xi \geq
3.3$ is required (Kallman and McCray 1982, their Fig."
 Lh)., 1b).
" Using the Lux of power law component reported by R02 (ic. 3.65«10.7 ergf/em/sa Φιν— 0.17) and a source distance of 6.6 kpc. the luminosity of the direct beam is L,=1.9«107 erg/s. ""his gives where 7 is the particle density anc x djs the size of the region enclosing the Fe NAVI line emitting. plasma."," Using the flux of power law component reported by R02 (i.e. $3.65 
\times 10^{-8}$ $^2$ /s at $\Phi_{35}=0.17$ ) and a source distance of $6.6$ kpc, the luminosity of the direct beam is $L_x=1.9 \times
10^{38}$ erg/s. This gives where $n$ is the particle density and $r$ is the size of the region enclosing the Fe XXVI line emitting plasma."
 A second constraint can be derived from the emission measure., A second constraint can be derived from the emission measure.
" Running a detailed. multi-temperature. plasma model. in XSTAR gives an emission measure of n?V.26.1107 em [or the 7 keV line. for a llux of 2«10.+ photons/cnr /s (the emission measure scales linearly with the line flux. and the latter varies between 1.2«1071 photons/eni"". see Figure 6))'"," Running a detailed multi-temperature plasma model in XSTAR gives an emission measure of $n^2 
V> 6.1 \times 10^{57}$ $^{-3}$ for the 7 keV line, for a flux of $2 \times 10^{-4}$ $^{2}$ /s (the emission measure scales linearly with the line flux, and the latter varies between $1-2 \times10^{-4}$ $^{2}$, see Figure \ref{fe_line})."
 The average temperature. computed. by using a main-on Custom ionization spectrum in ASTPAR. as well as HULLAC recombination emissivities (D. Liedahl. private comm.).," The average temperature, computed by using a main-on custom ionization spectrum in XSTAR as well as HULLAC recombination emissivities (D. Liedahl, private comm.),"
 is A~580 eV. Moreover. if the line emitting plasma is located outside the Alfvénn shell. ro3.6.10* em.," is $kT \sim 580$ eV. Moreover, if the line emitting plasma is located outside the Alfvénn shell, $r>3.6 \times
10^7$ cm."
 These three constraints define an allowed: region in the plane nr. which is shown in Figure 10..," These three constraints define an allowed region in the plane $n-r$, which is shown in Figure \ref{iron_reg}."
 The density of the hot coronal lavers in the disk models by Jimenez-Garate et al. (, The density of the hot coronal layers in the disk models by Jimenez-Garate et al. (
2001) is Lot1015 ? if the radius is r~103107 em. therefore within the limits shown in the figure.,"2001) is $10^{15}-10^{16}$ $^{-3}$ if the radius is $r \sim
10^8-10^{10}$ cm, therefore within the limits shown in the figure."
 In summary. the variabilitv of the Her N-1 spectrum lends support to the precession of the accretion disk.," In summary, the variability of the Her X-1 spectrum lends support to the precession of the accretion disk."
 The evidence for the disk identification relies on the moceled structure and. spectra from a photoionized. disk. CJimenez- ct al., The evidence for the disk identification relies on the modeled structure and spectra from a photoionized disk (Jimenez-Garate et al.
 2001). which is in agreement with the limit set spectroscopically on the density of the low-energy lines emitting region.," 2001), which is in agreement with the limit set spectroscopically on the density of the low-energy lines emitting region."
 Interestingly. the same model naturally provides a candidate for the region emitting the Fe NNVI," Interestingly, the same model naturally provides a candidate for the region emitting the Fe XXVI"
ealaxv a rather unique object of the nearby Universe.,galaxy a rather unique object of the nearby Universe.
 Cousidering that GRBs represent extremely rare events. our analvsis supports therefore the existence of a causal luk between the origin of 9980125 aud the activity of star formation in the WR region.," Considering that GRBs represent extremely rare events, our analysis supports therefore the existence of a causal link between the origin of 980425 and the activity of star formation in the WR region."
 More generally: aud takiug also account of the properties characterizing the whole population of LGRD hosts. our results corroborate +i6 common idea that the trigecring of long GRBs is mostly associated with very active star formation iu chemically-voung euvironmeuts.," More generally, and taking also account of the properties characterizing the whole population of LGRB hosts, our results corroborate the common idea that the triggering of long GRBs is mostly associated with very active star formation in chemically-young environments."
" This work was enabled based on the funding frou the IRS aud the MIPS projects which are supported by NASA through the Jet Propulsion Laboratory (subcontracts #11257181 469960755). and thanks to the efficient technical support provided by the Science Center,"," This work was enabled based on the funding from the IRS and the MIPS projects which are supported by NASA through the Jet Propulsion Laboratory (subcontracts 1257184 960785), and thanks to the efficient technical support provided by the Science Center."
 We want to thauk our referee for relevant sugecstious and a careful review of the mnauuscript. as well as Michal Michalowski for useful conmunueuts on our work.," We want to thank our referee for relevant suggestions and a careful review of the manuscript, as well as Michal owski for useful comments on our work."
 We are particularly erateful to Vianney Lebouteiller for fruitfil discussious and for providing us with updated IRS data reduction prior to the publication of the Cornell AtlaS of Spitzer IRS Sources (CASSIS.Lebouteilleretal.2011a).. and to To Seong Ilhvaug for sharing his results ou the SDSS data.," We are particularly grateful to Vianney Lebouteiller for fruitful discussions and for providing us with updated IRS data reduction prior to the publication of the Cornell AtlaS of Spitzer IRS Sources \citep[CASSIS,][]{Lebouteiller11b}, and to Ho Seong Hwang for sharing his results on the SDSS data."
 We also want to acknowledge Yauling Wu for her help in the IRS data reduction. as well as Jason Marshall for providing some of his IR SED iodoeliug and JD Sinith for making publicly available his PAIIFIT routines.," We also want to acknowledge Yanling Wu for her help in the IRS data reduction, as well as Jason Marshall for providing some of his IR SED modeling and JD Smith for making publicly available his PAHFIT routines."
 We ereatlv appreciated the help from Stephane Arnouts and Olivier Πρωτ when using the code LePhare., We greatly appreciated the help from Stephane Arnouts and Olivier Ilbert when using the code $LePhare$.
 We finally thauk Tio Díaz-Sautos. Eréddérric Galliano. David Elbaz. Suzanne Madden and Mare Sauvage for useful discussions.," We finally thank Tanio az-Santos, Fréddérric Galliano, David Elbaz, Suzanne Madden and Marc Sauvage for useful discussions."
redshifts are well distributed round the sky. but with some preference for lower LIatitucdes.,"redshifts are well distributed round the sky, but with some preference for lower latitudes."
 Our extinction maps started [rom the {του intensity mips of Rowan-Robinson (1991). binned into lune bins. with the τουfed ratio as given by Boulanger ancl Perault (1988). ancl ele/d ratio as given by Mathis (1990).," Our extinction maps started from the $I_{100}$ intensity maps of Rowan-Robinson (1991), binned into lune bins, with the $I_{100}/A_V$ ratio as given by Boulanger and Perault (1988), and $A_B/A_V$ ratio as given by Mathis (1990)."
 Because he dust temperature declines with galactic radius. this oocedure over-estimates extinctions towards the galactic centre. and. under-estimates them towards the anti-centre.," Because the dust temperature declines with galactic radius, this procedure over-estimates extinctions towards the galactic centre, and under-estimates them towards the anti-centre."
 We attempted to correct. for. this. dependence: of. dust emperature on position by devising a simple model for he cust and starlight in our galaxy., We attempted to correct for this dependence of dust temperature on position by devising a simple model for the dust and starlight in our galaxy.
" We assumed a coubly exponential. optically thin distribution of dust. and. stars. with standard LAW scalelengths for the solar radius and stellar distribution Gry=Ss5kpe.r,3X5kpe.z.0.35 kpc) and an assumed the dust to have the same radial and half the vertical scalelength."," We assumed a doubly exponential, optically thin distribution of dust and stars, with standard IAU scalelengths for the solar radius and stellar distribution $r_0=8.5 \kpc, r_{sc}=3.5 \kpc, z_{sc}=0.35 \kpc$ ) and an assumed the dust to have the same radial and half the vertical scalelength."
 The dust. properties were assumed to be uniform. and we assumed a dust temperature of 20A. at the solar radius. and proportional to the one fifth »ower of the local stellar density elsewhere.," The dust properties were assumed to be uniform, and we assumed a dust temperature of $20K$ at the solar radius, and proportional to the one fifth power of the local stellar density elsewhere."
 We then found. or each position on the celestial sphere. the ratio of 1007/2 emission to column density of dust. and normalised this ratio ον the Boulanger and Perault value for the δα.," We then found, for each position on the celestial sphere, the ratio of $100\mum$ emission to column density of dust, and normalised this ratio by the Boulanger and Perault value for the NGP."
 Subsequent to our definition of the PSCz catalogue and mask. Schlegel (1998) used the LIGAS ISSA and CODE DIRBE data to make beautiful high resolution maps of the dust emission and extinction in our Galaxy.," Subsequent to our definition of the PSCz catalogue and mask, Schlegel (1998) used the IRAS ISSA and COBE DIRBE data to make beautiful high resolution maps of the dust emission and extinction in our Galaxy."
 Comparison of our own maps and those of Schlegel show that (a) our emperature corrections across the galaxy are too small. and (b) there are several areas where cooler (1615) dust extends o [b]~30°," Comparison of our own maps and those of Schlegel show that (a) our temperature corrections across the galaxy are too small, and (b) there are several areas where cooler (16K) dust extends to $|b|\sim 30 \dg$."
 Overall. our extinction estimates may be in error by a [actor 1.5-2. in the sense of being too low towards the anticentre and too high towards the galactic centre.," Overall, our extinction estimates may be in error by a factor 1.5-2, in the sense of being too low towards the anticentre and too high towards the galactic centre."
 This has two principal cllects. (, This has two principal effects. (
"a) Towards the anticentre. galaxies may fall below the 6,=19.5"" limit because of extinction.","a) Towards the anticentre, galaxies may fall below the $b_J=19.5^m$ limit because of extinction."
" A subsequent program of A""-imaging has revealed a handful of extra nearby galaxies. and a significant number at redshilts 0.05 and greater. ("," A subsequent program of $K'$ -imaging has revealed a handful of extra nearby galaxies, and a significant number at redshifts 0.05 and greater. ("
b) Phe definition of the optically-selected catalogue in Section 2.3 depends on the extinction corrections.,b) The definition of the optically-selected catalogue in Section 2.3 depends on the extinction corrections.
 We will have selected too many optical galaxies towards the centre and too few towards the anticentre., We will have selected too many optical galaxies towards the centre and too few towards the anticentre.
 llowever. our matching of addscan ancl PSC fluxes. was designed to be robust to exactly this sort of error.," However, our matching of addscan and PSC fluxes was designed to be robust to exactly this sort of error."
 To test for any οσοι we have compared the simple. number-weighted dipole of the surface distribution of PSC galaxies. (a) using the normal catalogue and. (b) using purely PSC-derived fluxes as described in section 2.3.," To test for any effect, we have compared the simple, number-weighted dipole of the surface distribution of PSCz galaxies, (a) using the normal catalogue and (b) using purely PSC-derived fluxes as described in section 2.3."
 The dipole changes by 1 in direction and in amplitude when we do this. showing that anv bias caused by incorrect extinctions is negligible.," The dipole changes by $1\dg$ in direction and in amplitude when we do this, showing that any bias caused by incorrect extinctions is negligible."
 The error. quoted for PSC 607m Iluxes for genuine point sources is just. LEAGUES. VIL.D.2)., The error quoted for PSC $60\mum$ fluxes for genuine point sources is just (ES VII.D.2).
 We are more concerned with any non-fractional random. error component. since this may lead. to Aalmauist-type biases.," We are more concerned with any non-fractional random error component, since this may lead to Malmquist-type biases."
 The analysis in Lawrence (1999). based on the 12/6052. colours. of bright stars. finds an absolute error. of 0.059+0.007Jv. in addition to a fractional error of4.," The analysis in Lawrence (1999), based on the $12/60\mum$ colours of bright stars, finds an absolute error of $0.059
\pm 0.007 \Jy$, in addition to a fractional error of."
. Stars are better point sources than galaxies. so this may underestimate the absolute error for our sample.," Stars are better point sources than galaxies, so this may underestimate the absolute error for our sample."
 We have made an independent estimate. by investigating the scatter between our PSC and addscan [luxes.," We have made an independent estimate, by investigating the scatter between our PSC and addscan fluxes."
 OF course PSC and. addscan Uuxes start from the same raw data. but the processing and background estimation are entirely dillerent. while the actual photon noise is negligible.," Of course PSC and addscan fluxes start from the same raw data, but the processing and background estimation are entirely different, while the actual photon noise is negligible."
 Phe absolute component of the scatter is 0.06—0.07Jv. of which an estimated comes from the error in the adcdscan flux - confirming that the Lawrence value is a reasonable estimate of the PSC absolute error component.," The absolute component of the scatter is $0.06-0.07\Jy$, of which an estimated comes from the error in the addscan flux - confirming that the Lawrence value is a reasonable estimate of the PSC absolute error component."
 This error estimate leads to Malmequist biases in the source densities of order (dillerential. at the Hus limit) and (cumulative) (Murdoch. Crawford. ancl Jauncey 1973).," This error estimate leads to Malmquist biases in the source densities of order (differential, at the flux limit) and (cumulative) (Murdoch, Crawford and Jauncey 1973)."
 3116ΟΝ sky may have biases as large as (dilferential) and (cumulative). so the non-uniformity introduced into the catalogue should be no worse than (düllerential) and (cumulative).," 2HCON sky may have biases as large as (differential) and (cumulative), so the non-uniformity introduced into the catalogue should be no worse than (differential) and (cumulative)."
 This is borne out by the source counts in Figures 7 and S. but note that in any noise-imited catalogue such as the PSC. where the noise varies across the catalogue. there will always be a regime where Malmquist. biases are masked by incompleteness.," This is borne out by the source counts in Figures 7 and 8, but note that in any noise-limited catalogue such as the PSC, where the noise varies across the catalogue, there will always be a regime where Malmquist biases are masked by incompleteness."
 Lawrence (1999) find evidence for slight non-linearity in the PSC flux scale. but this should not allect any analysis except evolutionary. studies.," Lawrence (1999) find evidence for slight non-linearity in the PSC flux scale, but this should not affect any analysis except evolutionary studies."
 The οσο on evolution. was considered by Saunders (19:]., The effect on evolution was considered by Saunders (1997).
 on-uniformity can also come about as a result. of simple changes in the absolute Lux scale., Non-uniformity can also come about as a result of simple changes in the absolute flux scale.
 A small fractional error in he calibration will on average lead to an error half again as arge in the source density., A small fractional error in the calibration will on average lead to an error half again as large in the source density.
 “Phere are three obvious reasons or Εις scale variations: 1) The absolute calibration of the third 31ICON was revised wea few percent after the release of the PSC (ancl FSS)., There are three obvious reasons for flux scale variations: 1) The absolute calibration of the third 3HCON was revised by a few percent after the release of the PSC (and FSS).
 The etfect of this revision on PSC lluxes would be to change hose in the of the sky. covered by BLICONS byL&., The effect of this revision on PSC fluxes would be to change those in the of the sky covered by 3HCONS by.
. 2) Whenever the satellite entered the South Atlantic Anomaly. radiation hits altered. the sensitivity. of the detectors. and data taken curing such times were discarded (ES LLC4).," 2) Whenever the satellite entered the South Atlantic Anomaly, radiation hits altered the sensitivity of the detectors, and data taken during such times were discarded (ES III.C.4)."
 However. this leaves the possibility of AMalmauist effects. and also data taken near the boundaries ofthe SAA potentially sullers residual ellects.," However, this leaves the possibility of Malmquist effects, and also data taken near the boundaries ofthe SAA potentially suffers residual effects."
. We have investigated. this by checking the source counts for declinations 40°<ὃ«10°. where data taken is most ikelv to be alfected.," We have investigated this by checking the source counts for declinations $-40\dg < \delta < 10\dg$, where data taken is most likely to be affected."
 We see no evidence for any variation (ligure 7) above the few level in source counts., We see no evidence for any variation (Figure 7) above the few level in source counts.
 L) Whenever the satellite crossed. the Galactic Plane. or other verv bright sources. the detectors. sullerecd fron ivesteresis.," 3) Whenever the satellite crossed the Galactic Plane, or other very bright sources, the detectors suffered from hysteresis."
 This effect was investigated. by Strauss 990)., This effect was investigated by Strauss (1990).
 νου found that the likely error. is typically less han and always less than 2.2%.., They found that the likely error is typically less than and always less than .
 This is confirmed by, This is confirmed by
in any individual viewing period (VPs).,in any individual viewing period (VPs).
" The locations of low-conficlence EGRET sources are nol well-determined owing to the wide point spread Iunction (PSF) at photon energies >100 MeV (Thompsonetal.1993:Mattox,Hartman.&Reimer2001).. aud identification of EGRET sources with counterparts based on position alone is difficult."," The locations of low-confidence EGRET sources are not well-determined owing to the wide point spread function (PSF) at photon energies $> 100$ MeV \citep{thompson93,mhr01}, and identification of EGRET sources with counterparts based on position alone is difficult."
 Analvsis of EGRET data from 1991 1995 vielded a 7.7 6 detection ancl a bared spectrum with photon index a=τὸ0.15 (Ilarimanetal.1999).. and 3EG J1222+2841 was identified with ON 231 with “high confidence’.," Analysis of EGRET data from 1991 – 1995 yielded a 7.7 $\sigma$ detection and a hard spectrum with photon index $\alpha = 1.73 \pm 0.18$ \citep{hartman99}, and 3EG J1222+2841 was identified with ON 231 with “high confidence”."
 However. this identification was based on the >1 GeV position of Lamb&Macomb(1997) rather than the £>100 MeV position. as is standard practice for EGRET sources.," However, this identification was based on the $E > 1$ GeV position of \cite{lm97} rather than the $E > 100$ MeV position, as is standard practice for EGRET sources."
 Mattox.Hartman.&Reimer(2001). have recently done a quantitative re-evaluation of potential radio identifications for the 3EG radio sources bv calculating the probability of each identification., \cite{mhr01} have recently done a quantitative re-evaluation of potential radio identifications for the 3EG radio sources by calculating the probability of each identification.
 They note that based on its £z100 MeV. position. 3EG J12224-2841. cannot be included in their list of high confidence identifications.," They note that based on its $E > 100$ MeV position, 3EG J1222+2841 cannot be included in their list of high confidence identifications."
 Using the E>1 GeV position. Mattox.Hartman.&Reimer(2001) get the probability of identification to be 4%... which is classilied as “plausible”.," Using the $E > 1$ GeV position, \cite{mhr01} get the probability of identification to be 4, which is classified as “plausible”."
 Future observations with GLAST will be important in determining (he source position wilh more accuracy. and securing a more confident identification.," Future observations with GLAST will be important in determining the source position with more accuracy, and securing a more confident identification."
 For the purpose of our analvsis. we assume thal Wo Comae is the counterpart of 83EG J12224-2841.," For the purpose of our analysis, we assume that W Comae is the counterpart of 3EG J1222+2841."
 llowever. the low significance of the ~2.76 detection in May. 1998 and its non-simultaneitv io the BeppoSAX observations prevents us from deriving anv strong constraints from the EGRET this or spectrum.," However, the low significance of the $\sim 2.7 \, \sigma$ detection in May 1998 and its non-simultaneity to the BeppoSAX observations prevents us from deriving any strong constraints from the EGRET flux or spectrum."
 EGRET has observed 3EG J12224-2841 = W Comae several times since ils launch in 1991., EGRET has observed 3EG J1222+2841 = W Comae several times since its launch in 1991.
 Table 1. lists the VPs during which the source was observed. and the corresponding integral[Iuxes for energies greater Chan LOO MeV. Only photons with inclination angles less {han 30° were used for (he analvsis.," Table \ref{history} lists the VPs during which the source was observed, and the corresponding integralfluxes for energies greater than 100 MeV. Only photons with inclination angles less than $30^\circ$ were used for the analysis."
 For Phase 1 through Cvele 4 of the EGRET observations (1991 1995). Table 1 lists data from the 3EG catalog (Hartmanetal.1999).," For Phase 1 through Cycle 4 of the EGRET observations (1991 – 1995), Table \ref{history} lists data from the 3EG catalog \citep{hartman99}."
. Three additional observations were made in Cveles 5. 7 and 9.," Three additional observations were made in Cycles 5, 7 and 9."
 We have analvzed these data using the standard EGRET data processing technique. as described in Mattoxetal.(1996). and llarimanetal.(1999).. and included them in the table.," We have analyzed these data using the standard EGRET data processing technique, as described in \cite{mattox96} and \cite{hartman99}, and included them in the table."
 The light curve for 3EG J12224-2341 is shown in Fig. 1.., The light curve for 3EG J1222+2841 is shown in Fig. \ref{fluxhist}.
 The figure shows fluxes for all detections at a level greater than 26: lor detections below 26. upper limits at the 95 confidence level are shown.," The figure shows fluxes for all detections at a level greater than $2\sigma$; for detections below $2\sigma$, upper limits at the 95 confidence level are shown."
 We have computed the background-subtracted οταν spectra of 3EG J12224-2841 [or the strongest detections. as well as for March. 1998. close to the time of theBeppoSAX observations.," We have computed the background-subtracted $\gamma$ -ray spectra of 3EG J1222+2841 for the strongest detections, as well as for March 1998, close to the time of the observations."
 The spectra were determined by dividing the EGRET energy band of 30 MeV 10 GeV into 4 bins. and estimating the number of source photons in each interval. following the standard EGRET spectral analvsis technique (Nolanetal. 1993)..," The spectra were determined by dividing the EGRET energy band of 30 MeV – 10 GeV into 4 bins, and estimating the number of source photons in each interval, following the standard EGRET spectral analysis technique \citep{nolan93}. ."
" We have fitted a single power law of the [orm F(E)—kKGE/Ej)"" photons 7? ! MeV! to the data. where"," We have fitted a single power law of the form $F(E) = k \, (E/E_0)^{-\alpha}$ photons $^{-2}$ $^{-1}$ $^{-1}$ to the data, where"
As the orbital phase progresses [rom phase 0.0. the blueward rise rapidly disappears. and the flix peak in the spectrum moves to the blue.,"As the orbital phase progresses from phase 0.0, the blueward rise rapidly disappears, and the flux peak in the spectrum moves to the blue."
 It is unclear whether this blue feature completely disappears. or is simply hidden by the much higher flux levels associated with the phase 0.5 evelotron emission.," It is unclear whether this blue feature completely disappears, or is simply hidden by the much higher flux levels associated with the phase 0.5 cyclotron emission."
 During the transition to maximum lieht. a weak rise al the extreme red end of the Z-band appears.," During the transition to maximum light, a weak rise at the extreme red end of the -band appears."
 While the S/N of the data in this region is «quite low. the combination of the// and A-band data sets suggest (hat (his rise is real. and is due to a broad feature that peaks somewhere between the edges of theLf andA bandpasses. in (he middle of the region stronglv affected by telluric water vapor absorption.," While the S/N of the data in this region is quite low, the combination of the and -band data sets suggest that this rise is real, and is due to a broad feature that peaks somewhere between the edges of the and bandpasses, in the middle of the region strongly affected by telluric water vapor absorption."
 As the orbital phase continues (to increase on ils wav back to phase 0.0. the spectra retrace (heir preanasxinnin evolution.," As the orbital phase continues to increase on its way back to phase 0.0, the spectra retrace their pre-maximum evolution."
 The A-band spectra obtained wilh NURI are shown in Fig., The -band spectra obtained with NIRI are shown in Fig.
 3. and the A-band spectra obtained with NIRSPEC are presented in Fig.," 3, and the -band spectra obtained with NIRSPEC are presented in Fig."
 4., 4.
 Both data sets show similar evolution over an orbital evcle. suggesting that little has changed in the intervening nine months.," Both data sets show similar evolution over an orbital cycle, suggesting that little has changed in the intervening nine months."
 Near phase 0.0. the spectrum shows a continuum that is consistent wilh Chat of a late-tvpe star. with declining fixes at both blue ancl red ends. presumably. caused by water vapor absorption in the photosphere of the secondary star.," Near phase 0.0, the spectrum shows a continuum that is consistent with that of a late-type star, with declining fluxes at both blue and red ends, presumably caused by water vapor absorption in the photosphere of the secondary star."
 The co-addition of the NIRSPEC phase 0.0 spectra (Fig., The co-addition of the NIRSPEC phase 0.0 spectra (Fig.
 5) do not confirm (his. however. with no clear sign of the expected spectral features for a late-tvpe star.," 5) do not confirm this, however, with no clear sign of the expected spectral features for a late-type star."
" While a feature consistent with CO», (at 2,204 jmi) appears to be present. other features with a similar depth do not correspond to anv of the expected atomic features (e.g... Na [ αἱ 2.20 jun) for a cool star."," While a feature consistent with $^{\rm 12}$ $_{\rm (2,0)}$ (at 2.294 $\mu$ m) appears to be present, other features with a similar depth do not correspond to any of the expected atomic features (e.g., Na I at 2.20 $\mu$ m) for a cool star."
 By coadding the entire NIRSPEC data set. the first overtone feature of Ορ appears to remain mareinally detectable.," By coadding the entire NIRSPEC data set, the first overtone feature of $^{\rm 12}$ $_{\rm (2,0)}$ appears to remain marginally detectable."
 This would be the first detection of the secondary star in EF Eri., This would be the first detection of the secondary star in EF Eri.
PAIASS Web )).,2MASS Web ).
" FIRST 1413|4505 is biehly luminous in the radio. PiaoH,=2.2.107 W + and has a [lat radio spectral index. a= 0.09."," FIRST 1413+4505 is highly luminous in the radio, $P_{\rm 1.4 \ GHz} =
2.2 \times 10^{27}$ W $^{-1}$, and has a flat radio spectral index, $\alpha=-0.09$ ."
 From the ££ magnitude and using the ENG spectrum for the A-correction we obtained Alyy (1450 A)) — 25.7 (see equations in. Vigotti et al., From the $E$ magnitude and using the TNG spectrum for the $k$ -correction we obtained $M_{\rm AB}$ (1450 ) = $-25.7$ (see equations in Vigotti et al.
 2003)., 2003).
" At the position of the QSO. mid-infrared.: emission. is detected via 1D addition of LAS Scans (SCANDPL) at 12/0 and 255m. The automatic procedure obtained via NED gives an lle detection at 12/7 in the αστοι moces “detector weighted? and ""mean. 15e detection for ""rms weighted and non-detection for the ""median mode."," At the position of the QSO, mid-infrared emission is detected via 1D addition of IRAS Scans (SCANPI) at $\mu$ m and $\mu$ m. The automatic procedure obtained via NED gives an $\sigma$ detection at $\mu$ m in the addition modes `detector weighted' and `mean', $\sigma$ detection for `rms weighted' and non-detection for the `median' mode."
 The corresponding peak Dux densities in Jv are 0.22£0.02. 0.21+0.02and 0.31+0.02. respectively.," The corresponding peak flux densities in Jy are $0.22
\pm 0.02$, $0.21 \pm 0.02$and $0.31 \pm 0.02$, respectively."
" At 255m. coaddition gives 3.60 detection for ""detector. weighted? and ""mean modes. 5a for ""rms weighted! ancl non-detection. with the ""median? moce."," At $\mu$ m coaddition gives $\sigma$ detection for `detector weighted' and `mean' modes, $\sigma$ for `rms weighted' and non-detection with the `median' mode."
 The corresponding peak flux densities are 0.08d:0.02. 0.08+0.02 and 0.11+0.02.," The corresponding peak flux densities are $0.08 \pm 0.02$, $0.08 \pm
0.02$ and $0.11 \pm 0.02$."
 Mict-infrared cdetection at 2z is rare and we checked if the detection with SCANDPL could arise due to confusion with one of the two bright stars 52 aresce NE and 48 arcsec SW (binary) of the QSO. but neither of them was detected al 12jm or 254un. Assumingὃν that the mid-infrared. emission is physically associated with the QSO. and not a chance coincidence. the uminosity of the QSO in the rest [rame wavelength range 2200 6.1 pm (from z-band to 25 pim in the observer rame) would be Ls-10/7£.. Le. FIRST 1413|4505 would be one of the most luminous objects known.," Mid-infrared detection at $z \ge 2$ is rare and we checked if the detection with SCANPI could arise due to confusion with one of the two bright stars 52 arcsec NE and 48 arcsec SW (binary) of the QSO, but neither of them was detected at $\mu$ m or $\mu$ m. Assuming that the mid-infrared emission is physically associated with the QSO, and not a chance coincidence, the luminosity of the QSO in the rest frame wavelength range 2200 – 6.1 $\mu$ m (from $z$ -band to 25 $\mu$ m in the observer frame) would be $1.8
\times 10^{15} L_\odot$, i.e. FIRST 1413+4505 would be one of the most luminous objects known."
 To our knowledge there are five other svstems with such extreme uminosities. and although they were discovered in different wavs. and have dillerent SEDs and absorption/emission-line operties. the extreme Luminosity was in all cases ascribed o tux magnification due to gravitational lensine.," To our knowledge there are five other systems with such extreme luminosities, and although they were discovered in different ways, and have different SEDs and absorption/emission-line properties, the extreme luminosity was in all cases ascribed to flux magnification due to gravitational lensing."
 “These objects are the Sevfert 2 like object LIGAS FSC 10214|4724 (Rowan-Robinson et al., These objects are the Seyfert 2 like object IRAS FSC 10214+4724 (Rowan-Robinson et al.
 1991). the BAL QSO LI1413|117 (Magazin et al.," 1991), the BAL QSO H1413+117 (Magain et al."
 LOSS). the sub-mim galaxv SMM. 02399-)136 (vison et al.," 1988), the sub-mm galaxy SMM 02399-0136 (Ivison et al."
 1998). the optically bright QSO APM )8219|5255 (Irwin et al.," 1998), the optically bright QSO APM 08279+5255 (Irwin et al."
 1998) and the Lyman break galaxy AIS 1512-cD58 (Yee et al., 1998) and the Lyman break galaxy MS 1512-cB58 (Yee et al.
 1996)., 1996).
 Fig., Fig.
e 5 suggestsoo that no QSOs with zzc3.7 will have been missed by the colour selection O—E> 2., 5 suggests that no QSOs with $z \ge 3.7$ will have been missed by the colour selection $O-E >$ 2.
 The properties of the 10 z 3.7 OSOs are given in Table 6., The properties of the 10 $z >$ 3.7 QSOs are given in Table 6.
 Adopting the magnitude limit Lo<19.1. leaves 7 QSOs in the range B37moz€ 1. ," Adopting the magnitude limit $E \le 19.1$, leaves 7 QSOs in the range $3.7 \le z \le 4.4$ ."
The absolute magnitudes at rest-[rame1H250-A.. ΑΔνο (108 AJ). of the zc3.7 QSOs were calculated from the extinetion-corrected ££ magnitudes using the same procedure as in Vigotti ct al. (," The absolute magnitudes at rest-frame, $M_{\rm AB}$ (1450 ), of the $z \ge 3.7$ QSOs were calculated from the extinction-corrected $E$ magnitudes using the same procedure as in Vigotti et al. ("
2003). except [or a better estimation of the A-corrections. obtained here. from. the individual spectra rather than using an average.,"2003), except for a better estimation of the $k$ -corrections, obtained here from the individual spectra rather than using an average."
 We used the optical spectra from this work (FIRST 1110|4305). Benn ot al.," We used the optical spectra from this work (FIRST 1110+4305), Benn et al."
 2002 (FIRST 094115119) or from SDSS (remaining QSOs)., 2002 (FIRST 0941+5119) or from SDSS (remaining QSOs).
 Fig., Fig.
 Sa shows manL450(1ον or &-correction. [or the redshift range 3.7 0 4.4. using the spectra of the 10 QSOs.," 8a shows $m_{\rm AB} [1450 (1+z) ]- E$, or $k$ -correction, for the redshift range 3.7 – 4.4, using the spectra of the 10 QSOs."
 Fig., Fig.
 Sb shows the mean and. standard. deviation of may14500112}E for the seven QSOs with Ex19.1.," 8b shows the mean and standard deviation of $m_{\rm AB}
[1450 (1+z)] - E$ for the seven QSOs with $E \le 19.1$."
" The radio luminosities at rest-frame frequceney of 1.4 CGllz were caleulatecl assuming a spectral index a=0.8 (S, x7). which is the median spectral index obtained for 13 ο 3.6 FIRST-APAL QSOs in Holt et al. ("," The radio luminosities at rest-frame frequency of 1.4 GHz were calculated assuming a spectral index $\alpha = -0.3$ $S_\nu
\propto \nu^\alpha$ ), which is the median spectral index obtained for 13 $z >$ 3.6 FIRST-APM QSOs in Holt et al. ("
2004).,2004).
" The racio luminosities range from P4cg;=1077 to 1071 Ww to so all ten QSOs are ""radio loud’ QSOs. adopting the criterion £4GH,x107-- Wt (Ciregg ot al."," The radio luminosities range from $P_{\rm 1.4 ~GHz}=10^{25.7}$ to $10^{28.1}$ W $^{-1}$, so all ten QSOs are `radio loud' QSOs, adopting the criterion $P_{\rm 1.4 ~GHz} > 10^{25.5}$ W $^{-1}$ (Gregg et al."
 1996)., 1996).
 There are several sources of incompleteness. summarised below. (," There are several sources of incompleteness, summarised below. ("
1) Only 78 of the 94 candidates were spectroscopically classified. giving a fraction of 83 per cent. (,"1) Only 78 of the 94 candidates were spectroscopically classified, giving a fraction of 83 per cent. ("
2) The APAL completeness for ££x19.1 was estimated as the fraction of SDSS r<19.8. 2c 2 QSOs in the SDSSDR3 Quasar Catalog (Schneider et al.,"2) The APM completeness for $E \le 19.1$ was estimated as the fraction of SDSS $r \le 19.8$, $z \ge$ 2 QSOs in the SDSSDR3 Quasar Catalog (Schneider et al."
 2005) detected and starlike in APAL de., 2005) detected and starlike in APM $E$ .
 “Phe adopted. à band. limit was obtained from the average magnitude dillerence r£=0.69 (standard deviation 0.84) of the 10high-: QSOs in the sample., The adopted $r$ band limit was obtained from the average magnitude difference $r-E=0.69$ (standard deviation 0.34) of the 10$z$ QSOs in the sample.
 Average rf and standard. deviation for the 94 candidates are 0.55 and. 0.36. respectively., Average $r-E$ and standard deviation for the 94 candidates are 0.55 and 0.36 respectively.
 From the 2060, From the 2060
active life of the nucleus.,active life of the nucleus.
" We stress that the ""isotropy of broad line emission we refer to is relative to objects which. according to the current unification schemes. are observed within the opening angle of the putative obscuring torus (see e.g. Urry Padovani 1995 for a recent review)."," We stress that the `isotropy' of broad line emission we refer to is relative to objects which, according to the current unification schemes, are observed within the opening angle of the putative obscuring torus (see e.g. Urry Padovani 1995 for a recent review)."
 The outline of the paper is as follows., The outline of the paper is as follows.
 In Section 2. we describe the sample of sources. while in Sections 3 and 4 the methods applied ο estimate the luminosity in. broad lines and the jet kinetic* power are outlined.," In Section 2, we describe the sample of sources, while in Sections 3 and 4 the methods applied to estimate the luminosity in broad lines and the jet kinetic power are outlined."
 In Section 5 we examine the case for corisidering also BL Lacs objects in our study., In Section 5 we examine the case for considering also BL Lacs objects in our study.
 The results are oresented in Section 6 and. discussed in Section 7., The results are presented in Section 6 and discussed in Section 7.
 Section S summarizes our conclusions., Section 8 summarizes our conclusions.
 Cosmological parameters Lo=50 km/s/Mpce and qo=0 have been adopted throughout the paper., Cosmological parameters $_0$ =50 km/s/Mpc and $_0$ =0 have been adopted throughout the paper.
 We consider the sample of radioloud AGN assembled by Ghisellini et al. (, We consider the sample of radio–loud AGN assembled by Ghisellini et al. (
1993) 693]. which includes 105 sources with VLBI angular diameter data available in the literature. anc for which it is therefore. possible to apply the svnehrotron self.Compton (SSC) model and estimate the kinetic luminosity in the mas (milliaresec) scale jet.,"1993) [G93], which includes 105 sources with VLBI angular diameter data available in the literature, and for which it is therefore possible to apply the synchrotron self–Compton (SSC) model and estimate the kinetic luminosity in the mas (milliarcsec) scale jet."
 This is the only criterion. applied. to select. the sources. and therefore the sample is not. in any sense. completo.," This is the only criterion applied to select the sources, and therefore the sample is not, in any sense, complete."
 As already mentioned. we exclude from it radio galaxies. which we expect to be observed at large angles with respect to je jet axis and. whose broad line emission is therefore likely hidden by obscuring material.," As already mentioned, we exclude from it radio galaxies, which we expect to be observed at large angles with respect to the jet axis and whose broad line emission is therefore likely hidden by obscuring material."
 Among these radioloud. sources. we then consider the three subsamples of broad. line objects. namely: core dominated High anc Low Polarization Quasars (LIP) and LP respectively) and. lobe dominated: quasars (LDQ). obtaining a sample of G4 sources.," Among these radio–loud sources, we then consider the three subsamples of broad line objects, namely core dominated High and Low Polarization Quasars (HPQ and LPQ respectively) and lobe dominated quasars (LDQ), obtaining a sample of 64 sources."
 For a more detailed discussion about their classification as well as further data and references we refer to G93., For a more detailed discussion about their classification as well as further data and references we refer to G93.
 Fluxes at 1: keV for four sources without XN-rav data in G93 are derived. [rom theROSAT NGA catalogue (White. Giommi Angelini 1994) as described in Padovani. Ciommi Fiore (1996).," Fluxes at 1 keV for four sources without X-ray data in G93 are derived from the WGA catalogue (White, Giommi Angelini 1994) as described in Padovani, Giommi Fiore (1996)."
 This translates into a better determination of the Doppler factor (see below). which is otherwise based. for sources without A-ray data. on the optical Lux.," This translates into a better determination of the Doppler factor (see below), which is otherwise based, for sources without X-ray data, on the optical flux."
 We then have to estimate the total Iuminosity emitted. in broacl lines. Leer.," We then have to estimate the total luminosity emitted in broad lines, $\ll$."
 There is not any solidlv established procedure to derive Lyi and (unlike for example Faleke et al., There is not any solidly established procedure to derive $\ll$ and (unlike for example Falcke et al.
 1995) we adopt the following method. which involves three steps: 1) define a set of lines which dominates the total broad line emission: li) establish their relative Lux ratios: ill) extrapolate. using these ratios. from the luminosity in lines with measured [lux to the luminosity in all lines.," 1995) we adopt the following method, which involves three steps: i) define a set of lines which dominates the total broad line emission; ii) establish their relative flux ratios; iii) extrapolate, using these ratios, from the luminosity in lines with measured flux to the luminosity in all lines."
 i) Clearly. the more measured [ine lluxes are available. the more the estimate is correct.," i) Clearly, the more measured line fluxes are available, the more the estimate is correct."
 Furthermore. the use of fluxes from several lines increases the statistics in terni of number of sources in the sample and allows a better coverage in redshift.," Furthermore, the use of fluxes from several lines increases the statistics in term of number of sources in the sample and allows a better coverage in redshift."
 We consider fluxes for the following lines: Lya. € IV. Ale Lh Ue. ή. and. Ho. which are amongst the major contributors to the total Lyi» (making up in fact about GO per cent of it) and. are well identified in quasar spectra.," We consider fluxes for the following lines: $\alpha$ , C IV, Mg II, ${\gamma}$, ${\beta}$, and ${\alpha}$, which are amongst the major contributors to the total $\ll$ (making up in fact about 60 per cent of it) and are well identified in quasar spectra."
 We then collect. from the literature data on the broad. line Duxes., We then collect from the literature data on the broad line fluxes.
" In order to obtain clata as homogeneous as possible. ancl minimize the uncertainties. we only consider values of line fluxes (or. luminosities) either. when directly. given. or when the equivalent width and continuum [lux at the. corresponding line. frequency are reported by the same authors (for small differences between the line and continuum frequencies. we extrapolate thecontinuum. adopting a slope avy=0.5 between 2200 aand 2800.P(u) xu"")."," In order to obtain data as homogeneous as possible, and minimize the uncertainties, we only consider values of line fluxes (or luminosities) either when directly given or when the equivalent width and continuum flux at the corresponding line frequency are reported by the same authors (for small differences between the line and continuum frequencies, we extrapolate thecontinuum, adopting a slope $\alpha_{\rm UV}=0.5$ between 2200 and 2800, $F(\nu)\propto \nu^{-\alpha}$ )."
 We lind line Duxes for a subsample of 43 of the 64 objects. with an average of about two lines per source.," We find line fluxes for a subsample of 43 of the 64 objects, with an average of about two lines per source."
 When more han one value of the same line flux was found in the ilerature we either considered the most recent reference. which was likely to report values measured with more accurate techniques. or estimate the arithmetic average of he [lux (in almost all the cases the cülference was minimal).," When more than one value of the same line flux was found in the literature we either considered the most recent reference, which was likely to report values measured with more accurate techniques, or estimate the arithmetic average of the flux (in almost all the cases the difference was minimal)."
 We have not applied a reddening correction to the line D'uxes. cause the average data are not corrected.," We have not applied a reddening correction to the line fluxes, because the average data are not corrected."
 ii) We use the line ratios reported. by. Francis et al. (, ii) We use the line ratios reported by Francis et al. (
1991) (sce their Table or a complete list). mainly because of the large statistic and the range of lines incbuded*.,"1991) (see their Table 1 for a complete list), mainly because of the large statistic and the range of lines included."
. They refer to an optical sample. mostly. consisting of radio equie sources. but no major differences in the broad line Duxes between radio quiet and radio loud objects has been clearly determined. (e.g. Corbin 1992. see. also Steidel Sareen 1991. Boroson Creen 1992. Wills et al.," They refer to an optical sample, mostly consisting of radio quiet sources, but no major differences in the broad line fluxes between radio quiet and radio loud objects has been clearly determined (e.g. Corbin 1992, see also Steidel Sargent 1991, Boroson Green 1992, Wills et al."
 1993: Zheng et al., 1993; Zheng et al.
 1996)., 1996).
 iii) We then consider the sum of the line [uminosities of all the lines reported by. Francis et al. (, iii) We then consider the sum of the line luminosities of all the lines reported by Francis et al. (
1991) with respec o the Lva. to which we assign a reference. value of 100 (hereafter the asterisk refers to luminosities in the same units).,"1991) with respect to the $\alpha$, to which we assign a reference value of 100 (hereafter the asterisk refers to luminosities in the same units)."
" To this. we also add the contribution from Ly, (which is not inelucec in the list of Francis et al)."," To this, we also add the contribution from $\lha^*$ (which is not included in the list of Francis et al.),"
 with a value of 77 (from Gaskell ct al., with a value of 77 (from Gaskell et al.
 1981)., 1981).
 This gives a total μεις = 555.77~ 5.6 Ly...," This gives a total $\langle
\ll^* \rangle $ = $\sim$ 5.6 $\lla^*$ ."
 Therefore. given the sum of he observed luminosities in a certain numberof broad lines Nibioas. the total Leip can be calculatedas," Therefore, given the sum of the observed luminosities in a certain numberof broad lines $\Sigma_{\rm
i} L_{\rm i, obs}$ , the total $\ll$ can be calculatedas"
The KS test was then applied to both the distributions between the whole sample of GBM bursts and the 30 GBM bursts with measured redshift (for GRB 090519A and GRB 080928 a PL model fits the data best) and also to the distribution of bursts and the distribution of GBM bursts with measured redshift.,The KS test was then applied to both the distributions between the whole sample of GBM bursts and the 30 GBM bursts with measured redshift (for GRB 090519A and GRB 080928 a PL model fits the data best) and also to the distribution of bursts and the distribution of GBM bursts with measured redshift.
" In neither case the differences were statistically meaningful (P=24% and P=20%,, respectively)."," In neither case the differences were statistically meaningful $P=24$ and $P=20$, respectively)."
" Thus, all 3 histograms are very likely drawn from the same distribution."," Thus, all 3 histograms are very likely drawn from the same distribution."
" In conclusion, the sample of GBM GRBs with measured redshift presented here is representative for the whole population of GRBs which were ever observed by and GBM."," In conclusion, the sample of GBM GRBs with measured redshift presented here is representative for the whole population of GRBs which were ever observed by and GBM."
" However, it should be stressed that GBM cannot measure values which are lower than a certain limiting threshold."," However, it should be stressed that GBM cannot measure values which are lower than a certain limiting threshold."
" It is well known that values of GRBs can go as low as a few keV. ? for example find values as low as ~2 keV in GRBs that were observed by the High Energy Transient Explorer 2(HETE-2,, see e.g. ? and references therein)."," It is well known that values of GRBs can go as low as a few keV. \citet{pel08} for example find values as low as $\sim 2$ keV in GRBs that were observed by the High Energy Transient Explorer 2, see e.g. \citealt{hete2} and references therein)."
 These low energetic events have been classified as X-ray flashes (XRF) or X-ray rich bursts (XRB) (seee.g.??)..," These low energetic events have been classified as X-ray flashes (XRF) or X-ray rich bursts (XRB) \citep[see e.g.][]{heise01, sakamoto05}."
" However, it is very likely that XRFs and XRBs are nothing else than weak and long GRBs (see?,andreferencestherein)."," However, it is very likely that XRFs and XRBs are nothing else than weak and long GRBs \citep[see][and references therein]{kippen04}."
" The borderline value is obviously located somewhere near the low-energy sensitivity of the Nals, which has yet to be determined."," The borderline value is obviously located somewhere near the low-energy sensitivity of the NaIs, which has yet to be determined."
" Thus, in order to determine a potential bias in the distribution shown in Fig. 3,,"," Thus, in order to determine a potential bias in the distribution shown in Fig. \ref{fig:histoeprest},"
 it is important to understand and quantify the limits of the GBM to measure (be it either from the COMP model or Band function)., it is important to understand and quantify the limits of the GBM to measure (be it either from the COMP model or Band function).
" For this purpose, we created a set of simulated bursts with different initial spectral and temporal starting values."," For this purpose, we created a set of simulated bursts with different initial spectral and temporal starting values."
" We input the source lifetime (ts, 1 s, 5 s, 10 s, 100 s) and the photon flux (f) in the 10 keV to 1 MeV range (1, 3 and 10 ph cm? s!)."," We input the source lifetime $t_S$, 1 s, 5 s, 10 s, 100 s) and the photon flux $f$ ) in the 10 keV to 1 MeV range (1, 3 and 10 ph $^{-2}$ $^{-1}$ )."
" For the simulation the Band function was chosen as the photon model with varying (15,17, 25, 50, 100 keV) but fixed a=—0.8 and B=—2.4."," For the simulation the Band function was chosen as the photon model with varying (15,17, 25, 50, 100 keV) but fixed $\alpha=-0.8$ and $\beta=-2.4$."
 We simulate these bursts overlaid on real background data by using detector Nal 7 of GRB090926A‘., We simulate these bursts overlaid on real background data by using detector NaI 7 of GRB.
. This results in 60 different burst models., This results in 60 different burst models.
" For each model, we created 1000 bursts to account for Poissonian noise."," For each model, we created 1000 bursts to account for Poissonian noise."
" This results in 60000 spectra, each of which was then fitted with the Band function using the detector response matrix (DRM) of detector Nal 7 created for the location of GRB 0909264. After the fitting procedure, we reject those bursts which have ΔΕΡ/Ερ=0.3 and σαx0.4."," This results in 60000 spectra, each of which was then fitted with the Band function using the detector response matrix (DRM) of detector NaI 7 created for the location of GRB 090926A. After the fitting procedure, we reject those bursts which have $\Delta E_{\rm{p}} /{E_{\rm{p}}} \ge 0.3$ and $\sigma_\alpha \le 0.4$."
 These rejected spectra are then defined as unconstrained., These rejected spectra are then defined as unconstrained.
 We did not applythis criterion to the high-energy power law index f., We did not applythis criterion to the high-energy power law index $\beta$.
 A spectral fit that has a constrained and o but an unconstrained 6 is simply considered a COMP model., A spectral fit that has a constrained and $\alpha$ but an unconstrained $\beta$ is simply considered a COMP model.
" In Table 1 we report the mean and standard deviation of the output spectral parameters, and a, of the simulated bursts."," In Table \ref{tab:sim} we report the mean and standard deviation of the output spectral parameters, and $\alpha$ , of the simulated bursts."
 The conclusions of this exercise are:, The conclusions of this exercise are:
depth is below unity (the pairs set a diffusion time that is much larger than the dynamical time).,depth is below unity (the pairs set a diffusion time that is much larger than the dynamical time).
 On the other hand. (he rest lrame temperature in the expanding gas drops. and with it (he number of pairs and (he photon generation rate.," On the other hand, the rest frame temperature in the expanding gas drops, and with it the number of pairs and the photon generation rate."
 Therefore. the nmunber of photons in (he opaque expanding shells is roughly constant.," Therefore, the number of photons in the opaque expanding shells is roughly constant."
 The mildly relativistic temperature and the large optical depth ensure that photons and pairs interact many (limes (1.e.. sharing energy. annihilation and creation) over the dynamical time. thereby. keeping the pairs-photon gas in Compton pair equilibrium.," The mildly relativistic temperature and the large optical depth ensure that photons and pairs interact many times (i.e., sharing energy, annihilation and creation) over the dynamical time, thereby keeping the pairs-photon gas in Compton pair equilibrium."
 The radiation of the shell is confined to the gas during its expansion until the rest frame temperature drops enough so the pair loading becomes negligible., The radiation of the shell is confined to the gas during its expansion until the rest frame temperature drops enough so the pair loading becomes negligible.
" During"" (he expansion of a shell its rest frame temperature falls as 7""xV"",3 where V is the shell volume (in its rest frame) aud 122» is an effective adiabatic index."," During the expansion of a shell its rest frame temperature falls as $T' \propto V'^{-\beta}$, where $V'$ is the shell volume (in its rest frame) and $1+\beta$ is an effective adiabatic index."
 The value of 7 depends on temperature and it can drop slightly (by up to 30%)) below 1/3 because of pairs production and annihilation., The value of $\beta$ depends on temperature and it can drop slightly (by up to ) below $1/3$ because of pairs production and annihilation.
 ere we neglect this small deviation from the (vpical relativistic equation of state and approximate 3=1/3., Here we neglect this small deviation from the typical relativistic equation of state and approximate $\beta=1/3$.
 During the planar phase the volume of a fluid element grows as V'x//5., During the planar phase the volume of a fluid element grows as $V' \propto t/\g$.
 Therefore during acceleration Vxο77 while alter the acceleration ends Vx/.," Therefore during acceleration $V'
\propto t^{(3-\3)/2}$ while after the acceleration ends $V' \propto
t$."
" During the spherical phase Vx/ implving: The optical depth of the shell drops to unitv once (he pairs density. n2. is lower than the proton density. 5, (andtheir accompanied electrons). Le. ΠοπρI."," During the spherical phase $V' \propto t^3$ implying: The optical depth of the shell drops to unity once the pairs density, $n_\pm$, is lower than the proton density, $n_p$ (andtheir accompanied electrons), i.e., $n_\pm/n_p<1$."
" The initial value of this ratio is (e.g..Svensson1984:Budnikοἱal.2010): since pairs are in annihilation-creation equilibrium thei densitw drops exponentially with 7"" at this temperature range and the shell opacity is dominated by the electrons that were advected from the upstream once its temperature is: where the dependence on j is very weak."," The initial value of this ratio is \cite[e.g.,][]{Svensson84,Budnik10}: Since pairs are in annihilation-creation equilibrium their density drops exponentially with $T'$ at this temperature range and the shell opacity is dominated by the electrons that were advected from the upstream once its temperature is: where the dependence on $\g_0$ is very weak."
 Around this temperature. the breakout shell. and any less massive shell. becomes (iransparent and ils radiation escapes.," Around this temperature, the breakout shell, and any less massive shell, becomes transparent and its radiation escapes."
" We clefine /,544 as (he tme that the shell becomes optically thin. i.e.. its 7""= Ty."," We define $t_{th,0}$ as the time that the shell becomes optically thin, i.e., its $T'=T'_{th}$ ."
final metallicity of does not depend on their mass. because 3=f 0inthe starburst.,"final metallicity of does not depend on their mass, because $\beta=f=0$ in the starburst."
" In this case. we find that the final metallicity (Zr,~0.04 when Z2,~0—001."," In this case, we find that the final metallicity $\langle Z_{*}\rangle_{M_{*}}\sim 0.04$ when $Z_{cool}^{0}\sim
0-0.01$."
 ‘This value is almost the same as the final metallicities of ellipticals in Fig..., This value is almost the same as the final metallicities of ellipticals in \ref{fig:idcd}.
 Thus the results in the previous section are mostly explained by eqs.(16)) and (17))., Thus the results in the previous section are mostly explained by \ref{eqn:meanstz}) ) and \ref{eqn:meanstz2}) ).
 In this section. we investigate the parameter dependence of the CAIR on the star formation time-scale and the circular velocity dependence of the feedback.," In this section, we investigate the parameter dependence of the CMR on the star formation time-scale and the circular velocity dependence of the feedback."
 In Section ??.. we also investigate the clleet of the UV background radiation.," In Section \ref{sec:uv}, we also investigate the effect of the UV background radiation."
 In the previous section. the star formation time-scale. τι=20 Gye. is longer than the Llubble time.," In the previous section, the star formation time-scale, $\tau_{*}^{0}=20$ Gyr, is longer than the Hubble time."
 Next we investigate how the CMIt changes in the case of a short star formation⋅. time-scale., Next we investigate how the CMR changes in the case of a short star formation time-scale.
". We] adopt 7,=2 Gyr TNin this. section.", We adopt $\tau_{*}^{0}=2$ Gyr in this section.
. We show the CAIRs in the models I and E (Fig.S))., We show the CMRs in the models E and F \ref{fig:cmref}) ).
 The starburst model sbB is adopted in these mocdels., The starburst model sbB is adopted in these models.
 In the left panels (the model E). the dispersion becomes small compared. to the models A and € and the CALR reproduces the observations well.," In the left panels (the model E), the dispersion becomes small compared to the models A and C and the CMR reproduces the observations well."
 The colours of ellipticals come redder compared to those of the model A (Fig.2)) and those of bright ellipticals of the mocdel € (EFig.4)). and come bluer than those of dwarl ellipticals of the moclel C. Because the starburst model of the model LE is the same as that of the model €. these figures show that the star ormation time-scale is one of the essential parameters for he slope of the CM when we adopt the starburst mocoel sbDB. In the right. panels (the model E). luminous ellipticals are very red and the CAIR is lat at AdyX 20.," The colours of ellipticals become redder compared to those of the model A \ref{fig:cmrab}) ) and those of bright ellipticals of the model C \ref{fig:cmrcd}) ), and become bluer than those of dwarf ellipticals of the model C. Because the starburst model of the model E is the same as that of the model C, these figures show that the star formation time-scale is one of the essential parameters for the slope of the CMR when we adopt the starburst model sbB. In the right panels (the model F), luminous ellipticals are very red and the CMR is flat at $M_{V}\la-20$ ."
 The colour of the giant ellipticals nearly corresponds to the colour when Z.c0.04 and the age of galaxies is nearly equal to LO Car.," The colour of the giant ellipticals nearly corresponds to the colour when $Z_{*}\simeq
0.04$ and the age of galaxies is nearly equal to 10 Gyr."
 Moreover. we find that the dependence of CMIs on the starburst model becomes negligible.," Moreover, we find that the dependence of CMRs on the starburst model becomes negligible."
 Phe above results do not change when the starburst model sbj is adopted. (not shown)., The above results do not change when the starburst model sbA is adopted (not shown).
 These properties are discussed in Section ??.., These properties are discussed in Section \ref{sec:met}.
 In the same way as Section ??.. we investigate the metallicity and age of cllipticals next.," In the same way as Section \ref{sec:result}, we investigate the metallicity and age of ellipticals next."
 In order to understand the properties mentioned above. we show the metallicitv- and the age-luminosity cliagrams in Fig...," In order to understand the properties mentioned above, we show the metallicity- and the age-luminosity diagrams in \ref{fig:mtef}."
 Ages of ellipticals become older about 1-2 Cir than those inH the case ofH τιü=.20 Gyr1 on average., Ages of ellipticals become older about 1-2 Gyr than those in the case of $\tau_{*}^{0}=20$ Gyr on average.
 However. this dillerence corresponds to less than 0.1 mag in colour.," However, this difference corresponds to less than 0.1 mag in colour."
 We cannot explain the origin of the difference between the CMIBs in the models from A to D and the models E and E only by the age dillercnee., We cannot explain the origin of the difference between the CMRs in the models from A to D and the models E and F only by the age difference.
" The metallicity in the model E Fig.9((0)] increases in all range of magnitude compared to the model AX. Fig.3((c)]. and (Z,)5,.0.095 at Ay~23 and (Zr,~0.018 al Meo19."," The metallicity in the model E \ref{fig:mtef}( (c)] increases in all range of magnitude compared to the model A \ref{fig:mtab}( (c)], and $\langle Z_{*}\rangle_{L_{V}}\sim 0.028$ at $M_{V}\sim -23$ and $\langle Z_{*}\rangle_{L_{V}}\sim 0.018$ at $M_{V}\sim -19$."
 The reason of the metallicity. increase is considered: as. follows., The reason of the metallicity increase is considered as follows.
" Ln eq.(16)) it is only the case of fit,0ox that the final metallicity is determined by only £.", In \ref{eqn:meanstz}) ) it is only the case of $t/\tau_{*}\to\infty$ that the final metallicity is determined by only $F$.
 This is correct when the star formation time-scale is quite shorter than the Hubble time., This is correct when the star formation time-scale is quite shorter than the Hubble time.
 However. the value 7=20 Gyr is longer than the Llubble time and does not satisfy the above condition.," However, the value $\tau_{*}^{0}=20$ Gyr is longer than the Hubble time and does not satisfy the above condition."
" Phe mass-weighted mean stellar metallicity Z.(l)u, is à Monotonously increasing function about / in the deseription in the Section ??.."," The mass-weighted mean stellar metallicity $\langle
Z_{*}(t)\rangle_{M_{*}}$ is a monotonously increasing function about $t$ in the description in the Section \ref{sec:feedback}."
 So the actual final mean stellar metallicity is lower than 1.£7., So the actual final mean stellar metallicity is lower than $1-F$.
 Pherefore when we adopt a short star formation time-scale. mean stellar metallicities of galaxies increase totally and colours become redcder. compared to the model A. In contrast to the model A. in the model € with the same starburst model sbD as the model E. the metallicity is higher than that in the model LE. Ehe dillerence between the metallicities in the models C and I is emphasized at dwarl elliptical scale.," Therefore when we adopt a short star formation time-scale, mean stellar metallicities of galaxies increase totally and colours become redder, compared to the model A. In contrast to the model A, in the model C with the same starburst model sbB as the model E, the metallicity is higher than that in the model E. The difference between the metallicities in the models C and E is emphasized at dwarf elliptical scale."
 The reason that the metallicity of dwarf ellipticals is low in the model LE is considered. as follows., The reason that the metallicity of dwarf ellipticals is low in the model E is considered as follows.
 In the spiral state. the cold gas turns to stars in very. shor time-scale. 7;=2(1|2)3/2i6Gyr.," In the spiral state, the cold gas turns to stars in very short time-scale, $\tau_{*}=2(1+z)^{-3/2}$Gyr."
 In this star formation mode. supernova feedback can reheat the cold gas. anc most stars are formed in the spiral state owing to the shor," In this star formation mode, supernova feedback can reheat the cold gas, and most stars are formed in the spiral state owing to the short"
the disc of galaxy B is initially perpendicular to the disc of galaxy A (simulation 5).,the disc of galaxy B is initially perpendicular to the disc of galaxy A (simulation 5).
 The appearance of the velocity fields does of course differ for the different interaction scenarios., The appearance of the velocity fields does of course differ for the different interaction scenarios.
" However, the main conclusions drawn so far do not change."," However, the main conclusions drawn so far do not change."
" The merger is clearly visible in the asymmetric shape of the velocity field, the distortions are, however, smeared out at higher redshift."," The merger is clearly visible in the asymmetric shape of the velocity field, the distortions are, however, smeared out at higher redshift."
 Of course it is more probable to observe an interacting system in a pre- or post merger phase than in the relatively short period where the two galaxies permeate each other., Of course it is more probable to observe an interacting system in a pre- or post merger phase than in the relatively short period where the two galaxies permeate each other.
" Therefore, we study the same system 100 Myr after the direct encounter described above, i.e. after the first pericentre passage."," Therefore, we study the same system 100 Myr after the direct encounter described above, i.e. after the first pericentre passage."
 This is the period between the first passage and the complete merging of the two galaxies., This is the period between the first passage and the complete merging of the two galaxies.
 The small galaxy went through the gaseous disc of the massive galaxy and left a disturbed velocity field., The small galaxy went through the gaseous disc of the massive galaxy and left a disturbed velocity field.
 In the direct image the interaction is not visible anymore at redshift z=0.5., In the direct image the interaction is not visible anymore at redshift $z=0.5$.
 In Fig., In Fig.
 8 we show the distribution of the stellar mass in the galaxy as 2D image and as a radial profile.," \ref{direct_70}
 we show the distribution of the stellar mass in the galaxy as 2D image and as a radial profile."
 By assuming a certain constant mass-to-light ratio this could be translated to a light distribution., By assuming a certain constant mass-to-light ratio this could be translated to a light distribution.
 The profile can be fitted by an exponential law., The profile can be fitted by an exponential law.
 The velocity field shows still some signatures of the interaction (see Fig 9)) at lower redshift but at z—0.5 these are heavily smeared out., The velocity field shows still some signatures of the interaction (see Fig \ref{disturbed70}) ) at lower redshift but at z=0.5 these are heavily smeared out.
" Thus, the velocity information is an important complement to the morphological analysis when studying interactions of galaxies but only if the VF is sufficiently sampled."," Thus, the velocity information is an important complement to the morphological analysis when studying interactions of galaxies but only if the VF is sufficiently sampled."
 Kinematic disturbances from interactions are expected to fade within a few rotation cycles (€ 1 Gyr) (e.g. Dale et al., Kinematic disturbances from interactions are expected to fade within a few rotation cycles $\leq$ 1 Gyr) (e.g. Dale et al.
" 2001, Kronberger et al."," 2001, Kronberger et al."
 2006)., 2006).
 Only 200 Myr after the direct encounter presented above neither the direct image nor the VF show clear signatures of the interaction anymore., Only 200 Myr after the direct encounter presented above neither the direct image nor the VF show clear signatures of the interaction anymore.
" In a more detailed analysis, however, signs of the interaction can still be found."," In a more detailed analysis, however, signs of the interaction can still be found."
" About 1 Gyr after the first passage, when both galaxies are well separated but approaching each other again, galaxy A would be classified as undisturbed, even when observed at low redshift."," About 1 Gyr after the first passage, when both galaxies are well separated but approaching each other again, galaxy A would be classified as undisturbed, even when observed at low redshift."
" The analysis so far was made for the rather large galaxy A, which has a radial scale length of about 4.5 kpc."," The analysis so far was made for the rather large galaxy A, which has a radial scale length of about 4.5 kpc."
" Therefore, the velocity field is sampled by many pixels."," Therefore, the velocity field is sampled by many pixels."
" The situation is completely different for smaller galaxies, as e.g. model galaxy B. We investigate a ‘disturbed rotation’ of galaxy B using a snapshot of simulation 3, an equal mass merger of two small galaxies B. Equal mass mergers massively disturb the velocity field of the gas of the interacting galaxies."," The situation is completely different for smaller galaxies, as e.g. model galaxy B. We investigate a 'disturbed rotation' of galaxy B using a snapshot of simulation 3, an equal mass merger of two small galaxies B. Equal mass mergers massively disturb the velocity field of the gas of the interacting galaxies."
" Most of the gas is converted into stars by a merger induced starburst or lost to the intergalactic medium by tidal forces (e.g. Kapferer et al.,"," Most of the gas is converted into stars by a merger induced starburst or lost to the intergalactic medium by tidal forces (e.g. Kapferer et al.,"
 2005)., 2005).
" We choose a snapshot about 300 Myr after the first encounter, when again some regular rotation has established."," We choose a snapshot about 300 Myr after the first encounter, when again some regular rotation has established."
 In Fig., In Fig.
 10 the 2D velocity field of the galaxy for this snapshot is shown at redshift z=0.1 and at z=0.5 with an angular resolution of 0.4”.," \ref{equalvf} the 2D velocity field of the galaxy for this snapshot is shown at redshift $z=0.1$ and at $z=0.5$ with an angular resolution of 0.4""."
" While at z=0.1 the distortions in the VF are clearly visible, they are completely smeared out at z=0.5 leaving a regular VF, that would be classified as ’undisturbed rotation’."," While at $z=0.1$ the distortions in the VF are clearly visible, they are completely smeared out at $z=0.5$ leaving a regular VF, that would be classified as 'undisturbed rotation'."
" This classification is supported by the radial behaviour of the kinemetric quantities shown in Fig. 11,,"," This classification is supported by the radial behaviour of the kinemetric quantities shown in Fig. \ref{plotseq},"
 although at z=0.5 only two ellipses were fitted due to the small number of available pixels., although at $z=0.5$ only two ellipses were fitted due to the small number of available pixels.
 Such undetected distortions introduce an enormous source of systematic errors to distant Tully-Fisher studies., Such undetected distortions introduce an enormous source of systematic errors to distant Tully-Fisher studies.
 Therefore we conclude that current distant Tully-Fisher studies cannot give reliable results for low-mass systems if the velocity field is not sampled sufficiently., Therefore we conclude that current distant Tully-Fisher studies cannot give reliable results for low-mass systems if the velocity field is not sampled sufficiently.
 In Fig., In Fig.
" 12 we present the 2D velocity field of galaxy B for redshift z=0.05, z=0.2, z= 0.3, and z=0.4."," \ref{distevo2} we present the 2D velocity field of galaxy B for redshift $z=0.05$, $z=0.2$, $z=0.3$ , and $z=0.4$."
" Gradually the distortions are smeared out, leaving an almost regular VF at z=0.4."," Gradually the distortions are smeared out, leaving an almost regular VF at $z=0.4$."
 Also for this interaction scenario we investigate how long the 2D velocity field maintains the disturbed features., Also for this interaction scenario we investigate how long the 2D velocity field maintains the disturbed features.
 After the first pericentre passage of the two galaxies the velocity field settles again to an undisturbed state., After the first pericentre passage of the two galaxies the velocity field settles again to an undisturbed state.
" As for the unequal mass merger discussed above, the strongest features disappear after several hundred Myr while an undisturbed velocity field is again present after about 1 Gyr."," As for the unequal mass merger discussed above, the strongest features disappear after several hundred Myr while an undisturbed velocity field is again present after about 1 Gyr."
 Note that in this period the two galaxies are approachingeach other again for their, Note that in this period the two galaxies are approachingeach other again for their
however. there is a dramatic improvement between in the correspondence between pn=10.7 =0.1. and between n—1.2.,"however, there is a dramatic improvement between in the correspondence between $n=-1,
0$ , $n=0, 1$, and between $n=1, 2$."
 We can examine the morphology evolving from the same phase set by the characteristic scaleAy., We can examine the morphology evolving from the same phase set by the characteristic scale$k_{NL}$.
 For example. Avy Lor stage e corresponds to S computer erid units. thus the smoothing on scales bevond S grid. units erases. non-linearities. while the larger scale structure remains in the linear regime.," For example, $k_{NL}$ for stage $c$ corresponds to 8 computer grid units, thus the smoothing on scales beyond 8 grid units erases non-linearities, while the larger scale structure remains in the linear regime."
 The linear growth of the density [luctuations set up by the same initial phases depends only on time., The linear growth of the density fluctuations set up by the same initial phases depends only on time.
 The comparison therefore indicates the intrinsic. dillerence in clustering morphology arising from cdilferent initial power spectra., The comparison therefore indicates the intrinsic difference in clustering morphology arising from different initial power spectra.
 What we see on the scales bevond Ky; is that the correspondence is low between m=1 and à=1.2.," What we see on the scales beyond $k_{NL}$ is that the correspondence is low between $n=-1$ and $n=1, 2$."
 The correspondence deteriorates when the .dilference between the spectral indices increases. Le. the dillerence in intrinsic morphology is significant. or when the evolution goes into highly non-linear regime. where particles move away [rom their initial Lagrangian position and interact with each other non-linearlv.," The correspondence deteriorates when the difference between the spectral indices increases, i.e., the difference in intrinsic morphology is significant, or when the evolution goes into highly non-linear regime, where particles move away from their initial Lagrangian position and interact with each other non-linearly."
 Simulations evolving from the same initial phase configuration tend to develop nonlinear gaructures at or near the same spatial locations. but these structures appear with cillerent contrast when the initial spectra are dillerent.," Simulations evolving from the same initial phase configuration tend to develop nonlinear structures at or near the same spatial locations, but these structures appear with different contrast when the initial spectra are different."
 For example. filaments appear in the nonlinear regime in all cases. but for spectra with large n these tend to be less well defined. and. broken up into clumps.," For example, filaments appear in the nonlinear regime in all cases, but for spectra with large $n$ these tend to be less well defined and broken up into clumps."
 Our statistic S takes into account both the position and amplitude of structures that. form so it indicates a cleteriorating agreement for very cdillerent spectra., Our statistic $S$ takes into account both the position and amplitude of structures that form so it indicates a deteriorating agreement for very different spectra.
 Nevertheless. it is clear that there is strong imprint pon the morphology of the initial. phases resulting from 16 process of gravitational clustering.," Nevertheless, it is clear that there is strong imprint upon the morphology of the initial phases resulting from the process of gravitational clustering."
 This can also be seen visually in the pictures shown by Beacom et al. (, This can also be seen visually in the pictures shown by Beacom et al. (
1991).,1991).
 The cross-correlation coefficient being a single number. so one cannot infer [from it precisely how the phase-based reconstructions perform relative to the morphology of the original distribution.," The cross-correlation coefficient being a single number, so one cannot infer from it precisely how the phase-based reconstructions perform relative to the morphology of the original distribution."
 One intriguing question is how the density regions from reconstructions can be compared. to those from the original distributions., One intriguing question is how the density regions from reconstructions can be compared to those from the original distributions.
 Is it 0-x6° that produces high value of S7, Is it $\delta \propto \delta^{r}$ that produces high value of $S$?
 From rank-correlation coctlicicnt T. this is is not the case.," From rank-correlation coefficient $\tau$, this is is not the case."
 In order to get more information between the two distributions. we also made a grid-by-erid comparison between them.," In order to get more information between the two distributions, we also made a grid-by-grid comparison between them."
 A scatter diagram can he drawn in logarithmic scales for the comparison between each cell density from. phase-basecl reconstructions against the corresponding one from the original distribution., A scatter diagram can be drawn in logarithmic scales for the comparison between each cell density from phase-based reconstructions against the corresponding one from the original distribution.
" 1 ὃςXO'. the points scatter. along a straight line of slope equal to unity. anc the relative magnitude. of ""eventsis preserved. by a linear scale. factor."," If $\delta \propto \delta^{r}$, the points scatter along a straight line of slope equal to unity, and the relative magnitude of `events'is preserved by a linear scale factor."
 Fig., Fig.
 4 shows the scatter diagram. between original distribution and. phase-only. amplituce-swapped reconstructions lor nm=1 and n=].," \ref{scatter} shows the scatter diagram between original distribution and phase-only, amplitude-swapped reconstructions for $n=1$ and $n=-1$."
 Not all the points scatter along the straight line. as is suggested.," Not all the points scatter along the straight line, as is suggested."
 Instead. most. points scatter horizontally. particularly for phase-only cases.," Instead, most points scatter horizontally, particularly for phase-only cases."
 Also notice also that. in both reconstruction cases. the density values span only roughly one order of magnitude ancl the points that co not align well correspond to very. small Ductuations.," Also notice also that, in both reconstruction cases, the density values span only roughly one order of magnitude and the points that do not align well correspond to very small fluctuations."
" Ligh values of δν result. from phase configuration preserving the locations of high-density ""events! with large magnitudes in the original structure.", High values of $S$ result from phase configuration preserving the locations of high-density `events' with large magnitudes in the original structure.
 The relatively Lat) phase-only reconstructed densitv distribution is caused by the [lat power spectrum. in which the Fourier amplitude for cach mode is squasheck into unity.," The relatively flat phase-only reconstructed density distribution is caused by the flat power spectrum, in which the Fourier amplitude for each mode is squashed into unity."
 Lhe lower two panels are scatter clagrams of the amplitude-swapped reconstructions against the originals. Le. in Fi&.," The lower two panels are scatter diagrams of the amplitude-swapped reconstructions against the originals, i.e., in Fig."
 1 between (c) ancl (a) ancl between (d) and. (b)., \ref{demo} between (c) and (a) and between (d) and (b).
 The cdilference. is. amplituce-swapped reconstructions have the power spectra from the alternative sample distributions., The difference is amplitude-swapped reconstructions have the power spectra from the alternative sample distributions.
 The η=—1 amplitucle-swapped reconstruction has the power spectrum from the original distribution of &=1. which gives more power on small scales in the reconstruction. the p»=1 reconstruction. on the other hand. has more power on large scales.," The $n=-1$ amplitude-swapped reconstruction has the power spectrum from the original distribution of $n=1$, which gives more power on small scales in the reconstruction, the $n=1$ reconstruction, on the other hand, has more power on large scales."
 “Phis aclverse effect. causes the points of the low Illucetuations to spread on the scatter. diagram. which. however. doesn't decrease much. the cross-correlation coellicient S. between phasc-only and. amplitude-swapped reconstructions.," This adverse effect causes the points of the low fluctuations to spread on the scatter diagram, which, however, doesn't decrease much the cross-correlation coefficient $S$ between phase-only and amplitude-swapped reconstructions."
 This is in accord with with our visual impression that our eves pick up the maxima between the distributions for comparison., This is in accord with with our visual impression that our eyes pick up the maxima between the distributions for comparison.
 In Fig., In Fig.
 5. the scatter. diagrams for the realisations evolving from the same phase set are also produced., \ref{scatter1} the scatter diagrams for the realisations evolving from the same phase set are also produced.
 Llere only. (wo cases are chosen. Le. between n=1 and 0. and between n=l1and 2.," Here only two cases are chosen, i.e., between $n=-1$ and $0$, and between $n=-1$and $2$ ."
 The left two panels are comparison between n=l1 and n= 0. and the scattering is expected. which nonetheless follows the straight line.," The left two panels are comparison between $n=-1$ and $n=0$ , and the scattering is expected, which nonetheless follows the straight line."
 For the right panels. at early stage e. the deviation from the straight line," For the right panels, at early stage $c$ , the deviation from the straight line"
"Similarly, σεις","Similarly, _r^2."
was used.,was used.
 The rotation curve of the disk was computed following Casertauo(1993) aud Degeman(1987).., The rotation curve of the disk was computed following \citet{cas93} and \citet{beeg_phd}.
 The disk was asstuucd to have a vertical secl? distribution with a scale height zy=7/6 (wanderIxxuit&Searle1981)., The disk was assumed to have a vertical $^2$ distribution with a scale height $z_0 = h/6$ \citep{kruit81}.
. The rotation curves of the stellar component were resampled at the same radii as the ποσα curves., The rotation curves of the stellar component were resampled at the same radii as the smooth curves.
 We assume lis constant with radius., We assume is constant with radius.
" While oue expects some modest variation in Y, with radius (deJong1996). the color eradieuts in LSB galaxies tend to be small. so this effect is not likely to be significant."," While one expects some modest variation in $\MLstar$ with radius \citep{dejong_grad}, the color gradients in LSB galaxies tend to be small, so this effect is not likely to be significant."
 The surface density profiles prescuted im DMIT and vanderHulstetal.(1993) were used., The surface density profiles presented in BMH and \citet{vdh93} were used.
 They were scaled by a factor of 1.1 to take the contribution of hel aud metals iuto account., They were scaled by a factor of 1.4 to take the contribution of helium and metals into account.
 Their rotation curve was derived assumnidue the eas was distributed in a thin disk., Their rotation curve was derived assuming the gas was distributed in a thin disk.
 The gas rotation curves were resampled at the radii of the siiooth observed rotation curve., The gas rotation curves were resampled at the radii of the smooth observed rotation curve.
 The dark halo component differs from the previous two in that we are interested iu parametrizing this componucut assuniue some fiducial model., The dark halo component differs from the previous two in that we are interested in parametrizing this component assuming some fiducial model.
 The choice of this 110doel is the crus of most of the DM analyses in the literature. aud many models exist.," The choice of this model is the crux of most of the DM analyses in the literature, and many models exist."
 These can be broadly distinguished in two eroups: halo models with a core. aud. halo models with a cusp.," These can be broadly distinguished in two groups: halo models with a core, and halo models with a cusp."
 Au cxample of the first category is tlhe pseudo-isothermal halo. an example of the latter the CDM NEW halo.," An example of the first category is the pseudo-isothermal halo, an example of the latter the CDM NFW halo."
" As one of the goals of this paper is to assess the relevance of either category to the high-resolution LSB ealaxy rotation curves, we will present models derived using both models."," As one of the goals of this paper is to assess the relevance of either category to the high-resolution LSB galaxy rotation curves, we will present models derived using both models."
 We do realize there are many iutermediate mocels described iu the literature that probably can fit our data equally well., We do realize there are many intermediate models described in the literature that probably can fit our data equally well.
 However. our goal here is simply to see where the data lead us: is there a preference for models with a core or with a cusp?," However, our goal here is simply to see where the data lead us: is there a preference for models with a core or with a cusp?"
 We now describe the details of bothmodels., We now describe the details of bothmodels.
 The spherical pseudo-isothermal halo has a density profile where py is the central density of the halo. and ee: the core radius of the halo.," The spherical pseudo-isothermal halo has a density profile where $\rho_0$ is the central density of the halo, and $R_C$ the core radius of the halo."
 The corresponding rotation curve 1s given by The asvinptotic velocity of the halo. V4. is given by To characterize this halo only two of the three paramctors (pg.ReeVA) are needed. as equation (3) determines the value of the third parameter.," The corresponding rotation curve is given by The asymptotic velocity of the halo, $V_{\infty}$, is given by To characterize this halo only two of the three parameters $(\rho_0, R_C, V_{\infty})$ are needed, as equation (3) determines the value of the third parameter."
 The NEW iuass density distribution takes the form where Rods> the characteristic radius of the halo and pi is related to the deusitv of the universe at the time of collapse., The NFW mass density distribution takes the form where $R_s$ is the characteristic radius of the halo and $\rho_i$ is related to the density of the universe at the time of collapse.
 This mass distribution elves rise to a halo rotation curve where 65=RRogyy. It ds. characterized by a concentration parauecer e=Royy/Ry aud a radius Ag.," This mass distribution gives rise to a halo rotation curve where $x = R/R_{200}$ .It is characterized by a concentration parameter $c =
R_{200}/R_s$ and a radius $R_{200}$."
 These are directly related to Ry and p;. mt are used instead as they are a convenient wav to parametrize the rotation curve.," These are directly related to $R_s$ and $\rho_i$, but are used instead as they are a convenient way to parametrize the rotation curve."
 Tre radius Rogy is the radius where the density. contrast exceeds 200. roughly the virial radius (Navarro.Freuk&White1996).," The radius $R_{200}$ is the radius where the density contrast exceeds 200, roughly the virial radius \citep{NFW96}."
. The characteristic velocity Vou) of the halo is defined in the same way as Royy., The characteristic velocity $V_{200}$ of the halo is defined in the same way as $R_{200}$.
 These parauncters are not independent aud are set by the cosimolocsv., These parameters are not independent and are set by the cosmology.
 One of largest uncertainties in amv mass model is the value ofY., One of largest uncertainties in any mass model is the value of.
.. Though broad treuds in have been measured and modelled (c.g.Bottema1997:Bell&deJong 2000a).. the precise value for au individual ealaxy is not well kuown. aud depends ou extinction. star formation historv. Initial Mass Function. etc.," Though broad trends in have been measured and modelled \citep[e.g.][]{botje97,bell_pops}, the precise value for an individual galaxy is not well known, and depends on extinction, star formation history, Initial Mass Function, etc."
 Rotation curve fitting is a problem with too many free paramcters (vanAlbada&Sancisi1986:LakeFeinswoe1989) aud sone asstuuptions regarding must be made.," Rotation curve fitting is a problem with too many free parameters \citep{maxdisk86,lake89} and some assumptions regarding must be made."
" We therefore present disk-halo decompositions using four different assunptious for Y, for the galaxies in Sample I. For the galaxies in Sample IT ouly the mininmu-disk model is preseuted.", We therefore present disk-halo decompositions using four different assumptions for $\MLstar$ for the galaxies in Sample I. For the galaxies in Sample II only the minimum-disk model is presented.
 This uiel asstunes that the observed rotation curve is due cutirely to DM., This model assumes that the observed rotation curve is due entirely to DM.
 This gives an upper lit ou how concentrated the dark mass component can actually be aud is the versio1 of nininmun disk preferred in the CDM literature., This gives an upper limit on how concentrated the dark mass component can actually be and is the version of minimum disk preferred in the CDM literature.
 The contribution of the atomic eas and Πο) is taken iuto account. but lis asstuned to be zero.," The contribution of the atomic gas and He) is taken into account, but is assumed to be zero."
 This is the classical definition of niniuui disk as used in the rotation curve literature., This is the classical definition of minimum disk as used in the rotation curve literature.
 Tere lis set equal to à coustant value based on an Initial Mass Function aud a star formation listory appropriate for LSB ealaxies., Here is set equal to a constant value based on an Initial Mass Function and a star formation history appropriate for LSB galaxies.
 For the range in color 0.4«BoV<0.65 which LSB ealaxices normally exhibit (deBlok.vauderITulst.&Bothun1995) a value Y.(R}=Ll is a good estimate.," For the range in color $0.4 < B-V <
0.65$ which LSB galaxies normally exhibit \citep{edb_phot95} a value $\MLstar(R) = 1.4$ is a good estimate."
" For example. using the Bruzual&Charlot(1993) model with coustant star formation rate and Salpeter IME. we find that Y,(R)=L1 correspouds to BoV=0.16."," For example, using the \citet{ch_bruz} model with constant star formation rate and Salpeter IMF, we find that $\MLstar(R)=1.4$ corresponds to $B-V=0.46$."
 The model (Bocca-Volineraunge. cco.)," The model (Rocca-Volmerange, comm.)"
 gives avalue BoV=0.38. whereas the model by vields B.V— 0.67.," gives a value $B-V=0.38$, whereas the model by \citet{cole00} yields $B-V=0.67$ ."
" The models by Bell&ce give values around BVx0.6. The value Y,= L.Lis thusactually at the vlight-weieht™ eud of the plausible rauge. but this was deliberately chosen im order to eive niaxinmni opportunity for tlhe cuspy NEW models"," The models by \cite{bell_tfproc} give values around $B-V \simeq 0.6$ The value $\MLstar=1.4$ is thusactually at the “light-weight” end of the plausible range, but this was deliberately chosen in order to give maximum opportunity for the cuspy NFW models"
 , 
The new measurements show that the observed flares iu cdiffereut [frequencies operate under the,The new measurements show that the observed flares in different frequencies operate under the
mass function. then the SER. can be calculated following Ixennicutt(1998): The median FUR luminosity [or our sample is =210L. (the mean is ~S«107L.. with a large scatter) which is consistent with what has been found: previously for SMGs (e.g. Ixovaesetal. 2006)).,"mass function, then the SFR can be calculated following \citet{Kennicutt}: The median FIR luminosity for our sample is $\simeq2\times10^{12}\,\mathrm{L}_{\odot}$ (the mean is $\simeq8\times10^{12}\,\mathrm{L}_{\odot}$, with a large scatter), which is consistent with what has been found previously for SMGs (e.g. \citealt{Kovacs}) )."
 The median SER. for the 350 jpmi-observed SLLADES sources is 400M.vr.P as calculated from Equation 1..," The median SFR for the $350\,\mathrm{\mu m}$ -observed SHADES sources is $\simeq400\,\mathrm{M}_{\odot}\,\mathrm{yr}^{-1}$, as calculated from Equation \ref{SFR}."
 This is consistent with previous work. indicating that SMCs are significant contributors in the global star-Formation of the hieh-recshilt Universe (see E. Lillyetal.1999: Blainetal. 2002: [1 )5: Aretxagaetal. 2007)).," This is consistent with previous work, indicating that SMGs are significant contributors in the global star-formation of the high-redshift Universe (see e.g. \citealt{Lilly99}; \citealt{Blain}; \citealt{Chapman2005}; \citealt{paper4}) )."
 We can also constrain the cold dust mass. M4. in these obscured star-forming galaxies.," We can also constrain the cold dust mass, $M_\mathrm{d}$, in these obscured star-forming galaxies."
" The tux density of a galaxy at an observed frequeney. f. is given by the usual relation (e.g. Hughes.Dunlop&Rawlings 1997)): where Dp is the cosmological luminosity distance and D, is the Planck function evaluated at the emitted. frequency. smvafl|c)"," The flux density of a galaxy at an observed frequency, $\nu_\mathrm{obs}$, is given by the usual relation (e.g. \citealt{HDR97}) ): where $D_{\mathrm{L}}$ is the cosmological luminosity distance and $B_{\nu'}$ is the Planck function evaluated at the emitted frequency, $\nu'=\nu_\mathrm{obs}\,(1+z)$."
" The quantity αν is the wavelength- mass-absorption cocllicicnt (or ""elective area’ for Xxackbody emission by a certain mass of dust) and &4 can o extrapolated [rom an average iss,=2.64nrkg assuming 3=1.5 (Dunne.Eales&Edmunds20023)."," The quantity $\kappa_{\lambda}$ is the wavelength-dependent mass-absorption coefficient (or `effective area' for blackbody emission by a certain mass of dust) and $\kappa_{\lambda}$ can be extrapolated from an average $\kappa_{125_\mathrm{\mu
 m}}=2.64\,\mathrm{m}^{2}\,\mathrm{kg}^{-1}$ assuming $\beta=1.5$ \citep{Dunnekappa}."
. Using Equation 2.. the median dust mass implied. by he S50in observations is 95107M...," Using Equation \ref{mass}, the median dust mass implied by the $850\,\mathrm{\mu m}$ observations is $9\,\times10^{8}\,\mathrm{M}_{\odot}$."
 Uncertainties in he dust. masses are dominated by the uncertainty in the Xhotometrie redshifts Gvhen these are used). T4 and the lux densitv at S5Oμαι.," Uncertainties in the dust masses are dominated by the uncertainty in the photometric redshifts (when these are used), $T_\mathrm{d}$ and the flux density at $850\,\mathrm{\mu m}$."
 The uncertainty in αν is à factor ofa few and has not been included in the dust mass error xws. but the relative dust. masses in our sample will be correct i£ the same value holds for all SMCs (see Blainetal.2002 and references therein for a discussion).," The uncertainty in $\kappa_{\nu}$ is a factor of a few and has not been included in the dust mass error bars, but the relative dust masses in our sample will be correct if the same value holds for all SMGs (see \citealt{Blain} and references therein for a discussion)."
 Assuming that he maximum possible interstellar dust mass for a galaxy is about 1/500 of its total barvonic mass (see Eclmunels&Eales 1998)). the total barvonic mass of each SAIC is estimated to e Mise&5LOM ALS.," Assuming that the maximum possible interstellar dust mass for a galaxy is about 1/500 of its total baryonic mass (see \citealt{EdmundsEales}) ), the total baryonic mass of each SMG is estimated to be $M_\mathrm{bary}\simeq\!5\times10^{11}\,\mathrm{M}_{\odot}$ ."
 Por comparison. the barvonic mass of the Ιον Way is about LottAL.," For comparison, the baryonic mass of the Milky Way is about $10^{11}\,\mathrm{M}_{\odot}$."
. The best-fitting SEDs are given in Fig. 4..," The best-fitting SEDs are given in Fig. \ref{fig:fits2},"
 while Table eives the best-litting value of 24 for each SAIC. as well as the derived Lii. ία and SER.," while Table \ref{tab:tdm} gives the best-fitting value of $T_\mathrm{d}$ for each SMG, as well as the derived $L_\mathrm{FIR}$, $M_\mathrm{d}$ and SFR."
 Errors in the fit parameters (1.6. the position of the peak and the SED normalisation) are derived from the 68 per cent v2 confidence interval around the best-fitting values., Errors in the fit parameters (i.e. the position of the peak and the SED normalisation) are derived from the 68 per cent $\chi^{2}$ confidence interval around the best-fitting values.
 These errors have been carried. through to all the other derived «quantities in Fable 5.., These errors have been carried through to all the other derived quantities in Table \ref{tab:tdm}. .
 Obviously. for galaxies with only two FIR photometric points. the SEDs are not well-constrained and the minimuni reduced. X7. values are signilicantlv. less than one.," Obviously, for galaxies with only two FIR photometric points, the SEDs are not well-constrained and the minimum reduced $\chi^{2}$ values are significantly less than one."
 Photometric redshift errors have been folded into the uncertainties for derived quantities., Photometric redshift errors have been folded into the uncertainties for derived quantities.
 We assume that our adopted: model (a grevbody with ο= 1.5) is correct. and so the quoted. errors do not. include a contribution. from possible svstematic errors., We assume that our adopted model (a greybody with $\beta=1.5$ ) is correct and so the quoted errors do not include a contribution from possible systematic errors.
 Adopting a constant. value of :3 is euüivalent to assuming that the dust has similar properties in all galaxies., Adopting a constant value of $\beta$ is equivalent to assuming that the dust has similar properties in all galaxies.
 This is probably a reasonable assumption., This is probably a reasonable assumption.
 As additional precise photometric data become available (from for example) 23 will be able to be constrained. by direct fitting., As additional precise photometric data become available (from for example) $\beta$ will be able to be constrained by direct fitting.
 Note. however. that changing the value of ον +0.5 (between physically plausible values) has the effect of changing the derived dust temperature by about ES OMNI. which in turn allects the derived. luminosities of the SMCSs w c20 40 per cent.," Note, however, that changing the value of $\beta$ by $\pm{0.5}$ (between physically plausible values) has the effect of changing the derived dust temperature by about $\pm5$ K, which in turn affects the derived luminosities of the SMGs by $\pm20$ –40 per cent."
 The SALG population is therefore confirmed to x dominated bv. massive (—1017M... assuming οc 6). luminous (2«1077L. ) star-forming (SER 2400M.ve1 ) galaxies: with.. dust. temperatures o ~35lx.," The SMG population is therefore confirmed to be dominated by massive $>10^{12}\,\mathrm{M}_{\odot}$, assuming $M_\mathrm{tot}/M_\mathrm{baryons}\simeq6$ ), luminous $\simeq\!2\times10^{12}\,\mathrm{L}_{\odot}$ ) star-forming (SFR $\simeq400\,\mathrm{M}_{\odot}\,\mathrm{yr}^{-1}$ ) galaxies with dust temperatures of $\simeq35\,\mathrm{K}$."
" In contrast. local starburst galaxies with simular emperatures (Dunnectal.2000: 74,= 361xlx) are no usually classified as ULIRGs. since they are about. ten imes less luminous (Leu,~101 L.)."," In contrast, local starburst galaxies with similar temperatures \citealt{Dunne2000}: $\bar{T}_\mathrm{d}=36$ K) are not usually classified as ULIRGs, since they are about ten times less luminous $L_\mathrm{FIR}\sim10^{11}\,\mathrm{L}_{\odot}$ )."
 In addition. SMCs lave more than ten times more dust than local starburs galaxies. which could be explained by any of the following reasons: (1) SAIGs are hosted in galaxies ten times more massive than local starbursts: (2) SMGs are more gas- ane clust- rich for a given barvonicmass. being atan earlier stage in evolution with a higher eas fraction: or (3) the dus properties evolve with redshift and. high-redshift clust has a higher emissivity ancl so produces more submum emission for a given mass of dust.," In addition, SMGs have more than ten times more dust than local starburst galaxies, which could be explained by any of the following reasons: (1) SMGs are hosted in galaxies ten times more massive than local starbursts; (2) SMGs are more gas- and dust- rich for a given baryonicmass, being atan earlier stage in evolution with a higher gas fraction; or (3) the dust properties evolve with redshift and high-redshift dust has a higher emissivity and so produces more submm emission for a given mass of dust."
]xovacsetal.(2006). conducted. follow-up observations of bright. (25mv) racio-identified SCUBA sources with optical spectroscopic redshifts. including |LOCINSDQ.3. LOCINSSO.4. LOCINSDO.IS.. and τοςσσ.,"\citet{Kovacs} conducted follow-up observations of bright ${>}\,5\,\mathrm{mJy}$ ) radio-identified SCUBA sources with optical spectroscopic redshifts, including LOCK850.3, LOCK850.14, LOCK850.18, and LOCK850.30."
 Thev detected 12/15 of their sample with a mean z4.5 for the detections.," They detected 12/15 of their sample with a mean $\,{\simeq}\,4.5$ for the detections."
 Laurentetal.(2006). performed follow-up observations of 17 DBolocam l.lmunnme-selected. source candidates. in the LIE. with SILABRC-HI., \citet{Laurent06} performed follow-up observations of 17 Bolocam mm-selected source candidates in the LH with SHARC-II.
. Of. the 17. ten are. detected. including LOCIxS5O.1. LOCINS50.2. LOCINS50.3. LOCIxS50.12.. LOCINSS0.14.. LOCWS50.41. LOCINSDO.21. and (sce Laurentctal.2005. for counterpart id," Of the 17, ten are detected, including LOCK850.1, LOCK850.2, LOCK850.3, LOCK850.12, LOCK850.14, LOCK850.41, LOCK850.27, and (see \citealt{Laurent05} for counterpart )."
entifications)). ]xovacsctal.(2006) and Laurentetal.(2006). quote as a flux the peak flux density in an specified search radius around their target source position. which is taken from either racio or mm data.," \citet{Kovacs} and \citet{Laurent06} quote as a flux the peak flux density in an specified search radius around their target source position, which is taken from either radio or mm data."
 They do not deboost these Duxes. but their observations are at lower noise levels so this step is less important.," They do not deboost these fluxes, but their observations are at lower noise levels so this step is less important."
 Three SILADES sources. that we observe in this programme. LOCIx850.1.. LOCKS50.8 and οςΑΣAL have also been been observed by Kovacsetal.(2006). and Laurentetal.(2006).," Three SHADES sources that we observe in this programme, LOCK850.1, LOCK850.3 and LOCK850.41 have also been been observed by \citet{Kovacs} and \citet{Laurent06}."
. Because their observations are at lower noise we use their Duxes in Table 4. and in subsequent analvsis., Because their observations are at lower noise we use their fluxes in Table \ref{tab:sharc_photom} and in subsequent analysis.
 In the analysis section of this paper we use our best estimates of the 350pim Lux of cach source as photometric constraints.," In the analysis section of this paper we use our best estimates of the $350\,\mathrm{\mu m}$ flux of each source as photometric constraints."
 Llowever. for comparison to blank field surveys it is useful to consider how many of the sources are seen at a high chough S/N that they would be counted as detections.," However, for comparison to blank field surveys it is useful to consider how many of the sources are seen at a high enough S/N that they would be counted as detections."
 We have identified seven 350jam counterparts of 850jim. sources. as presented in Table >3..," We have identified seven $350\,\mathrm{\mu m}$ counterparts of $850\,\mathrm{\mu m}$ sources, as presented in Table \ref{tab:sharc_fluxes}."
 The positional olfsets of the detections from the SILXDISS positions are consistent with the positional distribution discussed above: 5/7 (71 per cent) of the detections lie within aarcsec., The positional offsets of the detections from the SHADES positions are consistent with the positional distribution discussed above; 5/7 (71 per cent) of the detections lie within arcsec.
 No trend is seen of decreasing posiional offset. with increasing S/N ratio of the seven associations. but this is not surprising given the low number statisics and the small dvnamic range in S/N. lt is striking hat the number ofdefeefed counterparts to SLLADES sources exactly matches the excess of positive peaks over noise (or negative peaks) in the full mapped area. as listed in Table 2..," No trend is seen of decreasing positional offset with increasing S/N ratio of the seven associations, but this is not surprising given the low number statistics and the small dynamic range in S/N. It is striking that the number of counterparts to SHADES sources exactly matches the excess of positive peaks over noise (or negative peaks) in the full mapped area, as listed in Table \ref{tab:peaks}."
 As we did with the number of peaks. we can use the corresponding composite negative catalogue to test if real associations are being found or just noisy peaks in the 350pim maps.," As we did with the number of peaks, we can use the corresponding composite negative catalogue to test if real associations are being found or just noisy peaks in the $350\,\mathrm{\mu m}$ maps."
" Counterparts to S50jii sources are searched for in an inverted map in the same wav. using a S/N threshold. of 2.5 ancl search radius of LOaaresec and no ""negative"" counterparts are identified."," Counterparts to $850\,\mathrm{\mu m}$ sources are searched for in an inverted map in the same way, using a S/N threshold of 2.5 and search radius of arcsec and no `negative' counterparts are identified."
 This result. reassures us that the 350/850jii associations are likely to be real.," This result reassures us that the $850\,\mathrm{\mu m}$ associations are likely to be real."
 Any of the Following mechanisms might cause an SM not to be detected at 350jm: (1) the 850jn source could be intrinsically less luminous. colder and/or have a cülferent SED shape than the strongly detected sources: (2) the SMG could [ie at a redshift in excess of about 3. this being more likely if there is also no radio counterpart. since the existing radio data are sullicientlv. deep to obtain a large fraction of counterparts for SMCs only for 23 (see Ivisonοἱal.2005)). and above z~3. 10 350jim band. is sampling the Wien side of the SED anc sulfers from. cosmological dimming without the benefit of the negative Ix-correction. or (3) The SAIC could. be spurious. although the false positive rate in the SILADIZS. οἳalogue is believed. to be very low.," Any of the following mechanisms might cause an SMG not to be detected at $350\,\mathrm{\mu m}$: (1) the $850\,\mathrm{\mu m}$ source could be intrinsically less luminous, colder and/or have a different SED shape than the strongly detected sources; (2) the SMG could lie at a redshift in excess of about 3, this being more likely if there is also no radio counterpart, since the existing radio data are sufficiently deep to obtain a large fraction of counterparts for SMGs only for $z\lesssim 3$ (see \citealt{Ivison05}) ), and above $z\sim 3$, the $350\,\mathrm{\mu m}$ band is sampling the Wien side of the SED and suffers from cosmological dimming without the benefit of the negative K-correction, or (3) The SMG could be spurious, although the false positive rate in the SHADES catalogue is believed to be very low."
 The catalogue was carclully constructed [rom a comparison of multiple data reductions in order to minimise false detections (see Coppinetal.2006))., The catalogue was carefully constructed from a comparison of multiple data reductions in order to minimise false detections (see \citealt{Coppin06}) ).
 We rely on the photometric luxes even for the 16 sources which are not detected. at. high significance., We rely on the photometric fluxes even for the 16 sources which are not detected at high significance.
 Therefore we have performed. à. stacking analysis to determine if a significantly. positive 350jim cux. density is associated with those sources.," Therefore we have performed a stacking analysis to determine if a significantly positive $350\,\mathrm{\mu m}$ flux density is associated with those sources."
 350pim Dux. densities ancl errors are measured. on the SUARC-LE maps at the radio position for the 11 of these 16 sources which have a radio counterpart and we find the mean flux is 16.7cdS manm.," $350\,\mathrm{\mu m}$ flux densities and errors are measured on the SHARC-II maps at the radio position for the 11 of these 16 sources which have a radio counterpart and we find the mean flux is $16.7 \pm 4.8$ mJy."
 This is our least biased. [lux estimator. and the agreement with our listed. deboosted κος is very encouraging.," This is our least biased flux estimator, and the agreement with our listed deboosted fluxes is very encouraging."
 The weighted mean deboosted lux for all 16 of these low significance sources is 16.643.8 mm.v., The weighted mean deboosted flux for all 16 of these low significance sources is $16.6 \pm 3.8$ mJy.
" We identify, one serendipitous 3.90 350ym source in our map of LOCW26/32: LOCIS350.1."," We identify one serendipitous $3.9\,\sigma$ $350\,\mathrm{\mu m}$ source in our map of LOCK26/32: LOCK350.1."
 LOCK350.1 is. not associated with any S50jii sources. since it is Z15arcsec away from any of the SLLADIES 850jn. positions in the map.," LOCK350.1 is not associated with any $850\,\mathrm{\mu m}$ sources, since it is $\gtrsim15\,\mathrm{arcsec}$ away from any of the SHADES $850\,\mathrm{\mu m}$ positions in the map."
 Based on Gaussian noise statistics and the number of independent beam sizes in the survey area with noise less thanντι. about 0.1 false positives are expected on average at as οἱ 3.9: LOCK350.1 istherefore likely a real blank-field 350iim source.," Based on Gaussian noise statistics and the number of independent beam sizes in the survey area with noise less than, about 0.1 false positives are expected on average ata S/N of 3.9; LOCK350.1 istherefore likely a real blank-field $350\,\mathrm{\mu m}$ source."
 “Phe position and [ux density of LOCK350.1 are.RA=110852™43°2.. (J2000) and 32.8£8.83mJ..," The position and flux density of LOCK350.1 are, (J2000) and $32.8\,\pm{8.3}\,\mathrm{mJy}$."
 This position corresponds to a region of positive [ux density in the S5Ojii map and a 2.36 peak is located &4arcsec from the position of LOCW350.1.," This position corresponds to a region of positive flux density in the $850\,\mathrm{\mu m}$ map and a $2.3\,\sigma$ peak is located $\simeq4\,\mathrm{arcsec}$ from the position of LOCK350.1."
 There also appears to be ada 1.4CGllIz source and a faint 24jm counterpart at these coordinates€ (Ivisonetal. 2007:: Ages&Lvison 2006:: see Fig. 2)).," There also appears to be a $4\,\sigma$ GHz source and a faint $24\,\mathrm{\mu m}$ counterpart at these coordinates \citealt{paper3}; ; \citealt{IvisonBiggs06}; ; see Fig. \ref{fig:lock350.1}) )."
 This source is therefore probably the third secure 350jin blank-field. detection (see μαςetal.2005 and Ixhanetal.2007 for the first two)," This source is therefore probably the third secure $350\,\mathrm{\mu m}$ blank-field detection (see \citealt{Khan05} and \citealt{Khan} for the first two)."
eas motions are dominated by stellar winds from. massive stars (sce Moelnick. Fenorio-Tagle. and Terlevich. 1999 for a recent review).,"gas motions are dominated by stellar winds from massive stars (see Melnick, Tenorio-Tagle, and Terlevich, 1999 for a recent review)."
 Lt is still not known if age is the only (or the dominant) parameter. or if environment also plavs an important role. but the fact that GIR with a wide range of ages fit the LOL?)—o relation suggests that the total mass ofthe objects is what determines m.," It is still not known if age is the only (or the dominant) parameter, or if environment also plays an important role, but the fact that GHR with a wide range of ages fit the $\lbeta-\sigma$ relation suggests that the total mass of the objects is what determines $\sigma$."
 The situation for HEEL galaxies is different., The situation for HII galaxies is different.
 Telles (1995) showed that these objects celine à. fundamental. plane that is remarkably similar to that. defined by norma elliptical galaxies., Telles (1995) showed that these objects define a fundamental plane that is remarkably similar to that defined by normal elliptical galaxies.
 This result. lends strong support to the interpretation of Terlevieh. ancl Melnick. (1981) and. NIE that the emission line-profile widths of Giant HIE galaxies directly measure the total mass of these svstems within the measuring radius., This result lends strong support to the interpretation of Terlevich and Melnick (1981) and MTM that the emission line-profile widths of Giant HII galaxies directly measure the total mass of these systems within the measuring radius.
 Therefore. besides systematic effects. tha are discussed below. the scatter in the L(IE2)—@ (notice that Telles used. continuum magnitudes and not L(1I:7)) depends among other things on the existence of à seconc parameter (see below). on possible variations of the LAL. on the importance of sources of broadening not related to a voung stellar component (c.g. rotation). and on the duration of the burst of star-formation that powers the emission lines.," Therefore, besides systematic effects, that are discussed below, the scatter in the $\lbeta-\sigma$ (notice that Telles used continuum magnitudes and not $\lbeta$ ) depends among other things on the existence of a second parameter (see below), on possible variations of the IMF, on the importance of sources of broadening not related to a young stellar component (e.g. rotation), and on the duration of the burst of star-formation that powers the emission lines."
 ALPML showed that this scatter can be reduced. by restricting the sample to objects with o«65km, MTM showed that this scatter can be reduced by restricting the sample to objects with $\sigma<65$.
s The same result was found by KooaL... (1995). for intermediate redshift objects.," The same result was found by Koo, (1995) for intermediate redshift objects."
 This cutoll can be understood if one assumes that HEEL ealaxies are powered by clusters of coeval stars (starbursts)., This cutoff can be understood if one assumes that HII galaxies are powered by clusters of coeval stars (starbursts).
 Phe cutolf results from imposing the condition that the time required for the clusters to form (e.g. the free-fall time) must be smaller than the life-time of the most massive stars., The cutoff results from imposing the condition that the time required for the clusters to form (e.g. the free-fall time) must be smaller than the life-time of the most massive stars.
 One of the two galaxies at z=8 plotted in Figure 2) appears to exceed this limit. but the measurement error (EPO 13) is still rather large.," One of the two galaxies at z=3 plotted in Figure \ref{lsigma} appears to exceed this limit, but the measurement error $\pm 20$ ) is still rather large."
 In order to minimize systematic effects due to the rapid evolution of the ionizing stars. NEM restricted their sample to galaxies with WiLL?)25A.," In order to minimize systematic effects due to the rapid evolution of the ionizing stars, MTM restricted their sample to galaxies with $\wbeta>25$."
. In fact. this restriction has à double purpose which is particularly relevant for high-z objects: it selects. the voung(est) starbursts. aud eliminates objects with significant underlving old(er) stellar populations.," In fact, this restriction has a double purpose which is particularly relevant for high-z objects: it selects the young(est) starbursts, and eliminates objects with significant underlying old(er) stellar populations."
 The latter is critical because an old. stellar population may widen the emission lines in a way that is ancorrelated with the luminosity of the voung component., The latter is critical because an old stellar population may widen the emission lines in a way that is uncorrelated with the luminosity of the young component.
 There are only a few objects in our intermecdiate recdshif sample with W113) 25A.., There are only a few objects in our intermediate redshift sample with $\wbeta>25$ .
 These are plotted with fillec symbols in Figure 2.., These are plotted with filled symbols in Figure \ref{lsigma}.
 Phe open symbols show the data for objects with weaker lines., The open symbols show the data for objects with weaker lines.
 As expected. these objects do no fit the correlation defined by the local LIL galaxies which have a mean line strength of <W113)>=105A.," As expected, these objects do not fit the correlation defined by the local HII galaxies which have a mean line strength of $<\wbeta>=105$."
. The luminosity evolution of a voung cocval starburs during the first 10vr. proceeds as a rapid. ἄοσαν of the emission line Hux after the first 3 Myr at roughly constan continuum flux until about 6 Myr., The luminosity evolution of a young coeval starburst during the first $10^7$ yr proceeds as a rapid decay of the emission line flux after the first 3 Myr at roughly constant continuum flux until about 6 Myr.
 Thus. in this range of ages the age-dimming in. LOL?) can be directly. estimated. [rom the change in equivalent widths (Lerlevich Moelnick. 1981: Copoetti. Pastoriza Dottori. 1986).," Thus, in this range of ages the age-dimming in $\lbeta$ can be directly estimated from the change in equivalent widths (Terlevich Melnick, 1981; Copetti, Pastoriza Dottori, 1986)."
 The mean equivalent width of the objects plotted as open squares in Figure 2.is, The mean equivalent width of the objects plotted as open squares in Figure \ref{lsigma} is
PolCor (Polarimeter and Coronograph) is a new combined imager. polarimeter. and coronagraph that provides sharp images (resolution down to 0722). and a well-defined point-spread function (PSF) resulting in a high image contrast.,"PolCor (Polarimeter and Coronograph) is a new combined imager, polarimeter, and coronagraph that provides sharp images (resolution down to 2), and a well-defined point-spread function (PSF) resulting in a high image contrast."
 In this paper we present a preliminary study to investigate its capability to image CSEs around AGB stars., In this paper we present a preliminary study to investigate its capability to image CSEs around AGB stars.
 The purpose is to study their structure and dynamical evolution., The purpose is to study their structure and dynamical evolution.
 We have observed three AGB stars using the PolCor instrument: the S-type binary AGB star W Agl. and the two detached shell sources DR Ser and U Cam.," We have observed three AGB stars using the PolCor instrument: the S-type binary AGB star W Aql, and the two detached shell sources DR Ser and U Cam."
 In Sect., In Sect.
 2. the imaging technique used in this work and previous investigations are briefly discussed., \ref{pol} the imaging technique used in this work and previous investigations are briefly discussed.
 In Sect., In Sect.
 3. the observations and the data reduction are described., \ref{obsdat} the observations and the data reduction are described.
 In Sect., In Sect.
 4. the observed sources are presented., \ref{sources} the observed sources are presented.
 In Sect., In Sect.
 5 the analysis is outlined., \ref{analys} the analysis is outlined.
 In Sect., In Sect.
 6. the results are presented and later discussed in Sect. 7.., \ref{resdis} the results are presented and later discussed in Sect. \ref{dis}.
 Finally. a summary is given and conclusions are drawn in Sect. 8..," Finally, a summary is given and conclusions are drawn in Sect. \ref{conc}."
 The PolCor instrument and its performance are presented in Appendix AppendixA:.., The PolCor instrument and its performance are presented in Appendix \ref{polcor}.
 When light is scattered by dust particles it becomes polarized., When light is scattered by dust particles it becomes polarized.
 The intensity and polarization of the scattered light can be used to determine the properties of the dust and to map the dust distribution., The intensity and polarization of the scattered light can be used to determine the properties of the dust and to map the dust distribution.
 The polarization degree is highest when the direction of the incident radiation is perpendicular to the viewing angle. but it will also depend on the wavelength. the grain size. and the grain composition.," The polarization degree is highest when the direction of the incident radiation is perpendicular to the viewing angle, but it will also depend on the wavelength, the grain size, and the grain composition."
 performed detailed calculations of the spectral polarization properties of optically thin and thick dust in different geometries., performed detailed calculations of the spectral polarization properties of optically thin and thick dust in different geometries.
 They found that the degree of linear polarization in the optical and near-infrared is high (=80% at 90° scattering angle) for wavelengths shorter than gm. decreases up to about m. stays constant at around or below and starts Increasing again long-ward of um. Only light that is scattered at an angle of 90° will be effectively polarized and polarized scattered stellar light hence probes the distribution of the dust in the plane of the sky.," They found that the degree of linear polarization in the optical and near-infrared is high $\approx$ at $^{\circ}$ scattering angle) for wavelengths shorter than $\mu$ m, decreases up to about $\mu$ m, stays constant at around or below and starts increasing again long-ward of $\mu$ m. Only light that is scattered at an angle of $90^{\circ}$ will be effectively polarized and polarized scattered stellar light hence probes the distribution of the dust in the plane of the sky."
 Imaging polarimetry has previously been proven to be a suitable technique. for mapping structures in. the close circumstellar environment around post-AGB stars(2)., Imaging polarimetry has previously been proven to be a suitable technique for mapping structures in the close circumstellar environment around post-AGB stars.
.. did ground-based near-infrared imaging polarimetry of 16 protoplanetary nebulae (PPNe)., did ground-based near-infrared imaging polarimetry of 16 protoplanetary nebulae (PPNe).
 They found that a large majority of their objects were extended in the polarized-intensity Images (as compared to the total-intensity images) showing that the objects are surrounded by dusty envelopes., They found that a large majority of their objects were extended in the polarized-intensity images (as compared to the total-intensity images) showing that the objects are surrounded by dusty envelopes.
 This work was then followed up by where 24 additional sources were observed., This work was then followed up by where 24 additional sources were observed.
 Also here polarization was detected in most of the observed sources., Also here polarization was detected in most of the observed sources.
 They suggest that the sources can be divided into objects with an optically thick disk (probably due to binary interaction) resulting in a bipolar morphology. or optically thin dust shells.," They suggest that the sources can be divided into objects with an optically thick disk (probably due to binary interaction) resulting in a bipolar morphology, or optically thin dust shells."
 found maximum polarization degrees up to in the bipolar objects., found maximum polarization degrees up to in the bipolar objects.
 High spatial resolution observations using NICMOS on the HST have also been used to investigate the morphologies of dust envelopes around PPNe(2). and found even higher polarization degrees around bipolar post-AGB stars.," High spatial resolution observations using NICMOS on the HST have also been used to investigate the morphologies of dust envelopes around PPNe, and found even higher polarization degrees around bipolar post-AGB stars."
 For stars on the AGB imaging polarimetry has been used when analyzing the large detached shells around carbon stars(??)., For stars on the AGB imaging polarimetry has been used when analyzing the large detached shells around carbon stars.
. To estimate the widths and radit of the detached shells as accurately as possible is important for a better insight into how the shells are formed., To estimate the widths and radii of the detached shells as accurately as possible is important for a better insight into how the shells are formed.
 It is also important in order to investigate whether the suggested He-shell flash scenario for the formation can be confirmed. and for getting a better understanding of the mass-loss processes during the thermally-pulsing AGB phase (TP-AGB).," It is also important in order to investigate whether the suggested He-shell flash scenario for the formation can be confirmed, and for getting a better understanding of the mass-loss processes during the thermally-pulsing AGB phase (TP-AGB)."
 Observations of the polarized light from the circumstellar environment of R Scl and U Ant have provided the possibility to measure the diameter and width of the dust shells with unprecedented accuracy., Observations of the polarized light from the circumstellar environment of R Scl and U Ant have provided the possibility to measure the diameter and width of the dust shells with unprecedented accuracy.
 The advantage of observing scattered stellar light in the optical in order to map circumstellar structure. especially compared to single-dish radio observations. is the high spatial resolution along with high-sensitivity detectors.," The advantage of observing scattered stellar light in the optical in order to map circumstellar structure, especially compared to single-dish radio observations, is the high spatial resolution along with high-sensitivity detectors."
 In the optical. the star-to-CSE flux ratio is - 10. requiring a very large dynamical range of the detector in order to image the weak emission from the CSE(?).," In the optical, the star-to-CSE flux ratio is $\sim$ $^{4}$ , requiring a very large dynamical range of the detector in order to image the weak emission from the CSE."
. Assuming that the direct light from the central star is essentially unpolarized. it should disappear in images of polarized light. reducing the required dynamical range of the detector.," Assuming that the direct light from the central star is essentially unpolarized, it should disappear in images of polarized light, reducing the required dynamical range of the detector."
 However. already a small amount of polarization will leave a significant signature of the stellar PSF also in the polarized images. and it is therefore sometimes necessary to also use a coronograph.," However, already a small amount of polarization will leave a significant signature of the stellar PSF also in the polarized images, and it is therefore sometimes necessary to also use a coronograph."
 The observations were carried out during three different observation runs at the mm Nordic Optical Telescope (NOT) on La Palma. The Canary Islands. using the PolCor instrument.," The observations were carried out during three different observation runs at the m Nordic Optical Telescope (NOT) on La Palma, The Canary Islands, using the PolCor instrument."
 In June 2006 we used a prototype of PolCor based on the same principles as the final instrument (described in detail in Appendix Appendix A:))., In June 2006 we used a prototype of PolCor based on the same principles as the final instrument (described in detail in Appendix \ref{polcor}) ).
 The observations are summarized in Table 1.., The observations are summarized in Table \ref{obs}.
 During each observation run we observed a number of polarization standard stars chosen from the lists of and?., During each observation run we observed a number of polarization standard stars chosen from the lists of and.
. The results of the calibration for the observations in July 2008 are shown in Fig. l.., The results of the calibration for the observations in July 2008 are shown in Fig. \ref{calibration}.
 The left panel shows the degree of polarization. and the right displays the polarization angle.," The left panel shows the degree of polarization, and the right displays the polarization angle."
 The formal error bars of the measurements are smaller than the symbols., The formal error bars of the measurements are smaller than the symbols.
 The scatter in the left panel is most likely due to time variations (most of the literature values are old) and/or it reflects differences in the effective wavelength., The scatter in the left panel is most likely due to time variations (most of the literature values are old) and/or it reflects differences in the effective wavelength.
 In the right panel a close relation between the observed polarization angle and that listed for the calibration stars is shown., In the right panel a close relation between the observed polarization angle and that listed for the calibration stars is shown.
 The off-set angle (22288) is due to the mechanical attachment of the instrument to the telescope., The off-set angle 8) is due to the mechanical attachment of the instrument to the telescope.
 The polarization of the NOT ts negligible (A. Djupvik. NOT Senior Staff Astronomer. private communication).," The polarization of the NOT is negligible (A. Djupvik, NOT Senior Staff Astronomer, private communication)."
 The “Lucky imaging’ technique was used to improve the image quality., The 'Lucky imaging' technique was used to improve the image quality.
 This means that only the sharpest frames were, This means that only the sharpest frames were
of the prior emission component. one needs to move the time zero point to ὧν. which we define as /=0.,"of the prior emission component, one needs to move the time zero point to $-T_\Delta$, which we define as $t=0$."
 Xccording to Ixobavashi Zhang (2007). the time zero point of external shock emission should coincide with the beginning of the corresponding central engine activity.," According to Kobayashi Zhang (2007), the time zero point of external shock emission should coincide with the beginning of the corresponding central engine activity."
 So for prior emission it should be at /=0. while for the main CRB it should beat Z7=05.," So for prior emission it should be at $t=0$, while for the main GRB it should be at $T=0$."
 The observed. X-ray afterelow is dominated bv the prior emission. component. which outshines the X-rav emission associated with the main GRB outflow during the shallow ancl normal decay phases.," The observed X-ray afterglow is dominated by the prior emission component, which outshines the X-ray emission associated with the main GRB outflow during the shallow and normal decay phases."
 A systematic study of NIE light curves from CRBs detected before mid-2009 (Liang et al., A systematic study of /XRT light curves from GRBs detected before mid-2009 (Liang et al.
 2009) suggests that the prior emission nmoclel provides a unified interpretation of both the canonical X-ray light curves and the light curves that seem to follow a single power law decay (which are the ones with negligibly short Z4)., 2009) suggests that the prior emission model provides a unified interpretation of both the canonical X-ray light curves and the light curves that seem to follow a single power law decay (which are the ones with negligibly short $T_\Delta$ ).
 In addition. Liang et al. (," In addition, Liang et al. ("
2009) also found that both tvpes of light curves in their sample can be roughly explained by the external shock mocol.,2009) also found that both types of light curves in their sample can be roughly explained by the external shock model.
 Ifa prior explosion indeed happened. this prior outllow must sweep up the ambient. medium and form an externa shock.," If a prior explosion indeed happened, this prior outflow must sweep up the ambient medium and form an external shock."
 The long-lasting prior X-ray emission that is invokes to interpret the X-ray. plateaus is most. likely from. this external shock., The long-lasting prior X-ray emission that is invoked to interpret the X-ray plateaus is most likely from this external shock.
 Lf this is the case. then the same externa shock would give rise to an optical emission. component that follows the same plateau behavior.," If this is the case, then the same external shock would give rise to an optical emission component that follows the same plateau behavior."
 This can be testec with the available optical data. especially during the earliest observational epoch.," This can be tested with the available optical data, especially during the earliest observational epoch."
 In. general. optical emission. observe during the prompt phase may be a mix of this prior emission component and several other emission components. including the optical counterpart of the prompt CRB emission. and the reverse shock and forward shock emission associated with the prompt outflow.," In general, optical emission observed during the prompt phase may be a mix of this prior emission component and several other emission components, including the optical counterpart of the prompt GRB emission, and the reverse shock and forward shock emission associated with the prompt outflow."
 As discussed in loka et al. (, As discussed in Ioka et al. (
2006). the ejecta associated with the prompt emission would stream into the trail of the prior component blastwave with modified medium density profile. leading to another pair of forward and reverse shocks (for the dynamics setting of the three shock system. see Zhang Alésszarros 2002).,"2006), the ejecta associated with the prompt emission would stream into the trail of the prior component blastwave with modified medium density profile, leading to another pair of forward and reverse shocks (for the dynamics setting of the three shock system, see Zhang Mésszárros 2002)."
 The relative importance of emission. contribution to the optical banc depends on many parameters., The relative importance of emission contribution to the optical band depends on many parameters.
 The. reverse shock component can dominate or be outshone bv. the [orward shock component in the optical band., The reverse shock component can dominate or be outshone by the forward shock component in the optical band.
 Phe detected prompt optical emission Εικ is therefore allowed to be higher than the predicted flux., The detected prompt optical emission flux is therefore allowed to be higher than the predicted flux.
 On the other hand. if the prompt optical ux or upper limit is already below the range of the predicted: optical Lux. then the external. shock prior emission model would be ruled out.," On the other hand, if the prompt optical flux or upper limit is already below the range of the predicted optical flux, then the external shock prior emission model would be ruled out."
 This is the motivation and strategy of this paper., This is the motivation and strategy of this paper.
 In order to create a sample. for this study. we need bursts with both X-rav. plateaus (which can be fit. within he prior emission model) ancl carly optical observations (cither detections or upper limits).," In order to create a sample for this study, we need bursts with both X-ray plateaus (which can be fit within the prior emission model) and early optical observations (either detections or upper limits)."
 We searched. through he literature to finc the GRBs which either clearly exhibit a shallow decay segment in their NICE light curve or have been it with a shallow decay segment., We searched through the literature to find the GRBs which either clearly exhibit a shallow decay segment in their XRT light curve or have been fit with a shallow decay segment.
 Specifically. the candidate oursts were taken from Table 1 of Liang et al. (," Specifically, the candidate bursts were taken from Table 1 of Liang et al. ("
2007) and Table 2 of Liane et al. (,2007) and Table 2 of Liang et al. (
2009).,2009).
 We further narrowed clown he sample to those GRBs with early optical observations. using Table 5 from Yost ct al. (," We further narrowed down the sample to those GRBs with early optical observations, using Table 5 from Yost et al. ("
2007a) ancl Table 2 from Yost et al. (,2007a) and Table 2 from Yost et al. (
2007b) for detections. and Table 3 from Yost et al. (,"2007b) for detections, and Table 3 from Yost et al. ("
2007b) for upper Limits.,2007b) for upper limits.
 Thus. the ellective cutoll date or our sample is December 2006. about two vears into the mission.," Thus, the effective cutoff date for our sample is December 2006, about two years into the mission."
 Altogether. our sample consists of eight bursts (Table 1) with a total of Ll early optical detections and 17 carly optical upper limits.," Altogether, our sample consists of eight bursts (Table 1) with a total of 11 early optical detections and 17 early optical upper limits."
 In the interest of comparing the optical light curves of these GRBs to those predicted: by he external shock prior emission miocdel. we have included as much of the available optical data for these bursts as ;xossible.," In the interest of comparing the optical light curves of these GRBs to those predicted by the external shock prior emission model, we have included as much of the available optical data for these bursts as possible."
" Lhe references. for the optical light curve data are given in ""Table. 1.", The references for the optical light curve data are given in Table 1.
 Since most of the prompt optical observations Crom Yost et al. (, Since most of the prompt optical observations from Yost et al. (
2007a.b) are in the A-band. we include only the ραπ optical light. curves.,"2007a,b) are in the -band, we include only the -band optical light curves."
 The only exception is CRB 060027. for which the included: prompt optical observations and optical light curve data are at à wavelength (SI90 21) near the zband (Ituiz-Velasco ct al.," The only exception is GRB 060927, for which the included prompt optical observations and optical light curve data are at a wavelength $8190~\mathring{A}$ ) near the -band (Ruiz-Velasco et al."
 2007)., 2007).
 Aeolore we can compare the prompt optical observations ancl optical light curves from the literature to the predictions of the external shock prior emission model. the optical data must be corrected for Galactic and host galaxy. extinction.," Before we can compare the prompt optical observations and optical light curves from the literature to the predictions of the external shock prior emission model, the optical data must be corrected for Galactic and host galaxy extinction."
 We corrected the optical data for Galactic extinction using he empirical Alilky Way (MW) extinction law. extinction coellicients. ancl Ay=3.08 from Pei (1992). along with the values of cle from Schlegel et al. (," We corrected the optical data for Galactic extinction using the empirical Milky Way (MW) extinction law, extinction coefficients, and $R_{V}=3.08$ from Pei (1992), along with the values of $A_B$ from Schlegel et al. ("
1998).,1998).
 In order to correct he optical data for host galaxy extinction. we searched the iterature on cach burst in our sample for the extinction law and ely or £g value which provided the best fit to its altcrglow spectral energy. distribution (SED).," In order to correct the optical data for host galaxy extinction, we searched the literature on each burst in our sample for the extinction law and $A_V$ or $E_{B-V}$ value which provided the best fit to its afterglow spectral energy distribution (SED)."
 We then used he appropriate results from the literature to calculate the rost galaxy extinction. once again relving on the empirical AIW. Large Magellanic Cloud (LAIC). and Small Magellanic Cloud (SAIC) extinction. laws ancl parameters from. Pei (1992).," We then used the appropriate results from the literature to calculate the host galaxy extinction, once again relying on the empirical MW, Large Magellanic Cloud (LMC), and Small Magellanic Cloud (SMC) extinction laws and parameters from Pei (1992)."
 The host galaxy extinction information and references or cach GRB are given in Table 1., The host galaxy extinction information and references for each GRB are given in Table 1.
 For CRB 060927 and CRB 0612224. the cited authors were unable to distinguish tween the three extinction laws.," For GRB 060927 and GRB 061222A, the cited authors were unable to distinguish between the three extinction laws."
 We chose the SAIC extinction law based on the conclusions of Schad et. al. (, We chose the SMC extinction law based on the conclusions of Schady et al. (
2007. 2010). that. in most cases. HE provides an acceptable it to the host galaxy extinction. profile.,"2007, 2010), that, in most cases, it provides an acceptable fit to the host galaxy extinction profile."
 It should. also be noted that the host galaxy extinction ofGIU 050401 is the subject of debate., It should also be noted that the host galaxy extinction of GRB 050401 is the subject of debate.
 We have decided to use the analysis of Watson et αἱ. (, We have decided to use the analysis of Watson et al. (
2006). although higher values of sly have been advocated by De Pasquale ct al. (,"2006), although higher values of $A_V$ have been advocated by De Pasquale et al. ("
2006) and Ixamble οἱ al. (,2006) and Kamble et al. (
2009).,2009).
" The calculated values (in magnitudes) of the host galaxy extinction zl,,,,, and Galactic extinction ον are eiven in Table 1.", The calculated values (in magnitudes) of the host galaxy extinction $A_{\lambda_{emit}}$ and Galactic extinction $A_{\lambda_{obs}}$ are given in Table 1.
" Phe error on zd,,,,, is due to propagating the error on the host galaxy. value of cy or Lev taken from the literature. ancl the error on Ze given by Pei (1992)."," The error on $A_{\lambda_{emit}}$ is due to propagating the error on the host galaxy value of $A_V$ or $E_{B-V}$ taken from the literature, and the error on $R_V$ given by Pei (1992)."
" Since the host galaxy chy values for GRB OSLIOOA and CARB OG61222A are limits. the errors on their ely), only take into"," Since the host galaxy $A_V$ values for GRB 051109A and GRB 061222A are limits, the errors on their $A_{\lambda_{emit}}$ only take into"
a distance of TOOpe and a temperature range 1000000019. For a low optical state the UV Luminosity reduces by a factor of 2Dto~1.5107to. which is still greater than the X-ray. luminosity bv a factor of 20.,"a distance of 700pc and a temperature range 10000–40000K. For a low optical state the UV luminosity reduces by a factor of 20 to $\sim1-5\times10^{32}$, which is still greater than the X-ray luminosity by a factor of 20."
 During the low state. the UV flux originates from the accreting white cwarl and the accretion disk.," During the low state, the UV flux originates from the accreting white dwarf and the accretion disk."
 Lt is therefore clillicult to determine the mass accretion rate (which is less than the mass transfer rate. cf.," It is therefore difficult to determine the mass accretion rate (which is less than the mass transfer rate, cf."
 Schoembs Hartmann 1983) during the low state., Schoembs Hartmann 1983) during the low state.
" However. if we use the high state UV. luminosity and £=CALMicRy. we obtain Alo~3105107 geís (—dL7—15.7LOM five. assuming A/,—0.6AZ... and a distance of 550-850. pc)."," However, if we use the high state UV luminosity and $L=GM_{1}\dot{M}_{\rm
acc}/R_{1}$, we obtain $\dot{M}_{\rm acc}\sim3-10\times10^{16}$ g/s $4.7-15.7\times10^{-10}$ /yr, assuming $M_{1}$ =0.6, and a distance of 550-850 pc)."
 Delove et al. (, Deloye et al. (
2007) predict. the mass aceretion rate as a function of the mass anc structure of the mass donor star and whether it is irracdated ancl find AM~520.Lor ος (which is slightly lower compared with that determined using the high state UV. luminosity) for an orbital period of 25 mins.,"2007) predict the mass accretion rate as a function of the mass and structure of the mass donor star and whether it is irradiated and find $\dot{M}_{\rm
acc}\sim5-20\times10^{15}$ g/s (which is slightly lower compared with that determined using the high state UV luminosity) for an orbital period of 25 mins."
 A greater understanding of the origin of the UV. emission during the outburst evele and the nature of the mass donor star is required. before the predicted: mass accretion rate can be usefully compared with observations., A greater understanding of the origin of the UV emission during the outburst cycle and the nature of the mass donor star is required before the predicted mass accretion rate can be usefully compared with observations.
 We model our optical lighteurve of KL Dra in the framework of the Disce Instability Model (hereafter DIM. c.g. Llameury et al.," We model our optical lightcurve of KL Dra in the framework of the Disc Instability Model (hereafter DIM, e.g. Hameury et al."
 1998) which has been adapted. for helium. dises by Lasota. Dubus Ixruk (2008). (," 1998) which has been adapted for helium discs by Lasota, Dubus Kruk (2008). ("
See Lasota 2001 for a review of the DIM).,See Lasota 2001 for a review of the DIM).
 We calculated: model lighteurves for. several dillerent. sets of parameters: ac (the viscosity parameter of the clise in the cold state). ag (the hot state). and the mass transler rate. Aly.," We calculated model lightcurves for several different sets of parameters: $\alpha_{\rm C}$ (the viscosity parameter of the disc in the cold state), $\alpha_{\rm H}$ (the hot state), and the mass transfer rate, $\dot{M}_{\rm tr}$."
 lo our simulations we fixed the mass of the accreting white dwarf at 0.6AL., In our simulations we fixed the mass of the accreting white dwarf at 0.6.
.. We also allowed the inner radius of the accretion dise. rj. to approach the radius of the white chwarl," We also allowed the inner radius of the accretion disc, $r_{\rm
in}$, to approach the radius of the white dwarf."
" Phe outer radius of the disc. mau. was allowed to vary around a mean of ray=1.2102""em."," The outer radius of the disc, $r_{\rm out}$, was allowed to vary around a mean of $r_{\rm
out}=1.2\times10^{10}\,\mathrm{cm}$."
" LorM our initialHEN simulations. we took Mi,t—3».107""l gs and we show a simulated light curve in Figure 5 using ας =0.025 and ay=0.026.", For our initial simulations we took $\dot{M}_{\rm tr}=3\times10^{16}$ g/s and we show a simulated light curve in Figure \ref{dim} using $\alpha_{\rm C}$ =0.025 and $\alpha_{\rm H}$ =0.026.
 It shows outbursts which repeat on à 60 cay timescale. an amplitude of ~3 mag and bright state lasting ~2 weeks.," It shows outbursts which repeat on a $\sim$ 60 day timescale, an amplitude of $\sim$ 3 mag and bright state lasting $\sim$ 2 weeks."
 In the simulations there is an increase in [lux of ~0.5 mag between the end of one outburst and the start of the next: this is a known celiciency. of the DIM (see comments in Lasota 2001)., In the simulations there is an increase in flux of $\sim$ 0.5 mag between the end of one outburst and the start of the next: this is a known deficiency of the DIM (see comments in Lasota 2001).
" In. contrast. if we assume AA,=l1.1y efs in our simulations then this low-state [ux increase is much greater."," In contrast, if we assume $\dot{M}_{\rm
tr}=1\times10^{17}$ g/s in our simulations then this low-state flux increase is much greater."
 Further. they give more asvnunetric outburst profiles (vith the decline being more extenced).," Further, they give more asymmetric outburst profiles (with the decline being more extended)."
 We also simulated. a set. of light. curves where we assumed that the secondary star was irradiated by X-rays and UM photons emitted in the accretion region and from he white dwark, We also simulated a set of light curves where we assumed that the secondary star was irradiated by X-rays and UV photons emitted in the accretion region and from the white dwarf.
 Llrradiation of the secondary can cause an enhancement of the mass transfer rate., Irradiation of the secondary can cause an enhancement of the mass transfer rate.
 We found. that although this gave values for the viscosity parameters similar o that of hvdrogen-dominated. accreting cdwarl novae (ce ac 00.02 and ay 0.1.0.2. Llameury ct al.," We found that although this gave values for the viscosity parameters similar to that of hydrogen-dominated accreting dwarf novae (eg $\alpha_{\rm C}\sim$ 0.02 and $\alpha_{\rm H}\sim$ 0.1–0.2, Hameury et al."
 1905. Smak 1999) 1ο simulated light. curves gave a large increase in xiehtness (over 1 mae) between the end of an outburst and he next.," 1998, Smak 1999) the simulated light curves gave a large increase in brightness (over 1 mag) between the end of an outburst and the next."
 A more detailed. investigation can help elucidate jow cdillerent. conditions (dillerent. critical values of X and T. smaller disc sizes) could. allect. cold. and hot. viscosity xwanmelters.," A more detailed investigation can help elucidate how different conditions (different critical values of $\Sigma$ and $T$, smaller disc sizes) could affect cold and hot viscosity parameters."
 Vhe AM. CVn systems CR Boo ancl VS03 Cen have been known for several decades and hence their optical light curves have been well studied they also have orbital periods close to that of KL Dra., The AM CVn systems CR Boo and V803 Cen have been known for several decades and hence their optical light curves have been well studied – they also have orbital periods close to that of KL Dra.
" The light curves of all three systems have shown a photometric signal on the ""uperhump' period (a signature of the precession period of", The light curves of all three systems have shown a photometric signal on the `superhump' period (a signature of the precession period of
"significant amounts of HD and will most likely cool to the temperature of the CMB before becoming Jeans-unstable, following the canonical pathway of Pop III.2 stars.","significant amounts of HD and will most likely cool to the temperature of the CMB before becoming Jeans-unstable, following the canonical pathway of Pop III.2 stars."
" Finally, triangles denote minihalos that have been disrupted by the SN remnant and are enriched to 1079Ze."," Finally, triangles denote minihalos that have been disrupted by the SN remnant and are enriched to $Z>10^{-6}~Z_{\odot}$ ."
" Two of these halos are even enriched to Z>10735Zo, and will almost certainly form Pop II stars."," Two of these halos are even enriched to $Z>10^{-3.5}~Z_{\odot}$ , and will almost certainly form Pop II stars."
" In Figure 8, we show a particularly interesting case that elucidates the importance of timing in regulating the strength of feedback."," In Figure 8, we show a particularly interesting case that elucidates the importance of timing in regulating the strength of feedback."
" Here, a minihalo has been partially photo-heated by a nearby star to 104K, and subsequently enriched with metals."," Here, a minihalo has been partially photo-heated by a nearby star to $10^{4}~{\rm K}$, and subsequently enriched with metals."
" In the process of its collapse, the core has fragmented into two distinct objects: one in which the molecule fraction has been significantly enhanced due to the photoionization, allowing the gas to cool to ~100K, and another in which the cooling is unaltered from the standard minihalo pathway."," In the process of its collapse, the core has fragmented into two distinct objects: one in which the molecule fraction has been significantly enhanced due to the photoionization, allowing the gas to cool to $\simeq 100~{\rm K}$, and another in which the cooling is unaltered from the standard minihalo pathway."
" In both cases, the metallicity is not high enough for metal fine-structure cooling to become important, although at higher densities dust-induced cooling may take over and lead to the formation of a small cluster of Pop II stars."," In both cases, the metallicity is not high enough for metal fine-structure cooling to become important, although at higher densities dust-induced cooling may take over and lead to the formation of a small cluster of Pop II stars."
 'This case suggests that the formation of low-mass stars need not await the formation of second-generation halos with My=105Mo., This case suggests that the formation of low-mass stars need not await the formation of second-generation halos with $M_{\rm vir}\ga 10^{8}~M_{\odot}$.
 What are the properties of the gas within the newly virialized galaxy?, What are the properties of the gas within the newly virialized galaxy?
" From Figures 5 and 6, it appears that ~10?Mo ofcold, dense gas at its center have been enriched to Z~10-?Zo."," From Figures 5 and 6, it appears that $\sim 10^5~M_{\odot}$ ofcold, dense gas at its center have been enriched to $Z\sim 10^{-3}~Z_{\odot}$."
" This is confirmed by Figure 9, where we show its distribution in density and temperature space within the virial radius of the galaxy."," This is confirmed by Figure 9, where we show its distribution in density and temperature space within the virial radius of the galaxy."
 It has become highly enriched by accretion from the surrounding IGM and is in a state of collapse., It has become highly enriched by accretion from the surrounding IGM and is in a state of collapse.
 We also find an interesting effect related to photoheating: previously ionized gas forms elevated H2 and HD abundances and coexists with unheated gas., We also find an interesting effect related to photoheating: previously ionized gas forms elevated $_{2}$ and HD abundances and coexists with unheated gas.
" Similar to the minihalo case discussed in Figure 8, the gas might therefore later fragment and form individual clumps."," Similar to the minihalo case discussed in Figure 8, the gas might therefore later fragment and form individual clumps."
" We note that in the density range resolved here, cooling is dominated by primordial molecules instead of metal fine-structure cooling etal.2007, 2009a)."," We note that in the density range resolved here, cooling is dominated by primordial molecules instead of metal fine-structure cooling \citep{jappsen07,jappsen09a}."
". At somewhat higher densities,(Jappsen we expect that the gas will cool to the temperature of the CMB before the equation of state hardens again."," At somewhat higher densities, we expect that the gas will cool to the temperature of the CMB before the equation of state hardens again."
" However, a more detailed investigation of the subsequent fragmentation must await high-resolution simulations that resimulate the central few kpc of the galaxy."," However, a more detailed investigation of the subsequent fragmentation must await high-resolution simulations that resimulate the central few kpc of the galaxy."
 We have performed a set of highly resolved SPH simulations that allow us to investigate the enrichment of the IGM by a PISN exploding in a high-redshift minihalo., We have performed a set of highly resolved SPH simulations that allow us to investigate the enrichment of the IGM by a PISN exploding in a high-redshift minihalo.
" We have incorporated a substantially higher degree of realism compared to our previous work in the form of radiation feedback, explicit chemical mixing, and metal line cooling."," We have incorporated a substantially higher degree of realism compared to our previous work in the form of radiation feedback, explicit chemical mixing, and metal line cooling."
" In particular, we have employed a physically motivated model for the mixing of metals between individual SPH particles based on diffusion (Greifetal. 2009a)."," In particular, we have employed a physically motivated model for the mixing of metals between individual SPH particles based on diffusion \citep{greif09a}."
". We have followed the distribution of metals as they are ejected into the IGM and then recollapse into the larger, My;~10°M potential well of a “first galaxy” assembling at z~ 10."," We have followed the distribution of metals as they are ejected into the IGM and then recollapse into the larger, $M_{\rm vir}\sim 10^{8}~M_{\odot}$ potential well of a “first galaxy” assembling at $z\simeq 10$ ."
 We have performed simulations with and without radiative feedback to assess the effects, We have performed simulations with and without radiative feedback to assess the effects
“EIT waves” were first observed in the solar corona using the Extreme ultraviolet Imaging Telescope (EIT:Mosesetal.1997) and analysed in detail by Thompsonetal.(1998)..,“EIT waves” were first observed in the solar corona using the )/Extreme ultraviolet Imaging Telescope \citep[EIT;][]{Moses:1997vn} and analysed in detail by \citet{Thompson:1998ab}.
 They have been a source of much controversy and debate in the solar physics community since this initial observation. with different authors suggesting that they are alternatively fast-mode MHD waves (e.g.:Wang2000:Warmuthetal.2004a:Long2008:Veronigetal.2008;Gopalswamy 2009).. a result of the re-structuring of the magnetic field during the eruption of a coronal mass ejection (CME:Chenetal.2002.2005:Attrill ora coronal MHD soliton (Wills-Daveyetal.2007)..," They have been a source of much controversy and debate in the solar physics community since this initial observation, with different authors suggesting that they are alternatively fast-mode MHD waves \citep[e.g.:][]{Wang:2000tg,Warmuth:2004rm,Long:2008eu,Veronig:2008ud,Gopalswamy:2009dn}, a result of the re-structuring of the magnetic field during the eruption of a coronal mass ejection \citep[CME;][]{Chen:2002rw,Chen:2005xe,Attrill:2006vn,Attrill:2007vn,Delannee:2007kx,Delannee:2008uq} or a coronal MHD soliton \citep{Wills-Davey:2007oa}."
" It should be noted at this point that the name “EIT wave"" is typically used for historical reasons.", It should be noted at this point that the name “EIT wave” is typically used for historical reasons.
 To reflect the uncertainty surrounding the physical interpretation of “EIT waves” we shall adopt the suggested nomenclature of Gallagher&Long(2010) and refer to these disturbances as coronal bright fronts (CBFs) throughout this paper., To reflect the uncertainty surrounding the physical interpretation of “EIT waves” we shall adopt the suggested nomenclature of \citet{Gallagher:2010ab} and refer to these disturbances as coronal bright fronts (CBFs) throughout this paper.
 The uncertainty in the nature of CBFs has arisen from conflicting results being drawn from. the same observations., The uncertainty in the nature of CBFs has arisen from conflicting results being drawn from the same observations.
 À pseudo-wave interpretation was proposed following observations of stationary bright fronts (DelannéeAulanier 1999).. a strong correlation with CMEs (Bieseckeretal.2002) and a lower than expected estimated pulse velocity (Wills-Daveyetal.2007)..," A pseudo-wave interpretation was proposed following observations of stationary bright fronts \citep{Delannee:1999ab}, a strong correlation with CMEs \citep{Biesecker:2002lq} and a lower than expected estimated pulse velocity \citep{Wills-Davey:2007oa}."
 However. observations of refraction (Wang2000:Ofman&Thompson2002;Veronigetal.2006) and reflection (Gopalswamyetal.2009) of CBFs at coronal hole boundaries would appear to suggest a wave interpretation.," However, observations of refraction \citep{Wang:2000tg,Ofman:2002ab,Veronig:2006fy} and reflection \citep{Gopalswamy:2009dn} of CBFs at coronal hole boundaries would appear to suggest a wave interpretation."
 A conclusive result has been hampered by the diffuse nature of CBFs and the relatively low temporal cadence of the observing instruments. both of which make it difficult to characterize their true nature.," A conclusive result has been hampered by the diffuse nature of CBFs and the relatively low temporal cadence of the observing instruments, both of which make it difficult to characterize their true nature."
 A full review of CBFs. their morphology. kinematics. relationship to other solar phenomena. and theoretical interpretations may be found in Wills-Davey&Attrill(2010) and Gallagher&Long(2010).," A full review of CBFs, their morphology, kinematics, relationship to other solar phenomena, and theoretical interpretations may be found in \citet{Wills-Davey:2010ab} and \citet{Gallagher:2010ab}."
 The first observations of CBFs were made using EIT with an effective cadence of ~ 12 minutes in the 195 ppassband., The first observations of CBFs were made using EIT with an effective cadence of $\sim$ 12 minutes in the 195 passband.
 Initial estimates of the kinematics of these disturbances. using a point-and-click methodology applied to running-difference Images. estimated the average velocity at ~189 km s! (Thompson&Myers2009)..," Initial estimates of the kinematics of these disturbances, using a point-and-click methodology applied to running-difference images, estimated the average velocity at $\sim$ 189 km $^{-1}$ \citep{Thompson:2009yq}."
 A higher velocity of 3114111 km s! was found by Warmuthetal.(20043) using additional passbands to compensate for the lack of 195 images., A higher velocity of $311\pm111$ km $^{-1}$ was found by \citet{Warmuth:2004rm} using additional passbands to compensate for the lack of 195 images.
 This is comparable to the range of Alfvénn speeds predicted by Wills-Daveyetal.(2007.~215-1500kms!)., This is comparable to the range of Alfvénn speeds predicted by \citet[][$\sim$215--1500~km~s$^{-1}$.
 The (STEREO: mission with the Extreme UltraViolet Imager (EUVI:etal.2004) instrument has led to new results., The \citep[\emph{STEREO}; mission with the Extreme UltraViolet Imager \citep[EUVI;][]{Wuelser:2004bs} instrument has led to new results.
 EUVI has an effective observing cadence of up to 1.5 minutes (ten times that of EIT). allowing for an improved estimate of the kinematies of these disturbances.," EUVI has an effective observing cadence of up to 1.5 minutes (ten times that of EIT), allowing for an improved estimate of the kinematics of these disturbances."
 A numerical differencing technique was applied by Longetal.(2008) to running-difference EUVI images to estimate a peak velocity range for a CBF of ~153 to 475 km s! with the acceleration of the disturbances estimated to be between —413 and 816 m s. depending on the cadence of the observations.," A numerical differencing technique was applied by \citet{Long:2008eu} to running-difference EUVI images to estimate a peak velocity range for a CBF of $\sim$ 153 to 475 km $^{-1}$ with the acceleration of the disturbances estimated to be between $-413$ and 816 m $^{-2}$, depending on the cadence of the observations."
 For the same event. Veronigetal.(2008) estimated the CBF velocity to be 460 km s! with an associated deceleration of 2160 m s by fitting a quadratic model to distance-time measurements.," For the same event, \citet{Veronig:2008ud} estimated the CBF velocity to be 460 km $^{-1}$ with an associated deceleration of $-160$ m $^{-2}$ by fitting a quadratic model to distance-time measurements."
 The results of Longetal.(2008) also indicated that the lower cadence of SOHO//EIT had resulted in the kinematics of CBFEs being previously underestimated. an observation confirmed by Veronigetal.(2008) and Maetal.(2009).," The results of \citet{Long:2008eu} also indicated that the lower cadence of /EIT had resulted in the kinematics of CBFs being previously underestimated, an observation confirmed by \citet{Veronig:2008ud} and \citet{Ma:2009fk}."
 The kinematics of CBFs provide an insight into the true nature of the disturbances., The kinematics of CBFs provide an insight into the true nature of the disturbances.
 Another useful physical indicator is the presence of pulse broadening. previously described by Warmuthetal.(2004b).," Another useful physical indicator is the presence of pulse broadening, previously described by \citet{Warmuth:2004ab}."
. Several authors have examined CBF pulse broadening using observations of multiple events from different passbands. including a combination of EUV and Ha observations (Warmuthetal.2001:Warmuth2010:Veronigal.2010) as well as observations of a disturbance in soft X- data from GOES//SXI (Warmuthetal.2005)..," Several authors have examined CBF pulse broadening using observations of multiple events from different passbands, including a combination of EUV and $\alpha$ observations \citep{Warmuth:2001ab, Warmuth:2010ab,Veronig:2010ab} as well as observations of a disturbance in soft X-ray data from /SXI \citep{Warmuth:2005vf}."
 In contrast. Wills-Davey(2006) observed no measurable increase m the FWHM of a pulse using high-cadence observations of a single CBE across the limited field-of-view of the (TRACE)).," In contrast, \citet{Wills-Davey:2006ab} observed no measurable increase in the FWHM of a pulse using high-cadence observations of a single CBF across the limited field-of-view of the )."
 This led Wills-Daveyetal. to propose that CBFs were soliton-like waves which exhibit no significant dispersion with propagation., This led \citet{Wills-Davey:2007oa} to propose that CBFs were soliton-like waves which exhibit no significant dispersion with propagation.
 Observations of CBFs typically show a decreasing front intensity with propagation., Observations of CBFs typically show a decreasing front intensity with propagation.
 This has been noted by many authors including Warmuthetal.(2001.2004b) and," This has been noted by many authors including \citet{Warmuth:2001ab, Warmuth:2004ab} and"
"The resulting Ks JJ-Ks colour-magnitude diagrams (CMD) for Wd1 and the off-field for all stars with DAOPHOT fitting errors <0.2 mmag are presented in refcmd,ll.",The resulting Ks J–Ks colour-magnitude diagrams (CMD) for Wd1 and the off-field for all stars with DAOPHOT fitting errors $\le$ mag are presented in \\ref{cmd_all}.
.T hedashedlineindicatesthesaturationlimit(seeT abl, The dashed line indicates the saturation limit (see Table \ref{obslog}) ).
"e fieldischaracterisedbyabluemain- sequenceof foreground stars, andaredsequenceof background Kignt sldeuteigt aeguatlaetpróatygai"," The off-field is characterised by a blue main-sequence of foreground stars, and a red sequence of background stars, most of which could be red giants located in the galactic bulge."
 hewd 1fi sequenceand pre—mainsequencemembersof theclusterinadditionseadecwop sparsatüth Jvilibleinfhemfifig. field., The Wd 1 field shows the main-sequence and pre-main sequence members of the cluster in addition to the two populations visible in the off-field.
Notethattheredbackground pear populationap stobelessreddenqqdclgsgato Wa ceumer eHatheastificial, Note that the red background population appears to be less reddened close to Wd 1 compared to the off-field.
 hisismostlikelyduetovaryingextinctionalongthelineo f Wd 1., This is most likely due to varying extinction along the line of sight towards the bulge at distances larger than the distance to Wd 1.
" In the following analysis, only stars with DAOPHOT photometric fitting errors of mmag or less are considered."," In the following analysis, only stars with DAOPHOT photometric fitting errors of mag or less are considered."
 Incompleteness simulations with the aim to assess the effect of crowding on our ability to detect faint sources were carried out both for the jitter combined off-field and the cluster field., Incompleteness simulations with the aim to assess the effect of crowding on our ability to detect faint sources were carried out both for the jitter combined off-field and the cluster field.
" In order not to change the crowding characteristics of the frames, for each run only 50 stars were added using under DAOPHOT."," In order not to change the crowding characteristics of the frames, for each run only 50 stars were added using under DAOPHOT."
" Artificial star magnitudes were allowed to scatter +0.5 mmag around a predefined value, and the stars were placed at random positions on the Ks-band frames."," Artificial star magnitudes were allowed to scatter $\pm$ mag around a predefined value, and the stars were placed at random positions on the Ks-band frames."
 The detection experiment was following the same steps and using the same PSF as the initial DAOPHOT analysis., The detection experiment was following the same steps and using the same PSF as the initial DAOPHOT analysis.
" Only those stars, whose recovered magnitudes match to within +0.5 mmag the input magnitude, were counted in the final analysis."," Only those stars, whose recovered magnitudes match to within $\pm$ mag the input magnitude, were counted in the final analysis."
" For each predefined magnitude, 10 such frames were created and analysed, i.e., 10x50 stars = 500 stars/mag bin."," For each predefined magnitude, 10 such frames were created and analysed, i.e., $\times$ 50 stars = 500 stars/mag bin."
" Then the magnitude was increased by 1.0mmag, and 10 new frames were created and analysed."," Then the magnitude was increased by mag, and 10 new frames were created and analysed."
" This procedure was repeated 6x, probing the magnitude range Ks = mmag to mmag for the cluster frame, and Ks = 13.9 to mag for the field frame."," This procedure was repeated $6 \times$, probing the magnitude range Ks = mag to mag for the cluster frame, and Ks = 13.9 to mag for the field frame."
" For the J-band incompleteness simulations, stars were placed at the same positions as on the Ks-band frames."," For the J-band incompleteness simulations, stars were placed at the same positions as on the Ks-band frames."
" J-band |)migeayges were computed assuming J-Ks = mmag, which is an intermediate colour between the (lower) main-sequence statarsnasWdstduleineb dlslbtirda stars ad"," J-band magnitudes were computed assuming J–Ks = mag, which is an intermediate colour between the (lower) main-sequence stars in Westerlund 1 with J–Ks $\approx$ mag, and the pre-main sequence stars with J–Ks $\approx$ mag."
ded were sighátgiacdistireeuaaatd, In total 240 frames with artificial stars added were analysed.
" arenaestaiaecdumhiedistanceto refinc,im.", The results are summarised in \\ref{inc_sim}.
.Notethatthelacko f brightstarsintheo f− fieldpushesthe50%completenes slimitto faintermagnitudescomparedtoti, Note that the lack of bright stars in the off-field pushes the completeness limit to fainter magnitudes compared to the cluster field.
" With Galactic coordinates 1 ~339.55? and b «—0.40°, Wd 1 is embedded in a rich population of fore- and background stars."," With Galactic coordinates l $\approx 339.55^\circ$ and b $\approx -0.40^\circ$, Wd 1 is embedded in a rich population of fore- and background stars."
" The more than 4000 stars detected in J and Ks in the comparison field with DAOPHOT fitting errors <0.2 mmag refcmd,ll, , right)giveagoode stimateof the field starpopulationmix."," The more than 4000 stars detected in J and Ks in the comparison field with DAOPHOT fitting errors $\le 0.2$ mag \\ref{cmd_all}, right) give a good estimate of the field star population mix."
T hisp ," This provides a firm basis to remove the field star contamination in the cluster CMD, and hence determine a clean, relatively unbiased cluster population for further analysis."
The statistical field subtraction is based on a comparison of the cluster and the field CMD., The statistical field subtraction is based on a comparison of the cluster and the field CMD.
 The CMDs are subdivided into grid cells with a step size of mmag in colour and magnitude., The CMDs are subdivided into grid cells with a step size of mag in colour and magnitude.
" The number of field stars within each cell is counted, normalised to the ratio of the sky areas covered by the image and the selected cluster annuli, and corrected for the differences in completeness fraction between the field and the cluster (see f refinc,im))."," The number of field stars within each cell is counted, normalised to the ratio of the sky areas covered by the image and the selected cluster annuli, and corrected for the differences in completeness fraction between the field and the cluster (see \\ref{inc_sim}) )."
"Finally, thesamenumbero f starsissubtractedatrandom f ro"," Finally, the same number of stars is subtracted at random from the corresponding grid cell in the cluster CMD."
"mti up(*fieldsubtracted"")clusterCM DsareshowninF emdiso..", Examples of the cleaned-up (“field subtracted”) cluster CMDs are shown in \\ref{cmd_iso}. .
selected from the redshift survey should be detected in (the lensing map.,selected from the redshift survey should be detected in the lensing map.
 To evaluate (he correspondence between systems in the redshift survey and significant peaks in (he convergence map. we lollow the approach of Geller οἱ al. (," To evaluate the correspondence between systems in the redshift survey and significant peaks in the convergence map, we follow the approach of Geller et al. ("
2010) for sampling the redshift survey.,2010) for sampling the redshift survey.
 We examine the redshift distribution in cones with radii of 3 and 6”., We examine the redshift distribution in cones with radii of $^\prime$ and $^{\prime}$ .
 The 3 sampling is similar to probes used in previous assessments of the efficiency of weak lensing based on counts (Schirmer et al., The $^\prime$ sampling is similar to probes used in previous assessments of the efficiency of weak lensing based on counts (Schirmer et al.
 2007). photometric redshilis (Gavazzi Soucail 2007) and spectroscopic redshifts (Geller et al.," 2007), photometric redshifts (Gavazzi Soucail 2007) and spectroscopic redshifts (Geller et al."
 2010)., 2010).
 The combination of 3 ancl 6’ probes samples (he virial radii of massive clusters throughout the redshilt range where we mieht expect the convergence map to detect these svstenis., The combination of $^{\prime}$ and $^{\prime}$ probes samples the virial radii of massive clusters throughout the redshift range where we might expect the convergence map to detect these systems.
 The 3° radius corresponds to 0.45 Mpe for z—0.15 and to 1.1 Alpe for z=0.53., The $^{\prime}$ radius corresponds to 0.45 Mpc for $z = 0.15$ and to 1.1 Mpc for $z = 0.53$.
" These radii are within reg). the radius where the enclosed average mass density. p(<Pou,=200p..."," These radii are within $_{200}$ , the radius where the enclosed average mass density, $\rho(<r)_{200} = 200\rho_c$."
" Ilere p, is the critical density.", Here $\rho_c$ is the critical density.
 The radius roo. a proxy for the virial radius. ranges from 1.0 to 2.0 Mpc for well-saanpled clusters with rest [rame Ime-of-sight. velocity dispersions in the range 500 km ! to 1000 kms | (Rines οἱ al.," The radius $r_{200}$, a proxy for the virial radius, ranges from 1.0 to 2.0 Mpc for well-sampled clusters with rest frame line-of-sight velocity dispersions in the range 500 km $^{-1}$ to 1000 km $^{-1}$ (Rines et al."
 2003: Rines Dialerio (2006)., 2003; Rines Diaferio (2006).
 According to the scaling relations of Rines Dialerio (2006). the corresponding range of masses. Map. within rao) is 1.3 x10! M. to 1.2x10P M...," According to the scaling relations of Rines Diaferio (2006), the corresponding range of masses, $_{200}$, within $_{200}$ is 1.3 $\times 10^{14}$ $_\odot$ to $\times 10^{15}$ $_\odot$."
 Our goal is [ist to assess theefficiency of (he convergence map in (he identilication of nassive Cluster halos., Our goal is first to assess the of the convergence map in the identification of massive cluster halos.
 Here we delineefficiency as the fraction of convergence map peaks with v>3.7 (the threshold chosen by Miyazaki οἱ al. (, Here we define as the fraction of convergence map peaks with $\nu > 3.7$ (the threshold chosen by Miyazaki et al. (
2007)) that. corresponds to an individual nassive cluster identilied within the redshift survey with a line-ol-sight velocity dispersion above the solid v=3.7 threshold curve of Figure 9..,2007)) that corresponds to an individual massive cluster identified within the redshift survey with a line-of-sight velocity dispersion above the solid $\nu=3.7$ threshold curve of Figure \ref{fig:sensitivity.ps}.
 In other words. the efficiency is the Traction of weak lensing peaks (hat correspond (to appropriately massive individual svstenis.," In other words, the efficiency is the fraction of weak lensing peaks that correspond to appropriately massive individual systems."
 We also seek an assessment of thecompleteness of the set of massive clusters among the veh confidence weak lensing candidate svstems., We also seek an assessment of the of the set of massive clusters among the high confidence weak lensing candidate systems.
 We deline as the fraction of all clusters in the foreground reclshilt survey with line-ol-sight velocity dispersions above the threshold curve that also correspond (o weak lensing peaks with »>3.7., We define as the fraction of all clusters in the foreground redshift survey with line-of-sight velocity dispersions above the threshold curve that also correspond to weak lensing peaks with $\nu > 3.7$.
 In other words. the completeness is the fraction of appropriately massive individual svstems in the redshift survev that are detected as weak lensing peaks.," In other words, the completeness is the fraction of appropriately massive individual systems in the redshift survey that are detected as weak lensing peaks."
 The delinilions ofefficiency anclcompleteness exclude superpositions of small groups from consideration., The definitions of and exclude superpositions of small groups from consideration.
 These superpositions may indeed produce a weak lensing signal. but the weak lensing peak should obviously not becounted as detection of a single massivesystem in a mass selected cluster catalog.," These superpositions may indeed produce a weak lensing signal, but the weak lensing peak should obviously not becounted as detection of a single massivesystem in a mass selected cluster catalog."
 The definitions also exclude coincidences between the, The definitions also exclude coincidences between the
 The processes of conversion of pulsar spiu-dowui energy iuto higl-enersv cussion in pulsar wind nebula are of ereat physical interest (e.g... 77:5 77).," The processes of conversion of pulsar spin-down energy into high-energy emission in pulsar wind nebula are of great physical interest (e.g., \cite{kc84}; ; \cite{che00}) )."
 Modern arcsec resolutiou imstruueuts such as ou aaud that are sensitive to photons up o 10 keV provide unique possibilities to study pulsar nebulae. because the non-thermal emissiou of the nebulae may be casily detected and studied iu this baud (e.g. Vela-X. ?7.. ??:: ολων 27: aud 3€58. ?77)). even when inuucrsed in the soft N-rav emission of the companion shell. as is often the Case.," Modern arcsec resolution instruments such as on and that are sensitive to photons up to 10 keV provide unique possibilities to study pulsar nebulae, because the non-thermal emission of the nebulae may be easily detected and studied in this band (e.g., Vela-X, \cite{hgh01}, \cite{pzs01}; G21.5-0.9, \cite{scs00}, \cite{wbb01}; and 3C58, \cite{bwm01}) ), even when immersed in the soft X-ray emission of the companion shell, as is often the case."
 Tudeed. the shell SNR IC. 113. once though to be mostly thermal in the N-ray baud (77: 73) has been discovered to enüt hard X-ray cussion by ??..," Indeed, the shell SNR IC 443, once though to be mostly thermal in the X-ray band \cite{pss88}; \cite{aa94}) ) has been discovered to emit hard X-ray emission by \cite*{wah92}."
 ASCA Cas TmaeineC»Oo Scitillator (CIS) observatious bv 77? discovered the localized character of the lard N-ray Cluission aud its non-thermal nature., ASCA Gas Imaging Scintillator (GIS) observations by \cite*{kpg97} discovered the localized character of the hard X-ray emission and its non-thermal nature.
 They concluded that WO of the 210 τον photons came from an isolated cnutting feature and from the South East clongatecd ridge of hard enmüssiou., They concluded that most of the 2–10 keV photons came from an isolated emitting feature and from the South East elongated ridge of hard emission.
 77 and ?? reported a hare component detected with the Phoswich Detector System (PDS) on BeppoSANX iud two compact N-ayv sources corresponding to the ASCA sources detected: with the BeppoSAX Aedimm-Eucrey Concentrator Spectrometer (MECS) (ISAN JOGL7.1|2221 aud ISAN JOGLS.0|2227)., \cite*{pre99} and \cite*{bb00} reported a hard component detected with the Phoswich Detector System (PDS) on BeppoSAX and two compact X-ray sources corresponding to the ASCA sources detected with the BeppoSAX Medium-Energy Concentrator Spectrometer (MECS) (1SAX J0617.1+2221 and 1SAX J0618.0+2227).
 Very recently. ISAN JOG1T.11:2] has been observed by aas reported by ??.. who also show a VLA observation at 1.16. [86 and 8.16 GIIz. aux a polarization measurement.," Very recently, 1SAX J0617.1+2221 has been observed by as reported by \cite*{ocw01}, who also show a VLA observation at 1.46, 4.86 and 8.46 GHz, and a polarization measurement."
 They argue that the hard radio spectral index. the amount of polarization and the overall N-vav aud racio morphology strongly sueeest that the source is a plerion nebula with a point source iu it. whose characteristic conietary shape is due to supersonic motion of the neutron star.," They argue that the hard radio spectral index, the amount of polarization and the overall X-ray and radio morphology strongly suggest that the source is a plerion nebula with a point source in it, whose characteristic cometary shape is due to supersonic motion of the neutron star."
 However. the limited counting statistics of the 10 ks observation do not allow a detailed spectral study of the jiebula. which is required to compare this new pleriou with current theoretical models.," However, the limited counting statistics of the 10 ks observation do not allow a detailed spectral study of the nebula, which is required to compare this new plerion with current theoretical models."
 Moreover. IC Τι is a vossible candidate for the πο EGRET x-ray source BEC JOGLT|2238 (77)) having a flux above 100 MeV. of ~H1* pls 1 tem? with a photon index of 2.01+1.06.," Moreover, IC 443 is a possible candidate for the CGRO EGRET $\gamma$ -ray source 3EG J0617+2238 \cite{hbb99}) ) having a flux above 100 MeV of $\sim 5 \times
10^{-7}$ ph $^{-1}$ $^{-2}$ with a photon index of $2.01 \pm 0.06$ ."
 More detailed iieasuremieuts of the nebula spectral xoperties are needed to study the relation between 3EC JOGLF|2238 and the nebula., More detailed measurements of the nebula spectral properties are needed to study the relation between 3EG J0617+2238 and the nebula.
 Iu thispaper an sstudy of the recently discovered plerion nebula in IC. 133 is presented., In this an study of the recently discovered plerion nebula in IC 433 is presented.
 Iu particular. we use the laree effective area of the EPIC instrmucut to address the svuchrotrom buru-off effect in the nebula. to resolve its structure and moeasure the dux. aud to constrain the thermal radiation of the central object in the nebula.," In particular, we use the large effective area of the EPIC instrument to address the synchrotron burn-off effect in the nebula, to resolve its structure and measure the flux, and to constrain the thermal radiation of the central object in the nebula."
 IC. 113 was observed as part of the Cal/PV phase of the OObservatorv (??7))., IC 443 was observed as part of the Cal/PV phase of the Observatory \cite{jla01}) ).
" ας woe have used the two observations centered ou 61772[5,3 |227267135 (T2000) performed ou 2000 September 27."," Here, we have used the two observations centered on $6^h17^m24^s.3$ $+22^d26^m43^s$ (J2000) performed on 2000 September 27."
 Data from the two MOS (27)) cameras and the PN (??)) camera are used., Data from the two MOS \cite{taa01}) ) cameras and the PN \cite{sbd01}) ) camera are used.
 The MOS aud PN cameras ave CCD arrays which collect rav photons between 0.1 and 15 keV aud lave a field of view of 30 diameter., The MOS and PN cameras are CCD arrays which collect X-ray photons between 0.1 and 15 keV and have a field of view of $30^\prime$ diameter.
" The pixel size is 1.1"" aud L1 for MOS and PN. respectively. while the mirror Point Spread Function is FFullwidth at half παπα,"," The pixel size is $1.1\arcsec$ and $4.1\arcsec$ for MOS and PN, respectively, while the mirror Point Spread Function is Full-width at half maximum."
 The data were acquired with the niediun filter and in full image mode. and therefore the temporal resolution is 2.5 s and 73 us for the MOS iux PN. respectively.," The data were acquired with the medium filter and in full image mode, and therefore the temporal resolution is 2.5 s and 73 ms for the MOS and PN, respectively."
 The poorer spatial resolution of the PN is compensated for wits ereater sensitivity. on the average more than the combined MOS caleras.," The poorer spatial resolution of the PN is compensated for by its greater sensitivity, on the average more than the combined MOS cameras."
 The Standard Analysis System (SAS) software used (version5.0.1. xiuusas-20001215) takes cares of most. of the required eveuts screening.," The Standard Analysis System (SAS) software used (version5.0.1, xmmsas-20001215) takes cares of most of the required events screening."
 However. we have further screened. the data to eliminate some residual hot pixels," However, we have further screened the data to eliminate some residual hot pixels"
"N,1 beige the nuuber o pairs of particles with separatious between r and r|Ar. V the volue considered. NV. the particle nuuber taken as centres. aud £p) the mean particle density.","$N_q$ being the number of pairs of particles with separations between $r$ and $r + \Delta r$, $V$ the volume considered, $N_c$ the particle number taken as centres, and $\langle \rho \rangle$ the mean particle density."
 We calculate this function at redshift Z=(0 for WD99 and DII aleorithius., We calculate this function at redshift $Z=0$ for WD99 and BH algorithms.
 The values we obtain are perfectly equal. and the substrucures we form Giuinber aud size) are identical.," The values we obtain are perfectly equal, and the substructures we form (number and size) are identical."
 To conclude oir WD99 description. we report the performances measured or the WD99 coe (Fie.," To conclude our WD99 description, we report the performances measured for the WD99 code (Fig."
 T) (1ucludingt10 boundary periodic conditions using the Ewald method )) and the perforLalces of the original DII aleoritlun., 7) (including the boundary periodic conditions using the Ewald method ) and the performances of the original BH algorithm.
 The imeastred performances lead us to the couclusion1 that when he system evolition Is οustered (Z= 0) the WD99 does not decrease the perornuauce as the BI algoritlin., The measured performances lead us to the conclusion that when the system evolution is clustered $Z=0$ ) the WD99 does not decrease the performance as the BH algorithm.
 This cffect is due to the natire of the WD99 algorithm. which las asructure that merevases the efficiency when clusters of paricles are well closed.," This effect is due to the nature of the WD99 algorithm, which has a structure that increases the efficiency when clusters of particles are well closed."
 This important eect alows us to n siuulations wihn verv Custered svsteus. obtainimg very good performance aiC uceligible errors.," This important effect allows us to run simulations with very clustered systems, obtaining very good performance and negligible errors."
 Moreover. the efficiency O the WD99 1icroases by a factor of up to five a the redshit Z--0.," Moreover, the efficiency of the WD99 increases by a factor of up to five at the redshift $Z=0$."
" The gun is enhanced when bieecr sinmilatious are run: a recent suumlation with 16 nüllion particles perormed on the Cray T3E system usine WD99 showed au increase in performace wea factor of 7 a the redshift Z="".", The gain is enhanced when bigger simulations are run: a recent simulation with 16 million particles performed on the Cray T3E system using WD99 showed an increase in performance by a factor of 7 at the redshift $Z=0$.
 We note that t1e gain obtained. m comparison wih that obtained by the original DII aleoritlin. is ercater using a lower critical level (5 or 6).," We note that the gain obtained, in comparison with that obtained by the original BH algorithm, is greater using a lower critical level (5 or 6)."
 The gain is increneuted using asphere critericn witli lower than the value we cousider. waving only a small iucremen in the eloba] error.," The gain is incremented using a Sphere criterion with lower than the value we consider, having only a small increment in the global error."
 The code WD99 is mainly used for LSS stilos. but it could be tested aud used for other," The code WD99 is mainly used for LSS studies, but it could be tested and used for other"
outbursts are eonerally attributed to a sizable increase in the dise accretion rate onto the stellar surface (Wartinann EKeuvou 1996).,outbursts are generally attributed to a sizable increase in the disc accretion rate onto the stellar surface (Hartmann Kenyon 1996).
 During the quiesceuce state FUors aud. EXors are normally accreting TTS. but because of thermal instability in the circiustellar disk accretion rate enhanced by a few orders of naenitude from ~LO TAL. up to —10 LAL./xi.," During the quiescence state FUors and EXors are normally accreting TTSs, but because of thermal instability in the circumstellar disk accretion rate enhanced by a few orders of magnitude from $\sim$ $^{-7}$$M_{\sun}$$/$ yr up to $\sim$ $^{-4}$$M_{\sun}$$/$ yr."
 FUor objects are characterized by AW IL-5 mag outburst amplitude. au F-C super-eiaut «ροζπα diviug outbursts association with reflection nebulae. location iu forming regions. a stroue Lil 6707 jin absorption. Πα aud Na I 58590. 5896 dadisplaviug νο profiles (Πας 1977: Reipiurth Aspin 2010).," FUor objects are characterized by $\Delta$$V$$\approx$ 4-5 mag outburst amplitude, an F-G super-giant spectrum during outbursts, association with reflection nebulae, location in star-forming regions, a strong LiI 6707 in absorption, $\alpha$ and Na I 5890, 5896 displaying P-Cyg profiles (Herbig 1977; Reipurth Aspin 2010)."
 The outburst of FUor objects last for several decades. iux the rise time is faster than the decline.," The outburst of FUor objects last for several decades, and the rise time is faster than the decline."
 EXor objects shows frequent (every few voars or decade). irregular aud relatively brief (à few mouths to a few veas) outburst of several magnitudes amplitude(AVz3-5).," EXor objects shows frequent (every few years or decade), irregular and relatively brief (a few months to a few years) outburst of several magnitudes amplitude."
 During such events. the cool spectra of the quiescence is veiled. aud strong cussion lines from single ionized metals are observed together with appearance of reversed P-C'e absorption componcuts (Herbig 2007).," During such events, the cool spectrum of the quiescence is veiled, and strong emission lines from single ionized metals are observed together with appearance of reversed P-Cyg absorption components (Herbig 2007)."
 A significant part of ITAEBE stars with a spectral type later than AO show stroug photometric variability with suddeu quasi-Aleolc» drops in briglituess aud amplitudesge up to 225 (V) (Natta et al., A significant part of HAEBE stars with a spectral type later than A0 show strong photometric variability with sudden quasi-Algol drops in brightness and amplitudes up to $2\fm5$ $V$ ) (Natta et al.
 1997. vau den Aucker et al.," 1997, van den Ancker et al."
 1998)., 1998).
 During the deep nuüuimauuus of brghtuess. au Increase da polarizaticueji and specific color variability are observed.," During the deep minimums of brightness, an increase in polarization and specific color variability are observed."
 The xototvpe of this eroup of PAIS objects with intermediate mass nane UXors is UN Orionis., The prototype of this group of PMS objects with intermediate mass named UXors is UX Orionis.
 The senueral explanation of its variability is a variable extinction from dust clumps of fibuunceuts passing through the line of sight (Cuiuin ot al., The general explanation of its variability is a variable extinction from dust clumps of filaments passing through the line of sight (Grinin et al.
 1991. Dullemoud et al.," 1991, Dullemond et al."
 2003)., 2003).
 The variable star CAL Cep was discovered: seveu decades ago on the photographic plates frou Souneberg observatory (Morgeeuroth 1939). but the exact icceliauiiulj of variability remains still παον discussion (Niao et al.," The variable star GM Cep was discovered seven decades ago on the photographic plates from Sonneberg observatory (Morgenroth 1939), but the exact mechanism of variability remains still under discussion (Xiao et al."
 2010)., 2010).
" CAL Cop lie iu the field of the voune open cluster ""Truupler 37 (£L Myr old) aud most likely is a 11e1iber of the cluster (Alarschall van Altena 1987. Sicilia-Aecuilar et al."," GM Cep lie in the field of the young open cluster Trumpler 37 $\sim$ 4 Myr old) and most likely is a member of the cluster (Marschall van Altena 1987, Sicilia-Aguilar et al."
 2005)., 2005).
 According to Sicilia-Aguilar et al. (, According to Sicilia-Aguilar et al. (
2008) GM Cop is a solar niass-star CAL — 2.1M.) from CTV-IKOV. spectral type. with radius between 3 aud 6 Rand with a strong IR excesses frou a very hues civetuustellar disk.,"2008) GM Cep is a solar mass-star $\it M$ $\sim$ $2.1 M_\sun$ ) from G7V-K0V spectral type, with radius between 3 and 6 $R_\sun$ and with a strong IR excesses from a very luminous circumstellar disk."
 The physical paramcters (uass. radius. spectral type) of GAL Cop defined it as member of the subclass of ETTS.," The physical parameters (mass, radius, spectral type) of GM Cep defined it as member of the subclass of ETTS."
 The spectimm of CAL Cep is dominated by a very strong and broad Wa eimission line with a strong P Cre profile (Sicilia-Aeuilar et al., The spectrum of GM Cep is dominated by a very strong and broad $\alpha$ emission line with a strong P Cyg profile (Sicilia-Aguilar et al.
 2008)., 2008).
 The equivalent width of Πα line vary significantly from 6.4 (2001 Jun.) to 1944 (2007 Apr.)., The equivalent width of $\alpha$ line vary significantly from $\AA$ (2001 Jun.) to $\AA$ (2007 Apr.).
 But the lack of photometric observations around these dates do impossible to determine a correlation between the brightness of the star and the equivalent width of Ilo line., But the lack of photometric observations around these dates do impossible to determine a correlation between the brightness of the star and the equivalent width of $\alpha$ line.
 The presence of a massive circumstellar disk and variable accretion rates (up to ~ 9 AL. /year) are also detected in the study of Sicilia-Aeuilar ct al. (, The presence of a massive circumstellar disk and variable accretion rates (up to $\sim$ $^{-6}$ $M_\sun$ /year) are also detected in the study of Sicilia-Aguilar et al. (
2008).,2008).
 Ou the basis of photographic monitoring Suvirkova L975) classified CAL Cep as a RW Aur variable (also known as extreme CTTSs)., On the basis of photographic monitoring Suyarkova (1975) classified GM Cep as a RW Aur variable (also known as extreme CTTSs).
 The registered photographic iuplitude for a nine vears period is 2.2 mae., The registered photographic amplitude for a nine years period is 2.2 mag.
 Similar brightness variations are reported by Kun (1986) (the ςserved amplitude is AV 22215)., Similar brightness variations are reported by Kun (1986) (the observed amplitude is $\Delta V$ $2\fm15$ ).
 The firs multicolor photometric study of CAL Cep based ou optical aud infrared observations was done by Sicilia-Agnilar et i (, The first multicolor photometric study of GM Cep based on optical and infrared observations was done by Sicilia-Aguilar et al. (
2008).,2008).
 The authors found the star much. brighter iu 2006 than in 2000 aud conclude that the most probable reason for brightuess increase is au outburst fomir EXor type., The authors found the star much brighter in 2006 than in 2000 and conclude that the most probable reason for brightness increase is an outburst from EXor type.
 Sicilia-Aguilar ct al. (, Sicilia-Aguilar et al. (
2008) areue their hwpothesis with the presence of hDnuuinous imuil-IR disk and with the high aud variable accretion rate.,2008) argue their hypothesis with the presence of luminous mid-IR disk and with the high and variable accretion rate.
 A long-term photometric study of GM. Cep for severa decades period was performed bv Xiao et al. (, A long-term photometric study of GM Cep for several decades period was performed by Xiao et al. (
2010).,2010).
 The photographic plate archives frou IEuvaird. College Observatory aud from Sounebere Observatory are usec construct the long-term D aud V light curves o the star., The photographic plate archives from Harvard College Observatory and from Sonneberg Observatory are used to construct the long-term $B$ and $V$ light curves of the star.
 The results sugeest that CM Cop do uot slow fast rises in brightuess typical of ENor variables aud the light curves secin to be dominated by dips superposec ou the quiescence state., The results suggest that GM Cep do not show fast rises in brightness typical of EXor variables and the light curves seem to be dominated by dips superposed on the quiescence state.
 Evidences for periodicity of observed dips in brigltucss were not found., Evidences for periodicity of observed dips in brightness were not found.
 Recent BVRI CCD photometric observatious of CA Cop are reported iu the preseut paper., Recent $BVRI$ CCD photometric observations of GM Cep are reported in the present paper.
 We try to collec regular observatious of the star in order to clarify the nature of this object., We try to collect regular observations of the star in order to clarify the nature of this object.
 The multicolor observations eive us the opportunity to determine the mechanism of the brightness variations., The multicolor observations give us the opportunity to determine the mechanism of the brightness variations.
 Our photometric CCD data were obtained in two observatories with three telescopes: the 2-11. Ritchey-Chircttien-Coudé aud the 50/70-c Schmidt telescopes of the National Astronomical Observatory Rozheu (Dulgaria) aud the 1.3-11 Ritchey-Créttien telescope of the Skinakas of the Iustitute of Astronomy University of Crete (Cireecce).," Our photometric CCD data were obtained in two observatories with three telescopes: the 2-m Ritchey-Chr\'{e}ttien-Coud\'{e} and the 50/70-cm Schmidt telescopes of the National Astronomical Observatory Rozhen (Bulgaria) and the 1.3-m Ritchey-Créttien telescope of the Skinakas of the Institute of Astronomy, University of Crete (Greece)."
 Observations were made with tree types of CCD camera - Vers Array 1300B at the ιν RCC telescope.," Observations were made with tree types of CCD camera - Vers Array 1300B at the 2-m RCC telescope,"
"K g-band images, or located within 20 pixels of the edge or other blanked pixels within individual K;-band frames.","$K^{}_{S}$ -band images, or located within 20 pixels of the edge or other blanked pixels within individual $K^{}_{S}$ -band frames."
" Following ? and ?,, the probability that a NIR counterpart with magnitude m is the true NIR counterpart of the radio source is given by the likelihood ratio with m the K-band magnitude of the NIR candidate, 0(«m) the a priori probability that the radio source has a NIR counterpart with a magnitude brighter than m, and p(«m) the surface number density of NIR sources with a magnitude smaller than m."," Following \citet{sutherland1992} and \citet{tasse2008}, the probability that a NIR counterpart with magnitude $m$ is the true NIR counterpart of the radio source is given by the likelihood ratio with $m$ the $K^{}_{S}$ -band magnitude of the NIR candidate, $\theta(<m)$ the a priori probability that the radio source has a NIR counterpart with a magnitude brighter than $m$, and $\rho(<m)$ the surface number density of NIR sources with a magnitude smaller than $m$."
" o2 and c2 are the quadratic sums of the uncertainties in the radio and NIR positions in right ascension and declination, respectively."," $\sigma^{2}_{\delta}$ and $\sigma^{2}_{\alpha}$ are the quadratic sums of the uncertainties in the radio and NIR positions in right ascension and declination, respectively."
" For the radio source positions and uncertainties we have taken the values from the 1.4 GHz WSRT catalogue, as the astrometric precision is better than the 153 MHz GMRT data."," For the radio source positions and uncertainties we have taken the values from the 1.4 GHz WSRT catalogue, as the astrometric precision is better than the 153 MHz GMRT data."
" For the FLAMEX survey we have adopted 0.3"" for the position uncertainty (this is half the CCD pixel size, as source positions are truncated to pixel positions)."," For the FLAMEX survey we have adopted $0.3\arcsec$ for the position uncertainty (this is half the CCD pixel size, as source positions are truncated to pixel positions)."
" 'The uncertainty-normalized angular distance is given by the parameter r=J/(A4/o,)?+(A;/o5)?, with A the positional difference in either o or ó between the possible NIR counterpart and radio source."," The uncertainty-normalized angular distance is given by the parameter $r = \sqrt{(\Delta^{}_{\alpha} / \sigma^{}_{\alpha})^2 + (\Delta^{}_{\delta} / \sigma^{}_{\delta})^{2}_{}}$, with $\Delta$ the positional difference in either $\alpha$ or $\delta$ between the possible NIR counterpart and radio source."
 The probability P(i) that candidate i is the true NIR counterpart is with j running over all possible NIR counterparts., The probability $P(i)$ that candidate $i$ is the true NIR counterpart is with $j$ running over all possible NIR counterparts.
" 0(mq,) is the fraction of radio sources having a NIR counterpart at the magnitude limit of the NIR survey.", $\theta(m^{}_\mathrm{lim})$ is the fraction of radio sources having a NIR counterpart at the magnitude limit of the NIR survey.
" We found @(miim) to be 0.62, where mj=19.2."," We found $\theta(m_\mathrm{lim})$ to be 0.62, where $m_\mathrm{lim} = 19.2$."
 The values for p(«m) (surface number density of NIR sources) were estimated from the data itself using bins of 0.1 magnitude across the full survey area., The values for $\rho(<m)$ (surface number density of NIR sources) were estimated from the data itself using bins of $0.1$ magnitude across the full survey area.
" The a-priori identification fraction 0(«m) was also estimated from the data itself, using the technique described by ?.."," The a-priori identification fraction $\theta(<m)$ was also estimated from the data itself, using the technique described by \citet{ciliegi2003}."
 This involves counting the surface number density of NIR sources around the radio sources as function of magnitude., This involves counting the surface number density of NIR sources around the radio sources as function of magnitude.
" This distribution is compared to a distribution of background objects, the overdensity of NIR sources around radio sources gives an estimate of 0(«m)."," This distribution is compared to a distribution of background objects, the overdensity of NIR sources around radio sources gives an estimate of $\theta(<m)$."
" We selected 368 radio sources that are present within both the 1.4 GHz and 153 MHz catalogs (matched within 6"") and have a simple morphology (fitted with a single gaussian)", We selected 368 radio sources that are present within both the 1.4 GHz and 153 MHz catalogs (matched within $6\arcsec$ ) and have a simple morphology (fitted with a single gaussian).
" We have computed the likelihood ratio of all NIR counterparts located within 20"" from the radio position.", We have computed the likelihood ratio of all NIR counterparts located within $20\arcsec$ from the radio position.
" We have defined a radio source to have a NIR counterpart if 5»;LR;(r,m)>0.75."," We have defined a radio source to have a NIR counterpart if $\sum_{j} \mathrm{LR}^{}_{j}(r,m) > 0.75$."
 The results are shown in Figure 6.., The results are shown in Figure \ref{fig:bootes_identification}.
" The identification fraction is ~70 percent for o1309>—0.7, while for o1399?<—0.7 the identification fraction drops to about 30 percent."," The identification fraction is $\sim 70$ percent for $\alpha^{1400}_{153} > -0.7$, while for $\alpha^{1400}_{153} < -0.7$ the identification fraction drops to about 30 percent."
" This result reproduces the expected (and previously observed) correlation between spectral index and K-band identification fraction (e.g.,?).."," This result reproduces the expected (and previously observed) correlation between spectral index and $K$ -band identification fraction \citep[e.g.,][]{wieringa1991}."
" We have presented the results from a deep (c1.0 mJy central noise), high-resolution (26""x22"") radio survey at 153 MHz, covering the full NOAO field and beyond."," We have presented the results from a deep $\sim 1.0$ mJy central noise), high-resolution $26\arcsec \times 22\arcsec$ ) radio survey at 153 MHz, covering the full NOAO field and beyond."
 This 11.3 square degree survey is amongst the deepest surveys at this frequency to date., This 11.3 square degree survey is amongst the deepest surveys at this frequency to date.
" We produced a catalog of 598 sources detected at 5 times the local noise level, with source flux densities ranging from 3.9 Jy down to 5.1 mJy."," We produced a catalog of 598 sources detected at 5 times the local noise level, with source flux densities ranging from 3.9 Jy down to 5.1 mJy."
" We estimate the catalog to be ~70 percent complete at the 5.1 mJy flux limit, ~92 percent for >10 mJy sources and >99 percent for >20 mJy sources, and less than 1 percent contaminated by false detections."," We estimate the catalog to be $\sim 70$ percent complete at the $5.1$ mJy flux limit, $\sim 92$ percent for $> 10$ mJy sources and $> 99$ percent for $> 20$ mJy sources, and less than $1$ percent contaminated by false detections."
" The on-source dynamic range (with noise measured near the source) is limited to ~600, while off-source (noise measured in central part of the image, away from bright sources) this rises to >1000."," The on-source dynamic range (with noise measured near the source) is limited to $\sim 600$, while off-source (noise measured in central part of the image, away from bright sources) this rises to $>1000$."
 We expect that residual RFI and residual direction-dependent calibration errors prevents reaching the thermal noise level of 0.2-0.3 beam-!., We expect that residual RFI and residual direction-dependent calibration errors prevents reaching the thermal noise level of 0.2-0.3 .
. The 153 MHz catalog presented in Section 3.1 allows for a detailed study of source populations in a relatively unexplored flux range., The 153 MHz catalog presented in Section \ref{sec:bootes_bdsm} allows for a detailed study of source populations in a relatively unexplored flux range.
 We have analyzed the source counts and spectral index distributions for our survey., We have analyzed the source counts and spectral index distributions for our survey.
" From this analysis we draw the following conclusions: (i) The Euclidean-normalized differential source counts, determined over a flux range from 10 mJy to 1 Jy, are well approximated by a single power-law slope of 0.91."," From this analysis we draw the following conclusions: (i) The Euclidean-normalized differential source counts, determined over a flux range from 10 mJy to 1 Jy, are well approximated by a single power-law slope of 0.91."
 The inconsistency with model source counts by ? and ? seems due to uncertainty in the predicted contribution of galaxies., The inconsistency with model source counts by \citet{jackson2005} and \citet{wilman2008} seems due to uncertainty in the predicted contribution of galaxies.
 Our survey is not deep enough to detect the flattening at the lowest flux densities as seen at higher frequencies (Section 4.1))., Our survey is not deep enough to detect the flattening at the lowest flux densities as seen at higher frequencies (Section \ref{sec:bootes_dsc}) ).
caused by overestimating the sky level in their images and/or the cllective radii of their galaxies.,caused by overestimating the sky level in their images and/or the effective radii of their galaxies.
 We conclude that the magnitudes of low-redshift. radio galaxy hosts have been systematically overestimated by no more than about mmag as a result of non-stellar radiation [from the obscured quasar nuclei. and that conclusions drawn so far concerning the cosmic evolution of radio galaxies LLillv Longair 1984: Eales Rawlings 1996) are therefore reliable.," We conclude that the magnitudes of low-redshift radio galaxy hosts have been systematically overestimated by no more than about mag as a result of non-stellar radiation from the obscured quasar nuclei, and that conclusions drawn so far concerning the cosmic evolution of radio galaxies Lilly Longair 1984; Eales Rawlings 1996) are therefore reliable."
 We are. grateful to Jo AleAllister for useful. cliscussions regarding the accuracy of the 2-D fitting results. and to the referee. Jim Dunlop. for suggesting improvements to the manuscript.," We are grateful to Jo McAllister for useful discussions regarding the accuracy of the 2-D fitting results, and to the referee, Jim Dunlop, for suggesting improvements to the manuscript."
 “Phe United Ixingdom Infrared Telescope 1s operated. by the Joint Astronomy Centre on behalf of the U. Ix. Particle Physies anc Astronomy Research Council. and we thank the UINIICE stall for their help.," The United Kingdom Infrared Telescope is operated by the Joint Astronomy Centre on behalf of the U. K. Particle Physics and Astronomy Research Council, and we thank the UKIRT staff for their help."
 We also thank Mike Coacl lor providing a coded. version of the simulated annealing minimization routine., We also thank Mike Goad for providing a coded version of the simulated annealing minimization routine.
" Phis work has made use of the NASA/IPAC Extragalactic Database (NED). operated bv the Jet. Propulsion Laboratory. California Institute. of ""Technology. under a contract with the National Acronautics and Space Administration."," This work has made use of the NASA/IPAC Extragalactic Database (NED), operated by the Jet Propulsion Laboratory, California Institute of Technology, under a contract with the National Aeronautics and Space Administration."
"caltecli.eduον, The past two decades have seen vast miprovenieuts in the sensitivity and detector counts of broadbaud (sub)nilhuuneter Πάσης cameras. from sinele pixel receivers τω backeround-limuited kilopixel arrays (6.8. Duncanetal.(1990):Tolland(1999):Ruuvauct (20063))."," The past two decades have seen vast improvements in the sensitivity and detector counts of broadband (sub)millimeter imaging cameras, from single pixel receivers to background-limited kilopixel arrays (e.g., \citet{duncan90, holland99, runyan03, dowell03, 
haig04, carlstrom09, swetz11, holland06}) )."
" Maux of these mstruineuts have relied ou the auets. along with a haudful of secondary. calibrators. o obtain 10flux calibrations (οσοι, IBlHetal. (20093.. Rudy.<etal. (1987)... ασandOrton.(1993.jiearafter (1921). aud Saucell (1991)))."," Many of these instruments have relied on the planets, along with a handful of secondary calibrators, to obtain $\lesssim 10$flux calibrations (e.g., \citet{hill09}, \citet{rudy87}, \citet{griffin93}, and \citet{sandell94}) )."
. Alternatively. some of the Luge-scale survey instruments have used cosluic nücrowave backeround (CMD) measurements o calibrate their data to even better precision (6.9.. Reichardtetal. (20093)).," Alternatively, some of the large-scale survey instruments have used cosmic microwave background (CMB) measurements to calibrate their data to even better precision (e.g., \citet{reichardt09}) )."
 ILowewer. calibrating from ιο CAB is only. practical for iustrunents that image =100 dee?. and therefore cannot be done in many cases.," However, calibrating from the CMB is only practical for instruments that image $\gtrsim 100$ $^2$, and therefore cannot be done in many cases."
 When plauets and secoudary calibrators are used. ιο accuracv of the flux. calibration is often lmnüted wouncertaintics in the brightuess of these calibrators rather than bw measurement uncertaiuties (6.9.. Saversetal. (2009))).," When planets and secondary calibrators are used, the accuracy of the flux calibration is often limited by uncertainties in the brightness of these calibrators rather than by measurement uncertainties (e.g., \citet{sayers09}) )."
 Consequently. accurate measurements of i6 absolute brightuess of the planets and secoudirv calibrators. along with a detailed uuderstaudiug of any cluporal variability iu these brightuesses. is critical to naxiuizing the scientific potential of eiuxreut and future (sub)nillinmeter cameras.," Consequently, accurate measurements of the absolute brightness of the planets and secondary calibrators, along with a detailed understanding of any temporal variability in these brightnesses, is critical to maximizing the scientific potential of current and future (sub)millimeter cameras."
 Until receutlv. svstematic uncertaiufies in the xiehtness of Urauus aud Neptune— lave been 25.—10% owing to uncertainties in the Martian brigltuess used o determine the absolute brightuess of cach of these dauets (Wrieht(1976). Rudyetal. (1987).. OsG)... 693).," Until recently, systematic uncertainties in the brightness of Uranus and Neptune have been $\simeq 5-10$, owing to uncertainties in the Martian brightness used to determine the absolute brightness of each of these planets \citet{wright76}, \citet{rudy87}, , G93)."
 However. recently published WALAP results lave significantly improved the constraints on the brightuess of Maus (~1% }}) and Uranus (2 3%)) by calibrating them relative to the CXMB (Pageetal.2003:Tall2009:Weilandctal.," However, recently published WMAP results have significantly improved the constraints on the brightness of Mars $\simeq 1$ ) and Uranus $\simeq 3$ ) by calibrating them relative to the CMB \citep{page03, hill09, weiland11}."
 POLL)... WALAP has also published CAIB-calibrated brightness measurements of Neptune accurate to (Weilandetal.2011)., WMAP has also published CMB-calibrated brightness measurements of Neptune accurate to \citep{weiland11}.
. Additionally. receutlv published results have shown evideuce for a 21054 decrease in the brightuess of Uranus at 8.6 and 90 CGIIz over the z20 vear period from the mid 1980s to the mid 2000s (σαιandHofstadter(2006.hearafterI106)./ aud Isrameretal.(2008.lear-after Ix08))). which have beeu attributed to chanecs iu the relative viewing augle of the planet from earth due to its large obliquity of 82 deg.," Additionally, recently published results have shown evidence for a $\simeq 10$ decrease in the brightness of Uranus at 8.6 and 90 GHz over the $\simeq 20$ year period from the mid 1980s to the mid 2000s \citet{klein06} and \citet{kramer08}) ), which have been attributed to changes in the relative viewing angle of the planet from earth due to its large obliquity of 82 deg."
 Note that. over the period of observations in Kl06 and Ίντο the SEP latitude of Uranus raneed from a iin of 82 deg iu 1985 to a maxi of 7 dee in 2005.," Note that, over the period of observations in Kl06 and Kr08 the SEP latitude of Uranus ranged from a minimum of $-82$ deg in 1985 to a maximum of $-7$ deg in 2005."
 These data indicate that the brightness of Uranus increases as the magnitude of the sub-carth poiut (SEP) latitude increases(.c.. the south pole of Uranus is brighter than the equatorial regions).," These data indicate that the brightness of Uranus increases as the magnitude of the sub-earth point (SEP) latitude increases (i.e., the south pole of Uranus is brighter than the equatorial regions)."
 Further evidence for this scenario comes from resolved cniwave iages of Uranus. which iudicate that the poles are brighter than the equatorial regions (Ortonctal. 2007).," Further evidence for this scenario comes from resolved cm-wave images of Uranus, which indicate that the poles are brighter than the equatorial regions \citep{orton07}."
 Bolocain is a larec-format bolometric camera with au S arcnun field of view capable of observing at either 113 or 268 GIIz., Bolocam is a large-format bolometric camera with an 8 arcmin field of view capable of observing at either 143 or 268 GHz.
 Dolocin began collecting scicutific data in early 2003 at the CaltechSubiillimeter Observatory (CSO). and was commissioned as a facility istrmucut later that vear (aigetal. 2001)..," Bolocam began collecting scientific data in early 2003 at the CaltechSubmillimeter Observatory (CSO), and was commissioned as a facility instrument later that year \citep{haig04}. ."
 In this manuscript we focus exclusively on 113 GIIz Bolocam data. which was," In this manuscript we focus exclusively on 143 GHz Bolocam data, which was"
Essentially. a skewness means that if vou walk (initially) north (for example) for 1000 tiles. or walk south for 1000 miles. vou will cover different amounts of map coordinate.,"Essentially, a skewness means that if you walk (initially) north (for example) for 1000 miles, or walk south for 1000 miles, you will cover different amounts of map coordinate."
 The skewness along a geodesic can be defined similarly to the flexion: As with the flexion. we may expand these out explicitly: A map with no skewness will have: Unlike the flexion. we know of no projections with zero skewuess everywhere.," The skewness along a geodesic can be defined similarly to the flexion: As with the flexion, we may expand these out explicitly: A map with no skewness will have: Unlike the flexion, we know of no projections with zero skewness everywhere."
 The Guomouic Projection is particularly interestiug., The Gnomonic Projection is particularly interesting.
 It bas a coordinate transformation: which is directly invertible to vield: We thus have the map metric:, It has a coordinate transformation: which is directly invertible to yield: The coordinate transformation is thus: We thus have the map metric:
In order to explore the possibility of central heating we calculated color temperatures based on the ratio of our 10.3 yan and 18.2 jm images.,In order to explore the possibility of central heating we calculated color temperatures based on the ratio of our 10.8 $\micron$ and 18.2 $\micron$ images.
 Temperature and emission optical depth maps from simple radiative transfer analysis provide a good first-order estimate of the sources of grain heating as well as the relative density of warm grains (Tresch-Fienberg et al., Temperature and emission optical depth maps from simple radiative transfer analysis provide a good first-order estimate of the sources of grain heating as well as the relative density of warm grains (Tresch-Fienberg et al.
 1987)., 1987).
 Figure 4 shows our temperature and emission optical depth maps., Figure 4 shows our temperature and emission optical depth maps.
 We calculate color temperatures rangine from 18545 Ix in the core to ~165415Ix. within the NLR (x~100 pc). consistent with the warm dust component (170 IX) as measured by ISO (RE9G).," We calculate color temperatures ranging from $\thicksim
185\pm 5$ K in the core to $\thicksim 165\pm 15$K within the NLR $%
r\thicksim 10 0 pc), consistent with the warm dust component (170 K) as measured by ISO (RE96)."
 The enission optical depth shows the density of these dust grains is enhanced along the direction ol the NLR., The emission optical depth shows the density of these dust grains is enhanced along the direction of the NLR.
 Assuming a simple uniform dust distribution. a first-order determination of the size of the region that could be heated by a central source can be made.," Assuming a simple uniform dust distribution, a first-order determination of the size of the region that could be heated by a central source can be made."
 Given that dust erains primarily absorb UV-optical radiation and re-emüit in the infrared. (he equilibrium temperature of dust in a strong UV [field is eiven by (Selleren et al.," Given that dust grains primarily absorb UV-optical radiation and re-emit in the infrared, the equilibrium temperature of dust in a strong UV field is given by (Sellgren et al."
 1983) In the above equation. T is the dust temperature. Lp) is the UV. luminosity of the central source. /2 is (he radius from (he source in parsecs. σ is (he Stefan-Doltzman constant. and ιν is the ratio of the Planck averaged UV absorption coefficient to the infrared enussion coefficient.," 1983) In the above equation, T is the dust temperature, $L_{UV}$ is the UV luminosity of the central source, $R$ is the radius from the source in parsecs, $\sigma $ is the Stefan-Boltzman constant, and $Q_{UV}/Q_{IR}$ is the ratio of the Planck averaged UV absorption coefficient to the infrared emission coefficient."
" Values of Ory/Qy, are dependent on (he dust grain size ancl composition and are given by Draine Lee (1984). Laor Draine (1993). and Weingartner Draine (2001) for graphite and “smoothecl astronomical” (SA) silicate."," Values of $Q_{UV}/Q_{IR}$ are dependent on the dust grain size and composition and are given by Draine Lee (1984), Laor Draine (1993), and Weingartner Draine (2001) for graphite and “smoothed astronomical” (SA) silicate."
 luminosity of NGC 4151 is ~LOM TheobservedUV-opticalL. (Penston et al.," The observed UV-optical luminosity of NGC 4151 is $\thicksim 10^{10}\ $ $_{%
\sun (Penston et al."
 1990)., 1990).
 Given (his huninositw. in order {ο heat dust to the observed temperature of ~165415 Ix ab a distance of e 100 pe. the inner sshould be composed of 0.004 pam graphite grains or 0.001 yan SA silicate grains.," Given this luminosity, in order to heat dust to the observed temperature of $\thicksim 165\pm 15$ K at a distance of $%
\thicksim 10 0 pc, the inner should be composed of $0.004$ $\micron$ graphite grains or $0.001$ $\micron$ SA silicate grains."
 These grain sizes fall near or below the limit of classical interstellar dust grains which range from 0.003 yam - 1 jan (Draine Lee 1934)., These grain sizes fall near or below the limit of classical interstellar dust grains which range from 0.003 $\micron$ - 1 $\micron$ (Draine Lee 1984).
 They. also are much smaller than the estimated grains sizes used for the centrally heated NLR. models of NGC! 1068 (~0.05 pam: Cameron et., They also are much smaller than the estimated grains sizes used for the centrally heated NLR models of NGC 1068 $\sim 0.05$ $\micron$; Cameron et.
 al., al.
 1993) and Cve A (~0.1 pam: Radomski et., 1993) and Cyg A $\sim 0.1$ $\micron$; Radomski et.
 al., al.
 2002)., 2002).
 Llowever Penston et al. (, However Penston et al. (
1990) proposed that the continuum emission in NGC 4151 is inherently anisotropic and that (he ionizing Iuminosity as seen within (he extended NLR (1-2NS kpc) may be on order of 13 (mes greater (han that observed from Earth.,1990) proposed that the continuum emission in NGC 4151 is inherently anisotropic and that the ionizing luminosity as seen within the extended NLR (1-2 kpc) may be on order of 13 times greater than that observed from Earth.
 Subsequent models bv Schulz Komossa (1993). Yoshida Obtani (1993). ancl Robinson et al. (," Subsequent models by Schulz Komossa (1993), Yoshida Ohtani (1993), and Robinson et al. ("
1994) also derive values for the anisotropy as high as 3-10.,1994) also derive values for the anisotropy as high as 3-10.
 Thus a better estimate of the Iuminosity in the NLR max be ~10! L..," Thus a better estimate of the luminosity in the NLR may be $%
\sim 10^{11} $_{\sun}$ ."
 This higher luminosity would increase the calculated grain sizes by an order of magnitude (0.04 jan graphite or 0.01 jan SA silicate). resulting in sizes more," This higher luminosity would increase the calculated grain sizes by an order of magnitude $0.04$ $\micron$ graphite or $0.01$ $\micron$ SA silicate), resulting in sizes more"
simulations provide the uncertainties in our estimate.,simulations provide the uncertainties in our estimate.
 The covariance matrix of Pj is ∇∇∐⊖↕⊖⋖⋗↕∨↕∁≺∼↕⊖∐⊖⇇⊖≋⇇∐⊖↕∖∕↕⊖∐⇇⊖−∁∂∐⊖∂∇⊖⊞≣↕∐≣∙↱∏↕⊖⊖∐⊙↕ , The covariance matrix of $\tilde{P}_b$ iswhere $\langle\cdot\rangle_{\text{MC}}$ denotes the Monte-Carlo averaging.
"bar in each P, is and the bin-bin correlation matrix is given by its standard definition The above sections describe the main features of the algorithm that provides an unbiased estimate of the power spectrum of datamap.", The error bar in each $\tilde{P}_b$ is and the bin-bin correlation matrix is given by its standard definition The above sections describe the main features of the algorithm that provides an unbiased estimate of the power spectrum of data.
. This power spectrum may however differ from the power spectrum., This power spectrum may however differ from the power spectrum.
" The map making process may indeed alter the statistical properties of the signal, together with data filtering and the convolution by the instrumental beam."," The map making process may indeed alter the statistical properties of the signal, together with data filtering and the convolution by the instrumental beam."
" For instance, the pixelization caused by the map making is equivalent to a convolution in real space by a square kernel and therefore translates into a multiplication in Fourier space by a factor sin,."," For instance, the pixelization caused by the map making is equivalent to a convolution in real space by a square kernel and therefore translates into a multiplication in Fourier space by a factor $\sin_c$."
" In the case of CMB anisotropy for which power spectrum estimation has been most extensively studied, to a good approximation, the transfer function F; of the map making and the data processing together reduces to a function of k that multiplies the data power spectrum Pi."," In the case of CMB anisotropy for which power spectrum estimation has been most extensively studied, to a good approximation, the transfer function $F_k$ of the map making and the data processing together reduces to a function of $k$ that multiplies the data power spectrum $P_k$."
" The determination of Fy, is performed by a set of Monte Carlo simulations of data processing and map making.", The determination of $F_k$ is performed by a set of Monte Carlo simulations of data processing and map making.
 The beam smearing effect is also described by a multiplicative function B;., The beam smearing effect is also described by a multiplicative function $B_k$.
" In the present framework, it is possible to be even more precise and to account for the exact beam shape and orientation because the beam can be completely described by its Fourier coefficients By» rather than its approximated annular average Dj."," In the present framework, it is possible to be even more precise and to account for the exact beam shape and orientation because the beam can be completely described by its Fourier coefficients $B_{mn}$ rather than its approximated annular average $B_k$."
 This may be of particular relevance for small fields over which the scanning directions are approximately constant and increase the effect of beam asymmetry., This may be of particular relevance for small fields over which the scanning directions are approximately constant and increase the effect of beam asymmetry.
" The map making together with the filtering transfer function is also likely to be more accurately represented by a function of both Fourier indices F5,4, such that Eq. (10))"," The map making together with the filtering transfer function is also likely to be more accurately represented by a function of both Fourier indices $F_{mn}$, such that Eq. \ref{eq:ps2cl}) )"
" is given by These new contributions can be incorporated in the definition of the convolution kernel Kw"", such that no additional modification of the algorithm is needed from Eq. (10))"," is given by These new contributions can be incorporated in the definition of the convolution kernel $K_{m,m'}^{n,n'}$ such that no additional modification of the algorithm is needed from Eq. \ref{eq:ps2cl}) )"
 onward., onward.
 An outline of the algorithm is presented in Fig. 6.., An outline of the algorithm is presented in Fig. \ref{fig:chart}. .
" To perform the complete process of power spectrum and statistical uncertainty estimation, one needs:"," To perform the complete process of power spectrum and statistical uncertainty estimation, one needs:"
asymptotically locally Euclidean (ALE) complex space which is algebraically given by which is singular at +=y:0.,asymptotically locally Euclidean (ALE) complex space which is algebraically given by which is singular at $x=y=z=0$.
" In two-dimensional \=2 linear sigma model with only (1) gauge symmetry, the resolution of this singularity is related to the D-term described by the following bosonic potential V(64.02.63) where H is the /(1) Fayet-Iliopoulos (FI) [21].."," In two-dimensional $N = 2$ linear sigma model with only $U(1)$ gauge symmetry, the resolution of this singularity is related to the D-term described by the following bosonic potential $V(\phi_1, \phi_2,
 \phi_3)$ where $R$ is the $U(1)$ Fayet-Iliopoulos (FI) \cite{W}."
" Geometrically, this corresponds to replace the singular point.r =y:0 by the 2-sphere 5? defined by τι©=0, which is the only non-trivial 2-cycle on which we can wrap D2-branes."," Geometrically, this corresponds to replace the singular point $x = y = z = 0 $ by the 2-sphere $S^2$ defined by $
V=\phi_2=0$, which is the only non-trivial 2-cycle on which we can wrap D2-branes."
 In order to geometrically engineer the SU(2) gauge symmetry only the compact piece containing the 5? is necessary [22].., In order to geometrically engineer the SU(2) gauge symmetry only the compact piece containing the $S^2$ is necessary \cite{KKV}.
" Now, the system consists of Type ΗΠΑ D2-branes wrapping around 57."," Now, the system consists of Type IIA D2-branes wrapping around $S^2$."
" This gives a pair of massive vectors 117, one for each of the two possible ways of the wrapping."," This gives a pair of massive vectors $W^{\pm}$, one for each of the two possible ways of the wrapping."
 The masses of these particles are proportional to the volume of the 2-sphere., The masses of these particles are proportional to the volume of the 2-sphere.
 They are charged under the {{1} gauge field obtained by decomposing the type IA three-form in terms of the harmonic 2-form on the 2-sphere and the 1-form gauge field in six dimensional space-time., They are charged under the $U(1)$ gauge field obtained by decomposing the type IIA three-form in terms of the harmonic 2-form on the 2-sphere and the 1-form gauge field in six dimensional space-time.
" In the limit where the shrinks, the 1!- particles become massless and, together with the one form gauge field, generate the 5U(2) adjoint representation."," In the limit where the 2-sphere shrinks, the $W^{\pm}$ particles become massless and, together with the one form gauge field, generate the $SU(2)$ adjoint representation."
 This will be identified with the gauge symmetry of our Yang monopole., This will be identified with the gauge symmetry of our Yang monopole.
" We have obtained the electrically charged sector, associated to D2-branes wrapping two vanishing cycles in the K3 surface."," We have obtained the electrically charged sector, associated to D2-branes wrapping two vanishing cycles in the K3 surface."
" Lifting consistently to 11 dimensions, the M2-brane is encountered."," Lifting consistently to 11 dimensions, the M2-brane is encountered."
" The magnetic Yang monopoles can be identified with D4-branes, totally wrapped on the K3 surface."," The magnetic Yang monopoles can be identified with D4-branes, totally wrapped on the K3 surface."
" As consequence, they generate the magnetic objects in the six dimensional space-time."," As consequence, they generate the magnetic objects in the six dimensional space-time."
 This is expected from the fact that the D4-brane is the only magnetic object in Type ITA superstring theory which can be obtained from the M5-brane and gives a particle after wrapping the K3 surface., This is expected from the fact that the D4-brane is the only magnetic object in Type IIA superstring theory which can be obtained from the M5-brane and gives a point-like particle after wrapping the K3 surface.
" Owing to the spherical symmetry of the six dimensional configuration and on the fact that, as seen above, the gauge group origin is linked to the singular limit of the geometry, we strongly believe that all Yang monopole properties should be derived from the K3 surface data."," Owing to the spherical symmetry of the six dimensional configuration and on the fact that, as seen above, the gauge group origin is linked to the singular limit of the geometry, we strongly believe that all Yang monopole properties should be derived from the K3 surface data."
" Schematically, the ten dimensional spacetime where Type HA lives is occupied as follows."," Schematically, the ten dimensional spacetime where Type IIA lives is occupied as follows."
" If we consider that the K3 surface extents along dimensions 6789, and the vanishing 2-cycle of K3 at 67 positions, then the D2-brane is at 067 and the D4-brane at 06769."," If we consider that the K3 surface extents along dimensions 6789, and the vanishing 2-cycle of K3 at 67 positions, then the D2-brane is at 067 and the D4-brane at 06789."
"that in the outer layers (r20.9Ro) of the Sun the kinetic energy is greatest at epochs of greatest activity, both in the low-latitude and in the high-latitude regions.","that in the outer layers $(r \ga 0.9 {\rm R}_{\odot})$ of the Sun the kinetic energy is greatest at epochs of greatest activity, both in the low-latitude and in the high-latitude regions."
" At intermediate depths (0.7Sr/Ro50.9) the high-latitude correlation with activity persists, but at low latitudes the correlation is reversed."," At intermediate depths $(0.7 \la
r/{\rm R}_{\odot} \la 0.9)$ the high-latitude correlation with activity persists, but at low latitudes the correlation is reversed."
" This trend is seen more clearly in Figure 3, in which the correlation coefficients are plotted as a function of radius, both for entire spherical shells and for the high- and low-latitude regions."," This trend is seen more clearly in Figure 3, in which the correlation coefficients are plotted as a function of radius, both for entire spherical shells and for the high- and low-latitude regions."
" The amplitude of the variation is the greater in the low-latitude region, so the behaviour there is an indicator of the behaviour of the kinetic energy, and the angular momentum, over the entire spherical shell."," The amplitude of the variation is the greater in the low-latitude region, so the behaviour there is an indicator of the behaviour of the kinetic energy, and the angular momentum, over the entire spherical shell."
 At the greatest depths it is more difficult to see the trends in Figure 2 because the uncertainties in the inversions are comparable with the trends themselves., At the greatest depths it is more difficult to see the trends in Figure 2 because the uncertainties in the inversions are comparable with the trends themselves.
" However, the covariance of the kinetic energy with the activity formally changes sign at r~0.7Ro, at both low and high latitudes, beneath which the low-latitude kinetic energy variation is again positively correlated with activity and the high-latitude variation is negatively correlated."," However, the covariance of the kinetic energy with the activity formally changes sign at $r \simeq 0.7{\rm R}_{\odot}$, at both low and high latitudes, beneath which the low-latitude kinetic energy variation is again positively correlated with activity and the high-latitude variation is negatively correlated."
" Again, this behaviour is visible more clearly in Figure 3."," Again, this behaviour is visible more clearly in Figure 3."
 These results are consistent with those of Komm et al. (, These results are consistent with those of Komm et al. (
"2003), who considered the angular-momentum variations over entire spherical shells.","2003), who considered the angular-momentum variations over entire spherical shells."
" The relative variations in angular momentum and rotational kinetic energy, and also, of course, the angular velocity, are of order0."," The relative variations in angular momentum and rotational kinetic energy, and also, of course, the angular velocity, are of order."
"1%.. Given that the total angular momentum and rotational kinetic energy of the convection zone are J,«2.8x10gcm?s! and T,~4.0x104! erg, this implies that the amplitudes of their variations are about 1x10“gcm?s! and 3x1035 erg respectively."," Given that the total angular momentum and rotational kinetic energy of the convection zone are $J_{\rm c} \simeq 2.8 
\times 10^{47} {\rm g\, cm}^2 {\rm s}^{-1}$ and $T_{\rm c} \simeq 4.0 
\times 10^{41}$ erg, this implies that the amplitudes of their variations are about $1 \times 10^{44} 
{\rm~g~cm}^2{\rm s}^{-1}$ and $3 \times 10^{38}~$ erg respectively."
" The change in sign at r~0.9Ro of the correlation of angular momentum with activity at low latitudes could be associated with the upward migration of the zonal flow pattern. reported by Basu Antia (2003): because the migration timescale from intermediate depths is about half a cycle period, the flow pattern at intermediate depths is anticorrelated with that at the surface."," The change in sign at $r
\simeq 0.9{\rm R}_{\odot}$ of the correlation of angular momentum with activity at low latitudes could be associated with the upward migration of the zonal flow pattern reported by Basu Antia (2003): because the migration timescale from intermediate depths is about half a cycle period, the flow pattern at intermediate depths is anticorrelated with that at the surface."
" At high latitudes, the angular-momentum variation is positively correlated with surface activity throughout the convection zone."," At high latitudes, the angular-momentum variation is positively correlated with surface activity throughout the convection zone."
" Beneath the convection zone the correlation appears to be weaker, possibly as a result of the greater uncertainties in the inferred angular velocity, and may not be significant."," Beneath the convection zone the correlation appears to be weaker, possibly as a result of the greater uncertainties in the inferred angular velocity, and may not be significant."
" To calculate the amplitude of the solar-cycle variation in kinetic energy, we fit a sinusoid with a period of 11 years to the Observed kinetic energy."," To calculate the amplitude of the solar-cycle variation in kinetic energy, we fit a sinusoid with a period of 11 years to the observed kinetic energy."
 Since the power spectra at most depths show a peak at frequencies corresponding to approximately this, Since the power spectra at most depths show a peak at frequencies corresponding to approximately this
of the mean collision time can be also obtained from the relation 7=L/noyc.,of the mean collision time can be also obtained from the relation $\tau =1/n\sigma _{0}v$.
" By assuming that (he cross section of quark-quark interactions is around σῃ2 1. mb and the quarks are travelling with the speed of light. we obtain 7=41.6 [m (=12.48x107?«) [or nj,=0.24 ? and 7=25 [m (=7.5xLO7s) for nj=0.4 bn*."," By assuming that the cross section of quark-quark interactions is around $%
\sigma _{0}\approx 1 mb and the quarks are travelling with the speed of light, we obtain $\tau =41.6$ fm $\left( =12.48\times 10^{-23}%
\textrm{ s}\right) for $n_{b}=0.24$ $^{-3}$ and $\tau =25$ fm $\left(
=7.5\times 10^{-23}\textrm{ s}\right) $ for $n_{b}=0.4$ $^{-3}$."
 The knowledgee of 7 allows the ealeulation of the spectrum of the photons emitted via the bremsstrahhimee mechanism from the quark star surface., The knowledge of $\tau $ allows the calculation of the spectrum of the photons emitted via the bremsstrahlung mechanism from the quark star surface.
 The applicability of the formulas for the spectral distribution aud intensity of the bremsstrahilung radiation we have derived in the previous Section is limited bv (he quantum condition that w be small as compared with the total kinetic οποίον of the quark. which is of the order of Ejzz(2ny)1/5w <<|Tp)QL/3," The applicability of the formulas for the spectral distribution and intensity of the bremsstrahlung radiation we have derived in the previous Section is limited by the quantum condition that $\omega $ be small as compared with the total kinetic energy of the quark, which is of the order of $E_{F}\approx \left( \pi ^{2}n_{b}\right) ^{1/3}$: $\omega <<\left( \pi
^{2}n_{b}\right) ^{1/3}$."
 since al the stars surlace the kinetic energv of the quarks is of the order of 300 MeV. the radiation formulae are valid only [or frequencies w<<300 MeV. In order to find the maximum [frequency cya; of the radiation in (the ultrarelativistic case we assume (hat il is of the order yac1/7(1—(07). that is it is inversely proportional to the duration of the collision aud (o the square of the energy of (he radiating particle1975).," Since at the star's surface the kinetic energy of the quarks is of the order of $300$ MeV, the radiation formulae are valid only for frequencies $\omega <<300$ MeV. In order to find the maximum frequency $\omega _{\max }$ of the radiation in the ultrarelativistic case we assume that it is of the order $\omega _{\max }\sim 1/\tau \left(
1-v^{2}\right) $, that is it is inversely proportional to the duration of the collision and to the square of the energy of the radiating particle."
". IL we assume that quarks have energies near the Fermi energy. (hen (he corresponding velocities are the Fermi velocities. which are of the order of vj=0.92 for a,=0.4 and ep=0.85 for a,=0.9."," If we assume that quarks have energies near the Fermi energy, then the corresponding velocities are the Fermi velocities, which are of the order of $v_{F}=0.92$ for $\alpha _{s}=0.4$ and $v_{F}=0.85$ for $\alpha _{s}=0.9$."
" Hence the maximum allowable frequency up to which the formalismcan be applied is given by For a quark-quark interaction cross sections of the order of 1 mb and for a,= 0."," Hence the maximum allowable frequency up to which the formalismcan be applied is given by For a quark-quark interaction cross sections of the order of $1$ mb and for $%
\alpha _{s}=0."
"4 we obtain «ya~30 MeV. In the following we shall restrict. our study only to quark bremsstrahlung frequencies satislving the condition vw««4,3530 MeV. The frequency distribution of the radiation emitted by a single quark moving in the dense matter at the surface of the quark star is presented in Fig.", 4 we obtain $\omega _{\max }\sim 30$ MeV. In the following we shall restrict our study only to quark bremsstrahlung frequencies satisfying the condition $\omega <\omega _{\max }\approx 30$ MeV. The frequency distribution of the radiation emitted by a single quark moving in the dense matter at the surface of the quark star is presented in Fig.
 3., 3.
 The total intensity Z of the bremsstrallung radiation emitted by a single quark. obtained by integratingEq. (13)," The total intensity $I$ of the bremsstrahlung radiation emitted by a single quark, obtained by integratingEq. )"
) over all possible frequencies is represented. as a function of the collision time 7. in Fig.," over all possible frequencies is represented, as a function of the collision time $\tau $ , in Fig."
 4., 4.
cluster RZ2109 is 10-3 Gyr.,cluster RZ2109 is $10 \pm 3$ Gyr.
 If the black hole is found to be in the globular cluster in the present epoch. its progenitor must have formed with the rest of the stars in the cluster.," If the black hole is found to be in the globular cluster in the present epoch, its progenitor must have formed with the rest of the stars in the cluster."
 This is because a globular cluster in a galaxy cannot capture any star that was previously unbound to it (e.g..2)..," This is because a globular cluster in a galaxy cannot capture any star that was previously unbound to it \citep[e.g.,][]{mieske_baumgardt07}."
 The cluster gravitational potential is too small compared to the kinetic energy of objects orbiting a large elliptical galaxy such as NGC 4472., The cluster gravitational potential is too small compared to the kinetic energy of objects orbiting a large elliptical galaxy such as NGC 4472.
 A typical velocity dispersion inside globular clusters is 1O km s7!. while orbital velocities in elliptical galaxies are in excess of 200 km st.," A typical velocity dispersion inside globular clusters is $\sim 10$ km $^{-1}$, while orbital velocities in elliptical galaxies are in excess of 200 km $^{-1}$."
 The lifetime of a progenitor star that led to the formation of the black hole in a type II supernova explosion is negligibly short., The lifetime of a progenitor star that led to the formation of the black hole in a type II supernova explosion is negligibly short.
 Stellar evolution models for a wide range of progenitor masses indicate main-sequence lifetimes smaller than a few times 10 yr., Stellar evolution models for a wide range of progenitor masses indicate main-sequence lifetimes smaller than a few times $10^7$ yr.
" As a result. the age of a 10M.. black hole in a globular cluster is practically equal to the age of the cluster itself,"," As a result, the age of a $10\, \Msun$ black hole in a globular cluster is practically equal to the age of the cluster itself."
 Using equation (1)) for the age of 10 Gyr and mass of 10M... we obtain an upper bound on the asymptotic radius of curvature of the extra dimensions L=0.003 mm.," Using equation \ref{eq:tau}) ) for the age of 10 Gyr and mass of $10 \
\Msun$, we obtain an upper bound on the asymptotic radius of curvature of the extra dimensions $L \lesssim 0.003$ mm."
 This bound is lower by an order of magnitude than previous table-top and astrophysical constraints., This bound is lower by an order of magnitude than previous table-top and astrophysical constraints.
 Figure 1. and Table 1. summarize the constraints on the maximum size of the extra dimensions from the black hole in RZ2109. the black holes in ΝΤΕ J1118+480 and A0620—00. and the laboratory measurements.," Figure \ref{fig:l} and Table \ref{tab:obs} summarize the constraints on the maximum size of the extra dimensions from the black hole in RZ2109, the black holes in XTE J1118+480 and $-$ 00, and the laboratory measurements."
 The new limit provided by the black hole in RZ2109 is much smaller than that for XTE J1118+480 because of the crucial difference in the black hole ages., The new limit provided by the black hole in RZ2109 is much smaller than that for XTE J1118+480 because of the crucial difference in the black hole ages.
 The shaded region shows the expected uncertainty in this limit due to the uncertainties of the black hole mass (between 5 and 20 M .) and its age (between 7 and 13 Gyr)., The shaded region shows the expected uncertainty in this limit due to the uncertainties of the black hole mass (between 5 and 20 $\Msun$ ) and its age (between 7 and 13 Gyr).
 Even for the worst combination of the parameters. the upper bound increases to L=0.01 mm. which is still an improvement by a factor of ~5 from the current laboratory limit.," Even for the worst combination of the parameters, the upper bound increases to $L\lesssim 0.01$ mm, which is still an improvement by a factor of $\simeq 5$ from the current laboratory limit."
 In this we placed a strong upper bound on the asymptotic radius of curvature of the extra dimensions. by estimating the age and mass of the black hole recently discovered in an extragalactic globular cluster.," In this we placed a strong upper bound on the asymptotic radius of curvature of the extra dimensions, by estimating the age and mass of the black hole recently discovered in an extragalactic globular cluster."
 The only potential caveats to our arguments would appear if the black hole were less massive when it formed and then subsequently merged with other stars or has grown significantly by aceretion., The only potential caveats to our arguments would appear if the black hole were less massive when it formed and then subsequently merged with other stars or has grown significantly by accretion.
 While such scenarios are possible in à. globular cluster. we argue here that they are unlikely to change our quantitative bound on the size of the extra dimensions.," While such scenarios are possible in a globular cluster, we argue here that they are unlikely to change our quantitative bound on the size of the extra dimensions."
 The black hole mass can change with time because of evaporation. accretion. and collisions with other stars: The rates of accretion and collisions are likely to. be variable in time. but to obtain practical estimates we make the following simplifying assumptions.," The black hole mass can change with time because of evaporation, accretion, and collisions with other stars: The rates of accretion and collisions are likely to be variable in time, but to obtain practical estimates we make the following simplifying assumptions."
 We take Mj to be constant in time and independent of the BH mass., We take $\dot{M}_{\rm coll}$ to be constant in time and independent of the BH mass.
 We take the accretion rate to be a constant fraction fe<| of the Eddington rate. where e is the radiative efficiency of accretion and ts typically ez0.1.," We take the accretion rate to be a constant fraction $f_E \le 1$ of the Eddington rate, where $\epsilon$ is the radiative efficiency of accretion and is typically $\epsilon \approx 0.1$."
" Thus Mj,xM.", Thus $\dot{M}_{\rm acc} \propto M$.
 With these assumptions the total mass derivative is a monotonic function of the mass: that is /M/df increases 1f increases. and vice versa.," With these assumptions the total mass derivative is a monotonic function of the mass: that is $dM/dt$ increases if $M$ increases, and vice versa."
 So If at one time in the past the derivativeM was positive the mass will only keep increasing. the rate of evaporation will diminish. and the black hole will survive.," So if at one time in the past the derivative was positive the mass will only keep increasing, the rate of evaporation will diminish, and the black hole will survive."
 Conversely. if the derivative was negative both accretion and collisions will become less important with time. and the black hole will evaporate in a finite amount time similar to του.," Conversely, if the derivative was negative both accretion and collisions will become less important with time, and the black hole will evaporate in a finite amount time similar to $\tau_{\rm ev}(M)$."
 Consider first the case of continuous accretion at a fraction fy of the Eddington limit., Consider first the case of continuous accretion at a fraction $f_E$ of the Eddington limit.
 The increase of mass by accretion would overwhelm the evaporation of the black hole if its mass, The increase of mass by accretion would overwhelm the evaporation of the black hole if its mass
On this basis. we can draw some interesting conclusions about the overall emission properties of pulsar DI944|17 Our analyses found no basis for determining the pulsar's carouscl circulation time: however. some speculations are appropriate.,"On this basis, we can draw some interesting conclusions about the overall emission properties of pulsar B1944+17— Our analyses found no basis for determining the pulsar's carousel circulation time; however, some speculations are appropriate."
 First. identification ofa circulation time may be dillieult or impossible in a pulsar with intermittent emission.," First, identification of a circulation time may be difficult or impossible in a pulsar with intermittent emission."
 Second. rough estimations of the star’s circulation time fron ltuderman Sutherland. (1975) or by extrapolaring from DB0943|LO (DIOI) both suggest that it could. beie. 10-15. P.," Second, rough estimations of the star's circulation time from Ruderman Sutherland (1975) or by extrapolaring from B0943+10 (DR01) both suggest that it could be, 10-15 $P_1$."
" HE there were then 10 or more “beamlets” in the carousel. the primary. ""drift modulation frequency could be aliased into the second order."," If there were then 10 or more “beamlets” in the carousel, the primary “drift” modulation frequency could be aliased into the second order."
" Such aliasing together with discrete changes in beamlet number could give rise o ""modes"" such as are observed in D1944|17.", Such aliasing together with discrete changes in “beamlet” number could give rise to “modes” such as are observed in B1944+17.
" For instance. were the star's carousel to circulate in some 12 £, with a configuration of 13. 7beamlets. drift modulation similar to he A mode would be produced."," For instance, were the star's carousel to circulate in some 12 $P_1$ with a configuration of 13 “beamlets”, drift modulation similar to the A mode would be produced."
 Similarly. 14. “beamlets” would modulate the PSs with a £5 like that of the D mocle. and. 12 would. produce a nearly stationary modulation very ike that of mode ο. We found it both fascinating and challenging to study a star such as DIO44|17 which exhibits so many distinct ποολο," Similarly, 14 “beamlets” would modulate the PSs with a $P_3$ like that of the B mode, and 12 would produce a nearly stationary modulation very like that of mode C. We found it both fascinating and challenging to study a star such as B1944+17 which exhibits so many distinct phenomena."
 Clearly. its complex modes ancl nulls still have much to αι.teach us.," Clearly, its complex modes and nulls still have much to teach us."
" Dillerent polarization measurements may be able to better define its sightline &cometry.. and interferometry could. well reveal a ""pedestal of continuous emission."," Different polarization measurements may be able to better define its sightline geometry, and interferometry could well reveal a “pedestal” of continuous emission."
 Finally. even longer observations are needed to further explore the stars carouscl-beam configuration and fully assess the frequeney of partial nulls.," Finally, even longer observations are needed to further explore the star's carousel-beam configuration and fully assess the frequency of partial nulls."
 We are pleased to acknowledge Vishal Gajjar. Tim Llankins. Jocri van Leeuwen and Ceollrey Wright for their critical readings of the manuscript anc Jelfrey Herfindal for assistance with aspects of the analysis.," We are pleased to acknowledge Vishal Gajjar, Tim Hankins, Joeri van Leeuwen and Geoffrey Wright for their critical readings of the manuscript and Jeffrey Herfindal for assistance with aspects of the analysis."
 We also sincerely thank both Tim Llankins for providing us with the original interpulse discovery observations. and Joel Weisbere for the ionospheric Faraday: rotation corrections.," We also sincerely thank both Tim Hankins for providing us with the original interpulse discovery observations, and Joel Weisberg for the ionospheric Faraday rotation corrections."
 One of us (IAL) sincerely thanks the Barry M. Goldwater Scholarship and. LExecllenee in [Education program ancl the UVM College of Arts ancl Sciences. for the APPLE Summer Fellowship. which together permitted her to complete this work.," One of us (IMK) sincerely thanks the Barry M. Goldwater Scholarship and Excellence in Education program and the UVM College of Arts and Sciences for the APLE Summer Fellowship, which together permitted her to complete this work."
 The other CJMIBR) thanks the Anton Pannekock Astronomical Institute of the University of Amsterdam for their generous hospitality ancl the Netherlands National Science Foundation and ASTRON for her Visitor Cirants., The other (JMR) thanks the Anton Pannekoek Astronomical Institute of the University of Amsterdam for their generous hospitality and the Netherlands National Science Foundation and ASTRON for her Visitor Grants.
 Portions of this work were carried out with support. [rom US National Science Foundation. Cuants AST. 99-ST654 ancl 08-07691., Portions of this work were carried out with support from US National Science Foundation Grants AST 99-87654 and 08-07691.
 Arecibo Observatory is operated: by Cornell University under contract to the US NSE., Arecibo Observatory is operated by Cornell University under contract to the US NSF.
 This work used the NASA ADS system., This work used the NASA ADS system.
"The results of the observations are eiven in Table 1 where we list the PN nane. the intensities Do, of the CO and PCO J=2 1 lines aud the derived isotopic ratio.","The results of the observations are given in Table 1 where we list the PN name, the intensities $I_{21}$ of the $^{12}$ CO and $^{13}$ CO $J=2$ –1 lines and the derived isotopic ratio."
 The sample of the six PNe studied in Πο shows little anission in CO: with the exception of NCC 6720. the values of D4 represent upper limits to the intensity aud have been estimated from the line widths decced from the expansion velocities listed in the catalog cf Acker et al. (," The sample of the six PNe studied in $^3$ He shows little emission in CO: with the exception of NGC 6720, the values of $_{21}$ represent upper limits to the intensity and have been estimated from the line widths deduced from the expansion velocities listed in the catalog of Acker et al. ("
1992).,1992).
 Bachiller et al. (, Bachiller et al. (
"1989) detected au extcucdec nolecular suvelope iu NGC 6720. the Ring ποία,","1989) detected an extended molecular envelope in NGC 6720, the Ring nebula."
 The CO emission reveals a chunipy rine. resenibliug to that of the ionized nebula.," The CO emission reveals a clumpy ring, resembling to that of the ionized nebula."
 We observed the 7—2 1 aud 1 0 lines of @CO in three positious aud detected emission in all cases., We observed the $J=2$ –1 and $1$ –0 lines of $^{13}$ CO in three positions and detected emission in all cases.
 Fie., Fig.
 ] shows he spectra of the four CO and CO lines observed toward the strongest peal: im the molecular envelope (offsetLO”. ο from the central position).," 1 shows the spectra of the four CO and $^{13}$ CO lines observed toward the strongest peak in the molecular envelope (offset, ] from the central position)."
" The measured isotopic ratio toward this position. where hne S/N is the highest. is =22,"," The measured isotopic ratio toward this position, where the S/N is the highest, is $=22$."
 This value is iu agreement with a previous estimate of Bachiller ct al. (, This value is in agreement with a previous estimate of Bachiller et al. (
1997) are with the values obtained oward the two oositious with lower S/N spectra.,1997) and with the values obtained toward the two positions with lower S/N spectra.
 Thus. our data provide 10 evidence for variations of the isotopic ratio across the rebula.," Thus, our data provide no evidence for variations of the isotopic ratio across the nebula."
 We have also detected a line around the ο. =1 frequency in the ceutral position of NCC 65123. but we aded to detect the J=2 1 line at relatively low levels.," We have also detected a line around the $^{12}$ CO $J = 1$ --0 frequency in the central position of NGC 6543, but we failed to detect the $J =
2$ –1 line at relatively low levels."
 This indicates that the line near the J=1 0 frequency is probably not duc to CO., This indicates that the line near the $J=1$ –0 frequency is probably not due to CO.
 It is interesting to recall that he [38a recombination line is only separated by 3 MIIZ CCS kins +) from the CO 7=1 0 line., It is interesting to recall that the $\alpha$ recombination line is only separated by 3 MHz (7.8 km $^{-1}$ ) from the $^{12}$ CO $J = 1$ –0 line.
 As discussed iu Bachiller et al. (, As discussed in Bachiller et al. (
1992). the II386 line cau dominate the enmüssiou around the -CO J=1 0 frequency in some uebulae with little or no molecular gas.,"1992), the $\alpha$ line can dominate the emission around the $^{12}$ CO $J = 1$ –0 frequency in some nebulae with little or no molecular gas."
 We believe that this is the case in the central position of NGC 6513., We believe that this is the case in the central position of NGC 6543.
 The pariuneters of the detected. sources aud the derived isotopic ratios are listed in Table 2., The parameters of the detected sources and the derived isotopic ratios are listed in Table 2.
 We have also included iu this table data for NGC 2316. NGC 7293. NGC 6781. AI [-9. and CRL 6158 from previous observatious by Dachiller et al. (," We have also included in this table data for NGC 2346, NGC 7293, NGC 6781, M 4-9, and CRL 618 from previous observations by Bachiller et al. ("
1989. 1997).,"1989, 1997)."
 The spectra of the remaining PNe are shown in Fie., The spectra of the remaining PNe are shown in Fig.
 1 (for M 1-16) aud iu Fig., 1 (for M 1-16) and in Fig.
 2., 2.
" Iu four cases (NGC F008, NCC 6853. ME 1-13. |3073639) we obtained teutative detections iu the CO lines. or ouly crude upper limits could be derived."," In four cases (NGC 7008, NGC 6853, M 1-13, $+30^\circ 3639$ ) we obtained tentative detections in the $^{13}$ CO lines, or only crude upper limits could be derived."
 The assmuptioun of optically thin enuüsson used to clotermine he isotopic ratio is likely to be accurate in the case of evolved nebulae like the Rine. the IHelix. NGC 2316. ete.," The assumption of optically thin emission used to determine the isotopic ratio is likely to be accurate in the case of evolved nebulae like the Ring, the Helix, NGC 2346, etc."
 Iu fact. Laree-Velocity-Gradicut (LVG) models confirm that the CO and CO line euission is optically thin iu these cases;," In fact, Large-Velocity-Gradient (LVG) models confirm that the $^{12}$ CO and $^{13}$ CO line emission is optically thin in these cases."
 Ou the other hand. the same asstuuption is not appropriate for the CO lines in voune objects such as CRE 2688. CRL 618. NGC 7027. aud M1-16.," On the other hand, the same assumption is not appropriate for the CO lines in young objects such as CRL 2688, CRL 618, NGC 7027, and M 1-16."
 Iu such cases. the derived. CÓ. column deusities andl ?CO/U CÓ cohuun deustv ratios represent oulv approximate lower limits.," In such cases, the derived CO column densities and $^{12}$ $^{13}$ CO column density ratios represent only approximate lower limits."
 One could also have weak emission arising fou small optically thick chumps very diluted within the CO aud 15CO beams., One could also have weak emission arising from small optically thick clumps very diluted within the $^{12}$ CO and $^{13}$ CO beams.
 This could be the case in some compact CO envelopes such as the Butterfly nebula. M 2-9.," This could be the case in some compact CO envelopes such as the Butterfly nebula, M 2-9."
 The CO in AI 2-9 is concentrated in an expanding clumpy ring which has a mean diameter of 6 arcsec (Zweigle ct al., The CO in M 2-9 is concentrated in an expanding clumpy ring which has a mean diameter of 6 arcsec (Zweigle et al.
 1997)., 1997).
 Tudividual chumps in this ring have sizes - Laresec., Individual clumps in this ring have sizes $<4$ arcsec.
 The weakness ofthe 12 CO aud CO lines we observe could be due to the iniportaut dilation of such small clamps witlin the 30-11 bean., The weakness of the $^{12}$ CO and $^{13}$ CO lines we observe could be due to the important dilution of such small clumps within the 30-m beam.
 The chuups could be optically thick in CO. and the reported 1° CO/P CO intensity ratio would just represent a lower lt to the abundance ratio iu these colpact PNe.," The clumps could be optically thick in CO, and the reported $^{12}$ $^{13}$ CO intensity ratio would just represent a lower limit to the abundance ratio in these compact PNe."
 Tn the case of extended. PNe (eg. NGC 6720. NGC 2516. NGC. 7293). the asstuuption of optically thin CO 21 cnussion is based ou detailed studies of individual nebulae which show that the intrinsic 21/10," In the case of extended PNe (e.g. NGC 6720, NGC 2346, NGC 7293), the assumption of optically thin CO 2–1 emission is based on detailed studies of individual nebulae which show that the intrinsic 2–1/1–0"
and the surface N(Li) reaches a value of ~0.9 (then the star has a total luminosity Log(L/L.) between ~3.0 and 3.6).,and the surface N(Li) reaches a value of$\sim$ 0.9 (then the star has a total luminosity $_{\odot}$ ) between $\sim$ 3.0 and 3.6).
 We thus confirm. but this time at solar metallicity. the finding by that thermohaline mixing does increase the surface Li abundance in low-mass TP-AGB stars.," We thus confirm, but this time at solar metallicity, the finding by that thermohaline mixing does increase the surface Li abundance in low-mass TP-AGB stars."
 Whether these Li-rich objects can simultaneously be carbon-rich will require further investigation of the TP-AGB phase including parametric convective overshoot as mentioned before., Whether these Li-rich objects can simultaneously be carbon-rich will require further investigation of the TP-AGB phase including parametric convective overshoot as mentioned before.
 As can be seen in Table 3.. the total stellar vields for lithium remain negative.," As can be seen in Table \ref{tableNLiTPAGB}, the total stellar yields for lithium remain negative."
 Let us now discuss the case of a 1.25 M. model that does take into account both thermohaline instability and rotation-induced mixing as described in 2.2 and 2.3., Let us now discuss the case of a 1.25 $_{\odot}$ model that does take into account both thermohaline instability and rotation-induced mixing as described in 2.2 and 2.3.
 As described above. the thermohaline instability induced by *He-burning sets in only on the RGB after the bump.," As described above, the thermohaline instability induced by $^3$ He-burning sets in only on the RGB after the bump."
 Before a star reaches that phase. the modifications of its internal and surface chemical abundances are thus driven by rotation-induced mixing on the main sequence and convective dilution during the first dredge-up episode on the subgiant branch and early-RGB.," Before a star reaches that phase, the modifications of its internal and surface chemical abundances are thus driven by rotation-induced mixing on the main sequence and convective dilution during the first dredge-up episode on the subgiant branch and early-RGB."
 The predictions of our rotating models up to that phase have been extensively tested in previous papers., The predictions of our rotating models up to that phase have been extensively tested in previous papers.
 They account nicely for the behaviour of lithium and beryllium at the surface of Population I main-sequence and subgiant stars22222002)., They account nicely for the behaviour of lithium and beryllium at the surface of Population I main-sequence and subgiant stars.
 Rotation-induced mixing has also an impact on the internal abundance profiles of heavier chemicals involved in hydrogen-burning at higher temperatures than the fragile Li and Be., Rotation-induced mixing has also an impact on the internal abundance profiles of heavier chemicals involved in hydrogen-burning at higher temperatures than the fragile Li and Be.
 This can be seen at the moment of the turnoff in Fig., This can be seen at the moment of the turnoff in Fig.
 2. for the 1.25 star computed for different initial rotation velocities., \ref{fig:profil_1.25_diffrot_Tof} for the 1.25 $_{\odot}$ star computed for different initial rotation velocities.
 In the rotating models. the abundance gradients are smoothed out compared to the standard case: *He. C. N. and |Ο diffuse outwards. while C and !*O diffuse inwards.," In the rotating models, the abundance gradients are smoothed out compared to the standard case: $^3$ He, $^{13}$ C, $^{14}$ N, and $^{17}$ O diffuse outwards, while $^{12}$ C and $^{18}$ O diffuse inwards."
 However. rotation-induced mixing is not efficient enough to noticeably change the surface abundances of these elements while on the main sequence for the 1.25 M.model?.. although it sets the scene for abundance variations in latter evolution phases.," However, rotation-induced mixing is not efficient enough to noticeably change the surface abundances of these elements while on the main sequence for the 1.25 $_{\odot}$, although it sets the scene for abundance variations in latter evolution phases."
 In particular the surface abundance variations during the first dredge-up are slightly strenghtened wher rotation-induced mixing is accounted for. as shownin Fig. 6..," In particular the surface abundance variations during the first dredge-up are slightly strenghtened when rotation-induced mixing is accounted for, as shownin Fig. \ref{fig:surfaceabundances1p25}."
 For example. more *He is brought into the stellar envelope. and the post dredge-up PeySC and 'O/'70 ratios are lower than in the non-rotating case.," For example, more $^3$ He is brought into the stellar envelope, and the post dredge-up $^{12}$ $^{13}$ C and $^{16}$ $^{17}$ O ratios are lower than in the non-rotating case."
 We note. however. that neither '°O nor ?Na are affected. and that no Να enhancement is expected at the surface of such a low-mass star during the first dredge-up.," We note, however, that neither $^{16}$ O nor $^{23}$ Na are affected, and that no $^{23}$ Na enhancement is expected at the surface of such a low-mass star during the first dredge-up."
 The modifications of the internal abundances due to rotation-induced mixing that we just discussed for the 1.25 M. model do induce slight variations of the overal stellar structure and of the evolutionary track., The modifications of the internal abundances due to rotation-induced mixing that we just discussed for the 1.25 $_{\odot}$ model do induce slight variations of the overal stellar structure and of the evolutionary track.
 In the case of such a low-mass star. the impact on the effective temperature and lummosity along the RGB is relatively modest as can be seen in Fig. 8..," In the case of such a low-mass star, the impact on the effective temperature and luminosity along the RGB is relatively modest as can be seen in Fig. \ref{fig:logL_logTeff_vs_t_1p25_v0_v110}."
 We note that in the case of the computation without rotation (black full line) the drop in lummosity at the RGB bump (that is associated with a slight increase in Teff at ~5.8«10° years) is clearly noticeable., We note that in the case of the computation without rotation (black full line) the drop in luminosity at the RGB bump (that is associated with a slight increase in Teff at $\sim 5.8 \times 10^9$ years) is clearly noticeable.
 In the rotating case the inflexion in luminosity at the bump ts less pronounced and occurs at slightly lower luminosity., In the rotating case the inflexion in luminosity at the bump is less pronounced and occurs at slightly lower luminosity.
 Consequently. the stellar luminosity at which the thermohaline i1stability reaches the convective envelope is also slightly lower in the rotating case (see Table 1).," Consequently, the stellar luminosity at which the thermohaline instability reaches the convective envelope is also slightly lower in the rotating case (see Table 1)."
 Figure 9 shows. at the evolutionpoint οὓς for the 1.25 M. model computed with an mitial rotation velocity of 110 km.see7!. the diffusion coefficient associated to the thermohaline instability. Dij (Eq.," Figure \ref{fig:Dtot_Dthc_1p25_105_deltaM} shows, at the evolutionpoint $_{1.25}$ for the 1.25 $_{\odot}$ model computed with an initial rotation velocity of 110 $^{-1}$ , the diffusion coefficient associated to the thermohaline instability, $_{\rm thc}$ (Eq."
" | for C,= 1000). and the total diffusion coefficient associated to rotation. Djq. that characterizes the transport of chemicals through the interaction of meridional circulation and shear turbulence ?"," 1 for $_{\rm t}=1000$ ), and the total diffusion coefficient associated to rotation, $_{\rm rot}$ , that characterizes the transport of chemicals through the interaction of meridional circulation and shear turbulence ."
mainBodyCitationEnd4754]Zahn92. Diye is five to six orders of magnitude higher than Dj. and this result is independent of the initial rotation velocity on the ZAMS.," $_{\rm thc}$ is five to six orders of magnitude higher than $_{\rm rot}$ , and this result is independent of the initial rotation velocity on the ZAMS."
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" Specifically, we solve the coupled linearized collisionless Boltzimanu equation for the distribution function of the primary stellar svstem aud Poissou's equation (cf."," Specifically, we solve the coupled linearized collisionless Boltzmann equation for the distribution function of the primary stellar system and Poisson's equation (cf."
 82) where the subscript 0 denotes the equilibrium quantities aud the subscript 1 denotes the first order perturbation of the distribution function f£. audWamiltonian Z7.," 2) where the subscript $0$ denotes the equilibrium quantities and the subscript $1$ denotes the first order perturbation of the distribution function $f$, andHamiltonian $H$."
" The perturbed potential. $4. is the sum of the tidal poteutial due to the perturbor. ®,. aud the gravitational potential of the response to the perturbation. $,,,,."," The perturbed potential, $\Phi_1$, is the sum of the tidal potential due to the perturber, $\Phi_p$, and the gravitational potential of the response to the perturbation, $\Phi_{resp}$."
 After performing the Laplace transform in time. the Fourier trausforii in the angle variables aud solving for fi. equation (1)) becomes ry where we denote the Laplace transtorm of a quantity y(t) bv g(5) aud the Fourier transform of a quantity yow) by Wy with y=(54.52.53) indicating the vector of integer indices for the discrete Fourier expansion in the angle variables.," After performing the Laplace transform in time, the Fourier transform in the angle variables and solving for $f_1$, equation \ref{veq}) ) becomes , where we denote the Laplace transform of a quantity $y(t)$ by $\tilde{y}(s)$ and the Fourier transform of a quantity $y({\bf w})$ by $y_{\gb}$ with ${ \gb}=(\g_1,\g_2,\g_3)$ indicating the vector of integer indices for the discrete Fourier expansion in the angle variables."
 The yequencics associated to the augle variables are denoted by w=offy/olk., The frequencies associated to the angle variables are denoted by ${\ob}=\partial H_0/ \partial {\bf I}$.
" Following the approach adopted iu Weinbere (1989). we will use a poteutial-deusitv biorthogonal basis. at,di to expand he perturbed quantities $4 and py."," Following the approach adopted in Weinberg (1989), we will use a potential-density biorthogonal basis, $u^l_i,~d^l_i$ to expand the perturbed quantities $\Phi_1$ and $\rho_1$."
 The pairs ih. di are solutious of the Poisson equation.," The pairs $u_i^l$, $d_i^l$ are solutions of the Poisson equation."
 The basis we have adopted is obtained by uunuerical solution of the associated Stuiu-Liouville problem iu which the lowest order basis functious are ailored. to the backeround equilibrimm model following Weinberg (1999)., The basis we have adopted is obtained by numerical solution of the associated Sturm-Liouville problem in which the lowest order basis functions are tailored to the background equilibrium model following Weinberg (1999).
 The solution to the coupled equation. which determines the respouse of the halo to the perturbation. then reduces to the solution of au algebraic iutegral equation for he coeffiicieuts of the expansion.," The solution to the coupled equation, which determines the response of the halo to the perturbation, then reduces to the solution of an algebraic integral equation for the coefficients of the expansion."
 The subsectious below sketch the main steps in the derivation of the algebraic equation or the coefficieuts of the expansion from equation (A)) (see c.g. Weinhere 1989 for additional details)., The subsections below sketch the main steps in the derivation of the algebraic equation for the coefficients of the expansion from equation \ref{lftr}) ) (see e.g. Weinberg 1989 for additional details).
" We begin by expaudius the perturbing poteutial. iu terms of spherical harmonics 375,,(0.0) (chere 0 aud o denotes the augular coordinates inside the primary stellar svstemi)."," We begin by expanding the perturbing potential, $\Phi_p$, in terms of spherical harmonics $Y_{lm}(\theta,\phi)$ (where $\theta$ and $\phi$ denotes the angular coordinates inside the primary stellar system)."
",. Up to quadrupole order. the potential of a point mass perturber can be written as where O(f)=avctan(Vf/p) is the position auele of the perturber defined as the angle between the axis from the ceuter of the system to the pericentiic position of the perturber aud the axis from the center to the position at time f."," Up to quadrupole order, the potential of a point mass perturber can be written as where $\Theta(t)=\arctan({Vt/p})$ is the position angle of the perturber defined as the angle between the axis from the center of the system to the pericentric position of the perturber and the axis from the center to the position at time $t$."
" We define the usual quantities re=Ην A). kc=πανί,R) where κ andΠare the radial coordinate imside the stellar svsteia and the distance between the ceuter of the svstem aud the perturber. respectively,"," We define the usual quantities $r_{<}=\min(r,R)$ , $r_{>}=\max(r,R)$ where $r$ and$R$are the radial coordinate inside the stellar system and the distance between the center of the system and the perturber, respectively."
halo component) also in the global population of galactic GCs (Harris 1996).,halo component) also in the global population of galactic GCs (Harris 1996).
" Several theoretical models (Larson 1987, Suntzeff and Kraft 1996; D'Ercole et al."," Several theoretical models (Larson 1987, Suntzeff and Kraft 1996; D'Ercole et al."
 2008; Decressin et al., 2008; Decressin et al.
 2008) suggest that currently observed GCs are only a small fraction in mass of the original parent structures where they formed., 2008) suggest that currently observed GCs are only a small fraction in mass of the original parent structures where they formed.
 In Carretta et al. (, In Carretta et al. (
2009c) we proposed a formation scenario for GCs where a so-called precursor raises the metallicity to a level which corresponds to the primordial abundances currently observed in a cluster.,2009c) we proposed a formation scenario for GCs where a so-called $precursor$ raises the metallicity to a level which corresponds to the primordial abundances currently observed in a cluster.
" In other words, in present-day GCs we are not seeing the contribution of massive SNe: the level of iron and α--εἰεπιεπίς is already homogeneously established in the gas by the explosion of core-collapse SNe in the precursor/proto-cluster."," In other words, in present-day GCs we are not seeing the contribution of massive SNe: the level of iron and $\alpha-$ elements is already homogeneously established in the gas by the explosion of core-collapse SNe in the precursor/proto-cluster."
 Of course the idea of pre-enrichment in large fragments seeding the present galactic GC population is not new (see Searle and Zinn 1978)., Of course the idea of pre-enrichment in large fragments seeding the present galactic GC population is not new (see Searle and Zinn 1978).
" In our scenario, the second effect of SNe in the precursor is to trigger a burst of star formation leading to the build up of the primordial population, a fraction of which we can still observe in GCs, making up about 1/3 of the total cluster stars (see Carretta et al."," In our scenario, the second effect of SNe in the precursor is to trigger a burst of star formation leading to the build up of the primordial population, a fraction of which we can still observe in GCs, making up about 1/3 of the total cluster stars (see Carretta et al."
 2009a)., 2009a).
" However, the correlation we see between iron spread and cluster mass tells us that clusters originating in more massive precursors are probably more capable to retain inhomogeneities in the metallicity plateau established in the previous phase of rapid enrichment."," However, the correlation we see between iron spread and cluster mass tells us that clusters originating in more massive precursors are probably more capable to retain inhomogeneities in the metallicity plateau established in the previous phase of rapid enrichment."
 This is just what is observed for the rms scatter in [Fe/H] as a function of My and of other cluster parameters strictly related to the cluster mass: the maximum temperature reached along the cluster HB (lower panel of Fig. 3;;, This is just what is observed for the $rms$ scatter in [Fe/H] as a function of $M_V$ and of other cluster parameters strictly related to the cluster mass: the maximum temperature reached along the cluster HB (lower panel of Fig. \ref{f:rmsmass};
 see Recio Blanco et al., see Recio Blanco et al.
 2006; Carretta et al., 2006; Carretta et al.
" 2009c), the slope of the cluster global mass function (see Djorgovski et al."," 2009c), the slope of the cluster global mass function (see Djorgovski et al."
" 2003), the velocity dispersion at half-mass radius (Gnedin et al."," 2003), the velocity dispersion at half-mass radius (Gnedin et al."
 2002)., 2002).
" We note that a (small) part of the higher values for the spread in more massive clusters may also be due to variations in He content, which are expected to be larger in GCs of larger masses (Gratton et al."," We note that a (small) part of the higher values for the spread in more massive clusters may also be due to variations in He content, which are expected to be larger in GCs of larger masses (Gratton et al."
 2009)., 2009).
" In this case, of course, the effect is due not only to the higher capability to retain ejecta but also to the enhanced He in some stars formed by gas heavily polluted in H-processed material, which lowers the denominator in the [Fe/H] ratio (see also Bragaglia et al."," In this case, of course, the effect is due not only to the higher capability to retain ejecta but also to the enhanced He in some stars formed by gas heavily polluted in H-processed material, which lowers the denominator in the [Fe/H] ratio (see also Bragaglia et al."
 2009)., 2009).
 We can check the relevance of this effect by computing the iron spread using only stars in the primordial component P (see Carretta et al., We can check the relevance of this effect by computing the iron spread using only stars in the primordial component P (see Carretta et al.
" 20092), which are not expected to be formed from matter polluted with additional He."," 2009a), which are not expected to be formed from matter polluted with additional He."
 We found that the difference in the rms of [Fe/H] (in the sense all stars minus P stars) is on average —0.002x0.003 (with o=0.012 for 19 GCs)., We found that the difference in the $rms$ of [Fe/H] (in the sense all stars minus P stars) is on average $-0.002\pm0.003$ (with $\sigma=0.012$ for 19 GCs).
" For the two most massive clusters in the sample (NGC 6388 and NGC 6441) this difference goes in opposite directions and the largest differences are found for the three GCs with lower quality spectra (NGC 6388 and NGC 6441 again, plus NGC 6752)."," For the two most massive clusters in the sample (NGC 6388 and NGC 6441) this difference goes in opposite directions and the largest differences are found for the three GCs with lower quality spectra (NGC 6388 and NGC 6441 again, plus NGC 6752)."
 We conclude that the contribution to the spread in iron caused by variations in the He content is probably very small., We conclude that the contribution to the spread in iron caused by variations in the He content is probably very small.
 A support for the proposed scenario comes from a second set of (anti)correlations involving the observed spread of iron and the position of GCs in the Galaxy (or related quantities; see Tab., A support for the proposed scenario comes from a second set of (anti)correlations involving the observed spread of iron and the position of GCs in the Galaxy (or related quantities; see Tab.
 2 and Fig. 4))., \ref{t:corrrmsg} and Fig. \ref{f:rmsposiz}) ).
" The dispersion in [Fe/H] is larger in GCs spending, on average, more time at larger distances from the Galactic centre (upper left panel)."," The dispersion in [Fe/H] is larger in GCs spending, on average, more time at larger distances from the Galactic centre (upper left panel)."
" At the same time, the spread decreases for older clusters and for increasing abundances of the a—element?."," At the same time, the spread decreases for older clusters and for increasing abundances of the $\alpha-$."
. In Carretta et al. (, In Carretta et al. (
"2009c) we emphasised that inner halo clusters preferentially include more massive clusters with respect to disc clusters; moreover, the first are also younger, on average, and with a slightly smaller level of a—elements than that found in disc/bulge GCs.","2009c) we emphasised that inner halo clusters preferentially include more massive clusters with respect to disc clusters; moreover, the first are also younger, on average, and with a slightly smaller level of $\alpha-$ elements than that found in disc/bulge GCs."
 All these features concur to explain the good correlation and the anti-correlations displayed in Fig. 4.., All these features concur to explain the good correlation and the anti-correlations displayed in Fig. \ref{f:rmsposiz}.
" Finally, in Fig."," Finally, in Fig."
" 5 we show the run of the rms scatter in Fe as a function of the fraction of first generation stars P, as well as the second generation components of intermediate modified (D), and extreme modified (E) composition (see Carretta et al."," \ref{f:rmspie} we show the run of the $rms$ scatter in Fe as a function of the fraction of first generation stars P, as well as the second generation components of intermediate modified (I), and extreme modified (E) composition (see Carretta et al."
 2009a for detailed definitions of these stellar populations)., 2009a for detailed definitions of these stellar populations).
" The dispersion in Fe increases along with the fraction of P stars still observed in GCs, again pointing toward a deeper potential well in more massive clusters, more able to retain both SN ejecta and low mass P stars."," The dispersion in Fe increases along with the fraction of P stars still observed in GCs, again pointing toward a deeper potential well in more massive clusters, more able to retain both SN ejecta and low mass P stars."
" On the other hand, we know that the fraction of E stars is larger in more massive clusters, and this explains also the correlation between the spread and the E component."," On the other hand, we know that the fraction of E stars is larger in more massive clusters, and this explains also the correlation between the spread and the E component."
 The complementarity between I and fractions (Carretta et al., The complementarity between I and E fractions (Carretta et al.
 2009c) well accounts for the anti-correlationE shown in the middle panel of Fig. 5.., 2009c) well accounts for the anti-correlation shown in the middle panel of Fig. \ref{f:rmspie}. .
"shown in Figures 8--10, are broadly consistent with those expected for SNIa and SNcc.","shown in Figures \ref{hosttypefig}- \ref{colormagfig}, are broadly consistent with those expected for SNIa and SNcc."
 Figure 8 shows AMs7 as a function of the photometric host types (Ilbertetal.2006)., Figure \ref{hosttypefig} shows $\Delta M_{570}$ as a function of the photometric host types \citep{ilbert2006}.
". As expected for a sample dominated by SNcc, the faint events have relatively fewer early-type hosts (19/152) compared to 24/69 for the bright events."," As expected for a sample dominated by SNcc, the faint events have relatively fewer early-type hosts (19/152) compared to 24/69 for the bright events."
 Figure 9 shows AMs;o as a function of τω/(1 z)., Figure \ref{falltimefig} shows $\Delta M_{570}$ as a function of $\tau_{fall}/(1+z)$.
" As with low redshift SNcc (Richardsonetal. 2002)., about half (47/108) the faint events have vfai/(1 z)> days, characteristic of plateau SNII and significantly longer than fall times for SNIa, 20«τω/(1+z)«30days."," As with low redshift SNcc \citep{richardson}, about half (47/108) the faint events have $\tau_{fall}/(1+z)>50\,{\rm days}$ , characteristic of plateau SNII and significantly longer than fall times for SNIa, $20<\tau_{fall}/(1+z)<30\,{\rm days}$."
" Finally, Fig."," Finally, Fig."
 10 shows the color-magnitude diagram using the AB magnitude at 450nm in the rest frame: The SNIa candidates have a narrower color distribution than the SNcc candidates.," \ref{colormagfig} shows the color-magnitude diagram using the AB magnitude at $450\,{\rm nm}$ in the rest frame: The SNIa candidates have a narrower color distribution than the SNcc candidates."
 SNLS did not have sufficient telescope time to obtain spectra of all SNIa candidates., SNLS did not have sufficient telescope time to obtain spectra of all SNIa candidates.
" Therefore, in order to define a more complete SNIa sample, the four-band light curves of all events were compared to SALT2 SNIa template light curves (Guyetal2007).."," Therefore, in order to define a more complete SNIa sample, the four-band light curves of all events were compared to SALT2 SNIa template light curves \citep{guysalt}."
" The SALT2 model characterizes light curves by four parameters: the date of maximum in the rest-frame B-band, the maximum flux in the rest-frame B band, a “color” parameter roughly equivalent to rest-frame B-V, and a “stretch” parameter that dilates the event time scale."," The SALT2 model characterizes light curves by four parameters: the date of maximum in the rest-frame B-band, the maximum flux in the rest-frame B band, a “color” parameter roughly equivalent to rest-frame B-V, and a “stretch” parameter that dilates the event time scale."
 The light curves were fit for these parameters imposing the host photometric redshift., The light curves were fit for these parameters imposing the host photometric redshift.
 Events were “photometrically” classified as SNIa if the four-band fit was reasonable (y?/dof< 10) and if fit parameters corresponded to normal SNIa., Events were “photometrically” classified as SNIa if the four-band fit was reasonable $\chi^2/dof<10$ ) and if fit parameters corresponded to normal SNIa.
" In particular, cuts were appliedto the rise and fall times, the color, c, and to the position in the two color magnitude diagrams, (g'—i"")vs.g' and (r'— z)vs.z."," In particular, cuts were appliedto the rise and fall times, the color, $c$ , and to the position in the two color magnitude diagrams, $(g^\prime-i^\prime)\, vs.\, g^\prime$ and $(r^\prime-z^\prime)\, vs.\, z^\prime$ ."
jet is at its optimum anele. resulting in maximum apparent velocity. which will generally not be the case.,"jet is at its optimum angle, resulting in maximum apparent velocity, which will generally not be the case."
 We have established above that it will be practically impossible to do more than place a lower limit on the Lorentz [actor of a relativistic jet (rom proper motions. whether one- or two-sided.," We have established above that it will be practically impossible to do more than place a lower limit on the Lorentz factor of a relativistic jet from proper motions, whether one- or two-sided."
 However. the proper motions themselves can be used to make a distance estimate to the source. more accurate the more relativistic the jet intrinsically is.," However, the proper motions themselves can be used to make a distance estimate to the source, more accurate the more relativistic the jet intrinsically is."
 What else can we learn from the proper motions?, What else can we learn from the proper motions?
 As already stated. the ratio of proper motions is also the ratio of Doppler factors.," As already stated, the ratio of proper motions is also the ratio of Doppler factors."
 This may be useful in associated unidentified lines with a jet. even though the absolute value of the Doppler shift cannot be predicted. (rom the proper motions.," This may be useful in associated unidentified lines with a jet, even though the absolute value of the Doppler shift cannot be predicted from the proper motions."
" For example. to check if an unidentified. feature at wavelength A, is a Doppler shifted. line of from. the approaching Jet. look at wavelength Ap=Auc5flap[νου LO try and find the line from the receding jet (obviously more dillieult since it will be Doppler de-boosted)"," For example, to check if an unidentified feature at wavelength $\lambda_a$ is a Doppler shifted line of from the approaching jet, look at wavelength $\lambda_r = \lambda_a
\times \mu_{\rm app} / \mu_{\rm rec}$ to try and find the line from the receding jet (obviously more difficult since it will be Doppler de-boosted)."
 There is even a possibility to achieve the goal of limiting the Lorentz factor. in the case of a jet whose angle to the line of sight changes. for example due to precession.," There is even a possibility to achieve the goal of limiting the Lorentz factor, in the case of a jet whose angle to the line of sight changes, for example due to precession."
 “Phis can be seen from Fig 2. where at lower Lorentz factors the (auus is quite à strong function of angle. whereas it is not at all for the higher Lorentz factors.," This can be seen from Fig 2, where at lower Lorentz factors the $d_{\rm max}$ is quite a strong function of angle, whereas it is not at all for the higher Lorentz factors."
 This is illustrated. in Fig 5. in which the product. fpτος is plotted for varving angles. for different Lorentz factors.," This is illustrated in Fig 5, in which the product $\mu_{\rm
app} \mu_{\rm rec}$ is plotted for varying angles, for different Lorentz factors."
 Apart from the smallest angles to the line of sight. at which two-sided proper motions are anvway unlikely to be detected. the most. relativistic jets have an almost constant value of this product. whereas the slower jets have a significantly varving product.," Apart from the smallest angles to the line of sight, at which two-sided proper motions are anyway unlikely to be detected, the most relativistic jets have an almost constant value of this product, whereas the slower jets have a significantly varying product."
 For example. a jet with a mean angle to the line of sight. of 60 degrees. precessing with a hallopenine angle of 20 degrees would. produce a 25% change in the product frpftoec Over its precession period if it had. an intrinsic Lorentz factor of 2.," For example, a jet with a mean angle to the line of sight of 60 degrees, precessing with a half-opening angle of 20 degrees would produce a $\sim 25$ change in the product $\mu_{\rm app} \mu_{\rm
rec}$ over its precession period if it had an intrinsic Lorentz factor of 2."
 We the jet has a Lorentz factor of five or more. the fractional change in the product over the precession evcle is or less.," If the jet has a Lorentz factor of five or more, the fractional change in the product over the precession cycle is or less."
 This approach. albeit almost certainly limited to galactic sources (due to the necessity of tracking in time a periodic precession evele) presents an interesting possibility for limiting the Lorentz factors.," This approach, albeit almost certainly limited to galactic sources (due to the necessity of tracking in time a periodic precession cycle) presents an interesting possibility for limiting the Lorentz factors."
 In this paper the uses and limitations of relativistic jet proper motions have been explored. under the assumption of intrinsically symmetric ejection velocities.," In this paper the uses and limitations of relativistic jet proper motions have been explored, under the assumption of intrinsically symmetric ejection velocities."
 The main results derived. are: As stated above. everything calculated in this paper is only strictly. valid: under the assumption of svmmetrie ejection events.," The main results derived are: As stated above, everything calculated in this paper is only strictly valid under the assumption of symmetric ejection events."
 Observations of jets from the neutron star ARB Sco X-l (Fomalont ct al., Observations of jets from the neutron star XRB Sco X-1 (Fomalont et al.
 2001a. 2001b) have shown us that the resolved sites of radio emission may in some cases simplv be the regions of jetISM interaction ancl may not rellect the underlying bulk velocity of the How.," 2001a, 2001b) have shown us that the resolved sites of radio emission may in some cases simply be the regions of jet–ISM interaction and may not reflect the underlying bulk velocity of the flow."
 This is even more dramatically demonstrated by observations of larec-scale decelerating jets from the black hole transient NTTIZ J1550-564 (Corbel et al., This is even more dramatically demonstrated by observations of large-scale decelerating jets from the black hole transient XTE J1550-564 (Corbel et al.
 2002: Ixaaret. ct al., 2002; Kaaret et al.
 2003: Tomsick et al., 2003; Tomsick et al.
 2003)., 2003).
 As a result. the application of the results in this paper. e.g. estimating the distance dias should.," As a result the application of the results in this paper, e.g. estimating the distance $\sim
d_{\rm max}$ should,"
"which can be used to correct the observed polarization values according to the input signal-to-noise ratio, measured polarization and number of HWP positions.","which can be used to correct the observed polarization values according to the input signal-to-noise ratio, measured polarization and number of HWP positions."
" In general, the bias effect is present even at reasonably high values of SNR when the polarization is small and op and AP tend to become similar, so that the systematic bias correction is comparable to the random uncertainty of the polarization."," In general, the bias effect is present even at reasonably high values of $SNR$ when the polarization is small and $\sigma_P$ and $\Delta P$ tend to become similar, so that the systematic bias correction is comparable to the random uncertainty of the polarization."
" This is better seen in Fig. 2,,"," This is better seen in Fig. \ref{fig:bias},"
 where we have plotted the ratio AP over cp as a function of Po/op=SNRp deduced from our simulations., where we have plotted the ratio $\Delta P$ over $\sigma_P$ as a function of $P_0/\sigma_P\equiv SNR_P$ deduced from our simulations.
" As anticipated, for low values of S.NRp, the ratio between AP and op tends to unity, with some variations among different estimators."," As anticipated, for low values of $SNR_P$, the ratio between $\Delta P$ and $\sigma_P$ tends to unity, with some variations among different estimators."
" For S.NRp 73 the bias correction is less than of the expected accuracy, and it is, therefore, negligible."," For $SNR_P\geq$ 3 the bias correction is less than of the expected accuracy, and it is, therefore, negligible."
" Moreover, above that threshold, all estimators give practically identical results."," Moreover, above that threshold, all estimators give practically identical results."
 In Fig., In Fig.
" 2 we have plotted, for comparison, the function computed by Simmons&Stewart who have used a Maximum Likelihood estimator in (1985),,order to evaluate the bias."," \ref{fig:bias} we have plotted, for comparison, the function computed by \cite{simmons}, who have used a Maximum Likelihood estimator in order to evaluate the bias."
" As these authors have shown, this is the best estimator for SNRp «0.7, while for SNRp >0.7 the mode, first used by Wardle&Kronberg(1974),, should be used."," As these authors have shown, this is the best estimator for $SNR_P\leq$ 0.7, while for $SNR_P>$ 0.7 the mode, first used by \cite{wardle}, should be used."
 Until now we have assumed that one is able to perfectly subtract the background contribution., Until now we have assumed that one is able to perfectly subtract the background contribution.
" This is most likely the case when one is to perform polarimetric measurements on point-like sources, since in that situation local background subtraction is in most cases straightforward."," This is most likely the case when one is to perform polarimetric measurements on point-like sources, since in that situation local background subtraction is in most cases straightforward."
" We remark that the background, whatever its nature is, must be subtracted before the calculation of normalized Stokes parameters (seealsoTinbergen1996),, so that possible background polarization is vectorially removed."," We remark that the background, whatever its nature is, must be subtracted before the calculation of normalized Stokes parameters \cite[see also][]{tinbergen}, so that possible background polarization is removed."
" If we assume that the object is characterized by P, and xo and the background by P, and χι, the two polarization fields can be expressed using the Stokes vectors defined as S,(1,IP;cos2X,IoP,sin2x5) and (Is,IyPycos2x4,IjPsin2x5), where we have neglected S,any circular polarization."," If we assume that the object is characterized by $P_o$ and $\chi_o$ and the background by $P_b$ and $\chi_b$, the two polarization fields can be expressed using the Stokes vectors defined as $\vec{S_o}(I_o,
I_o P_o \cos 2\chi_o, I_o P_o \sin 2\chi_o)$ and $\vec{S_b}(I_b, I_b
P_b \cos 2\chi_b, I_b P_b \sin 2\chi_b)$, where we have neglected any circular polarization."
" Since Stokes vectors are additive (seeforexampleChandrasekhar1950), the resulting polarization field is described by S=S,+S; and, therefore, the total polarization is given by the following formula: where r=(IyPj)/(I5;P5), ie. the ratio between the polarized fluxes of background and object."," Since Stokes vectors are additive \cite[see for
example][]{chandra}, the resulting polarization field is described by $\vec{S}=\vec{S_o}+\vec{S_b}$ and, therefore, the total polarization is given by the following formula: where $r=(I_b P_b)/(I_o P_o)$, i.e. the ratio between the polarized fluxes of background and object."
" The corresponding polarization angle is Clearly, the background is going to influence significantly the object when r=1."," The corresponding polarization angle is Clearly, the background is going to influence significantly the object when $r\gtrsim1$."
" For r~1 one can write: which implies that, for comparable polarized fluxes, the resulting polarization is nulled when the polarization fields areperpendicular (|xo—xp|=7/2)."," For $r\sim1$ one can write: which implies that, for comparable polarized fluxes, the resulting polarization is nulled when the polarization fields are $|\chi_o-\chi_b|=\pi/2$ )."
 One of the basic problems one has to face when reducing the data produced by dual-beam instruments is the flat-fielding., One of the basic problems one has to face when reducing the data produced by dual-beam instruments is the flat-fielding.
" Due to the fact that image splitting occurs after the focal mask, the collimator and the HWP, one would in principle need to obtain flat exposures with all optical components in the light path."," Due to the fact that image splitting occurs after the focal mask, the collimator and the HWP, one would in principle need to obtain flat exposures with all optical components in the light path."
" Unfortunately, in all practical conditions, this introduces strong artificial effects due to the strong polarization typical of flat-field sources (either twilight sky or internal screens)."," Unfortunately, in all practical conditions, this introduces strong artificial effects due to the strong polarization typical of flat-field sources (either twilight sky or internal screens)."
 In principle one can reduce this effect using the continuous rotation of the HWP as a depolarizer., In principle one can reduce this effect using the continuous rotation of the HWP as a depolarizer.
" This is implemented, for example, in EFOSC2, currently mounted at the ESO-3.6m telescope (Patat and it is effective only if the HWP rotation time is much1999) shorter than the required exposure time."," This is implemented, for example, in EFOSC2, currently mounted at the ESO-3.6m telescope \citep{efosc} and it is effective only if the HWP rotation time is much shorter than the required exposure time."
 The depolarizing effect can also be achieved by averaging flats taken the same set of HWP angles used for the scientific exposures., The depolarizing effect can also be achieved by averaging flats taken the same set of HWP angles used for the scientific exposures.
" In fact, with the usage of the optimal angle set (see Sec. 2)),"," In fact, with the usage of the optimal angle set (see Sec. \ref{sec:basics}) ),"
" one has that and a similar expression for fz, which do not contain any polarization information."," one has that and a similar expression for $f_{E,i}$, which do not contain any polarization information."
" The problem is that this is true only if the source is stable in intensity, which is surely not the case for the twilight sky and probably not really true for most lamps."," The problem is that this is true only if the source is stable in intensity, which is surely not the case for the twilight sky and probably not really true for most lamps."
 An alternative solution (at least for imaging) is the usage of a set of twilight flats obtained without HWP and WP., An alternative solution (at least for imaging) is the usage of a set of twilight flats obtained without HWP and WP.
" If on the one hand this eliminates source polarization, on the other hand it does not allow fora proper flat field correction."," If on the one hand this eliminates source polarization, on the other hand it does not allow fora proper flat field correction."
" In fact, while the pixel- variations are properly taken into account, the"," In fact, while the pixel-to-pixel variations are properly taken into account, the"
of a planet like GJ4436b and this adds to the ambiguity of specific interpretations because the age of GJ4436 is quite uncertain.,of a planet like 436b and this adds to the ambiguity of specific interpretations because the age of 436 is quite uncertain.
 More theoretical and observational work is obviously needed to improve our understanding of 4436b and. by extension. general models of planet formation and evolution.," More theoretical and observational work is obviously needed to improve our understanding of 436b and, by extension, general models of planet formation and evolution."
 From an observational standpoint. obtaining and analyzing an additional high-precision transit light curve of 4436b with a different instrument than already used would still be valuable.," From an observational standpoint, obtaining and analyzing an additional high-precision transit light curve of 436b with a different instrument than already used would still be valuable."
 Such data would preferably be obtained in the near-infrared to reduce the impact of stellar variability and uncertainties 1 the limb darkening., Such data would preferably be obtained in the near-infrared to reduce the impact of stellar variability and uncertainties in the limb darkening.
 Another interesting observational study would be to make a direct radius measurement of the host star with interferometry. similar to what has been done for the transiting planet host star HD 189733 (?)..," Another interesting observational study would be to make a direct radius measurement of the host star with interferometry, similar to what has been done for the transiting planet host star HD 189733 \citep{baines07}."
 We have found a similar ratio of the planet and star radit as the previous studies. and so an additional constraint οἱ the stellar radius would give a tighter constraint on the planet radius.," We have found a similar ratio of the planet and star radii as the previous studies, and so an additional constraint on the stellar radius would give a tighter constraint on the planet radius."
 What can be definitively ascertained from comparisor of the observational results and theoretical models 1s. that 4436b has a significant H/He envelope similar to Uranus anc Neptune., What can be definitively ascertained from comparison of the observational results and theoretical models is that 436b has a significant H/He envelope similar to Uranus and Neptune.
 Such similarities in structure suggest that they formec in à similar environment., Such similarities in structure suggest that they formed in a similar environment.
 Additionally. 4436b and the solar system Ice giants are similar in that they likely could not have formed in their current locations.," Additionally, 436b and the solar system ice giants are similar in that they likely could not have formed in their current locations."
 Uranus and Neptune are thought to be too far away from the Sun (?).. while 4436b is likely too close to its host star.," Uranus and Neptune are thought to be too far away from the Sun \citep{levison01}, while 436b is likely too close to its host star."
 One hypothesis for the formation and evolution of Uranus and Neptune is that they were originally in the same region of the protoplanetary disk as Jupiter and Saturn (?).., One hypothesis for the formation and evolution of Uranus and Neptune is that they were originally in the same region of the protoplanetary disk as Jupiter and Saturn \citep{thommes99}.
 The later two planets accreted gas much more quickly and scattered the ice giants to-be outward towards their current locations. which limited their growth.," The later two planets accreted gas much more quickly and scattered the ice giants to-be outward towards their current locations, which limited their growth."
 Uranus and Neptune are therefore really gas giant planet cores that did not accrete gas quickly enough to complete formation., Uranus and Neptune are therefore really gas giant planet cores that did not accrete gas quickly enough to complete formation.
 In contrast. the structure of 4436b likely does not require a comparably violent past because it orbits à M dwarf.," In contrast, the structure of 436b likely does not require a comparably violent past because it orbits a M dwarf."
 ? have shown that formation of gas giants around low mass stars is severely hindered because these stars are expected to have correspondingly low surface density protoplanetary disks and the dynamical timescale of orbiting bodies is longer., \citet{laughlin04} have shown that formation of gas giants around low mass stars is severely hindered because these stars are expected to have correspondingly low surface density protoplanetary disks and the dynamical timescale of orbiting bodies is longer.
 ? predict a dearth of Jovian mass planets due to the resulting slow accretion of gas. but plenty of Neptune and lower mass planets around M dwarfs.," \citet{laughlin04} predict a dearth of Jovian mass planets due to the resulting slow accretion of gas, but plenty of Neptune and lower mass planets around M dwarfs."
 In this context 4436b could be considered a failed gas giant akin to Uranus and Neptune., In this context 436b could be considered a failed gas giant akin to Uranus and Neptune.
 Furthermore. 4436b could have experienced scattering due to another body like Uranus and Neptune even though this probably wasn't required to limit its growth into a larger body.," Furthermore, 436b could have experienced scattering due to another body like Uranus and Neptune even though this probably wasn't required to limit its growth into a larger body."
 The evidence for this comes from its current location. which is far away from its likely formation site; and its observed orbital eccentricity (e = 0.16). which is in direct contradictior to predictions of tidal eircularization. theories (2)...," The evidence for this comes from its current location, which is far away from its likely formation site; and its observed orbital eccentricity $e$ = 0.16), which is in direct contradiction to predictions of tidal circularization theories \citep{maness07}."
 A major dynamical interaction event after formation leading to inward scattering is an explanation that unifies both of these properties into à single evolutionary picture., A major dynamical interaction event after formation leading to inward scattering is an explanation that unifies both of these properties into a single evolutionary picture.
 This interaction could have occurred with the outer object in the system that is indicatec by the long term trend in 4436s radial velocities., This interaction could have occurred with the outer object in the system that is indicated by the long term trend in 436's radial velocities.
 An alternative explanation for 4436b's orbital properties ts disk interaction leading to migration and eccentricity excitation (e.g.??)..," An alternative explanation for 436b's orbital properties is disk interaction leading to migration and eccentricity excitation \citep[e.g.][]{goldreich80, goldreich03}."
 Continued observations of the 4436 system are needed to detect and characterize all of the objects it contains., Continued observations of the 436 system are needed to detect and characterize all of the objects it contains.
 Complementary theoretical studies based on the system census and the likelihood that 4436 had a low-mass disk are also needed to specifically assess the plausibility of the conceivable evolutionary scenarios., Complementary theoretical studies based on the system census and the likelihood that 436 had a low-mass disk are also needed to specifically assess the plausibility of the conceivable evolutionary scenarios.
 The results of our analysis also allow us to search for deviations in transit times that might arise from the gravitational perturbations of another planet in the system., The results of our analysis also allow us to search for deviations in transit times that might arise from the gravitational perturbations of another planet in the system.
 The times determined for each of the five partial transit observations. which are given in Table 2.. exhibit à maximum deviation from the mean ephemeris of ss. The average uncertainty on these times ts ss so the deviations are fully consistent with regular transit times over this observational period.," The times determined for each of the five partial transit observations, which are given in Table \ref{t4}, exhibit a maximum deviation from the mean ephemeris of s. The average uncertainty on these times is s so the deviations are fully consistent with regular transit times over this observational period."
 To search for possible long-term variations we compare our mean transit time with the published transit times from two other epochs., To search for possible long-term variations we compare our mean transit time with the published transit times from two other epochs.
 The previous results come from the initial discovery data presented by ?.. and the analyses of the light curve by GO7. DO7. and SO8 that were discussed above.," The previous results come from the initial discovery data presented by \citet{gillon07b}, and the analyses of the light curve by G07, D07, and S08 that were discussed above."
 It should be noted that the time stamps of the photometry presented by 07. which was also analyzed by SOS. are too late by 33ss (M. Gillon private communication).," It should be noted that the time stamps of the photometry presented by G07, which was also analyzed by S08, are too late by s (M. Gillon private communication)."
 The corrected central transit times for the GO7 and S08 results are 2454280.78148 and HHJD respectively., The corrected central transit times for the G07 and S08 results are 2454280.78148 and HJD respectively.
 We utilized these corrected times for all the investigations described below., We utilized these corrected times for all the invesigations described below.
 The transit times determined by DO7 and GO7 are quite consistent. but disagree with the SOS determined value by more than expected from the given uncertainties (~207).," The transit times determined by D07 and G07 are quite consistent, but disagree with the S08 determined value by more than expected from the given uncertainties $\sim 2\sigma$ )."
 To compare the transit times with the other data we collapse the G07. DO7. and 908 times into a single value by first averaging the GO7 and S08 results. and then averaging this with the DO7 result.," To compare the transit times with the other data we collapse the G07, D07, and S08 times into a single value by first averaging the G07 and S08 results, and then averaging this with the D07 result."
 We adopted the largest given uncertainty from the three analyses dd from D07) as the error in the collapsed transit times., We adopted the largest given uncertainty from the three analyses d from D07) as the error in the collapsed transit times.
 A plot of the observed transit time residuals from our ephemeris (transit time and orbital period) for the three available epochs is shown in Fig. 3.., A plot of the observed transit time residuals from our ephemeris (transit time and orbital period) for the three available epochs is shown in Fig. \ref{f3}.
 Our determined orbital period is consistent with. but more precise than. the value found by ? from analyzing the radial velocities of 4436 only.," Our determined orbital period is consistent with, but more precise than, the value found by \citet{maness07} from analyzing the radial velocities of 436 only."
 We find a difference of 128ss between the representative value from the observation and our own., We find a difference of s between the representative value from the observation and our own.
 The deviations show a nearly even trend with orbital epoch when also considering the lower precision time given by ?., The deviations show a nearly even trend with orbital epoch when also considering the lower precision time given by \citet{gillon07b}.
", This could be a result of an error in the period because the deviations are well within the range expected by the uncertainty in this parameter.", This could be a result of an error in the period because the deviations are well within the range expected by the uncertainty in this parameter.
 Alternatively. the transit timing deviations could be due to perturbations of the host star arising from an additional object in the system.," Alternatively, the transit timing deviations could be due to perturbations of the host star arising from an additional object in the system."
exaniple. assuming 7—30.000 IX. as would be appropriate for an Of star. the bolometric Iuninosity would bo adsser factor of 7 higher.,"example, assuming $T$ =30,000 K, as would be appropriate for an Of star, the bolometric luminosity would be a factor of 7 higher."
 Although aad II2 are very dunudnous stars. it is nonetheless surprising that they are so bright in the radio. mid-mfrared. anc N-ravs.," Although and H2 are very luminous stars, it is nonetheless surprising that they are so bright in the radio, mid-infrared, and X-rays."
 Free-free emission iu the winds of massive. evolving stars stars should make them 12Jx radio sources frou the distance to the Galactic center (e.8..Scuderetal.1998:Duncan&White2002).," Free-free emission in the winds of massive, evolving stars stars should make them mJy radio sources from the distance to the Galactic center \citep[e.g.,][]{scu98,dw02}."
. The poiut-like radio source consistent with lis consistent with a stellar wind., The point-like radio source consistent with is consistent with a stellar wind.
 We can estimate the required mass-loss rate by inverting Eq., We can estimate the required mass-loss rate by inverting Eq.
" Lin Seuderi (1998):: Tf we asstuue the electron temperature iu units of 10! IK is T,=1. the mean atomic weight per free clectron is pns= 1.3. and the terminal velocity of the wind is ος=200 kins as appropriate for an LBV (c.¢..PCreinSeuderiet7al. 1998).. then οςοι]=2.3 ιν nuplies AL=2«10? is"," 4 in \citet{scu98}: If we assume the electron temperature in units of $10^4$ K is $T_4 = 1$, the mean atomic weight per free electron is $\mu_e = 1.3$ , and the terminal velocity of the wind is $v_\infty = 200$ km $^{-1}$ as appropriate for an LBV \citep[e.g., P Cyg in][]{scu98}, then $S_{8.4~{\rm GHz}} = 2.3$ mJy implies $\dot{M} = 2\times10^{-5}$ $^{-1}$."
 This is similar to the mass loss rates observed for LBVs (seealso.e.g...Clarketal.2003).," This is similar to the mass loss rates observed for LBVs \citep[see also, e.g.,][]{cla03}."
.. À higher velocity of οςτο1000 Jan +. as appropriate for an Of stax. would imply a uiass loss rate >10.| B," A higher velocity of $v_{\infty} \approx 1000$ km $^{-1}$, as appropriate for an Of star, would imply a mass loss rate $>10^{-4}$ $^{-1}$."
" This rate is much ligher than is typically seen frou, Of stars (Seuderietal.1998).", This rate is much higher than is typically seen from Of stars \citep{scu98}.
.. However. that we have not measured the radio spectral index of1516.1... so it is possible that most of the flux is nou-thermal. the above mass-loss rates are upper lnüts.," However, that we have not measured the radio spectral index of, so it is possible that most of the flux is non-thermal, the above mass-loss rates are upper limits."
 To determine the nature of the radio emission fromX171516.1.. either the radio spectral index must be measured. or the mass-loss rate mst be obtained from a high-resolution infrared spectrin.," To determine the nature of the radio emission from, either the radio spectral index must be measured, or the mass-loss rate must be obtained from a high-resolution infrared spectrum."
 The radio dux of the extended uchula around IT2 is orders of magnitude too bright to be produced by its stellar wind., The radio flux of the extended nebula around H2 is orders of magnitude too bright to be produced by its stellar wind.
" Zhaoetal.(1993) model it as an nuusually deuse (9,~105 7) ΠΠ regiou ilhuuinated bv a Lyiuan contimuun flux equivalent to that of an OF age malu sequence star.", \citet{zhao93} model it as an unusually dense $n_e \sim 10^3$ $^{-3}$ ) HII region illuminated by a Lyman continuum flux equivalent to that of an O7 zero-age main sequence star.
 The detection of both stars in the MSX imidaufrared survey demonstrates that they are surrounded by wari eas and dust., The detection of both stars in the MSX mid-infrared survey demonstrates that they are surrounded by warm gas and dust.
 Eganctal.(2002) and Clarketal.(2003) have identified dusty ring nebulae arouud several known and candidate LBV stars. and hivpothesize that they were produced by past episodes of rapid lass loss.," \citet{egan02} and \citet{cla03} have identified dusty ring nebulae around several known and candidate LBV stars, and hypothesize that they were produced by past episodes of rapid mass loss."
" Surprisingly. the müd-dufrared counterparts to our Galactic center stars are Z10 and 100 times nore luminous at 11.7 and 21.3 gan. respectively, than the uchulae around the other candidate LDVs."," Surprisingly, the mid-infrared counterparts to our Galactic center stars are $\ga$ 10 and 100 times more luminous at 14.7 and 21.3 $\mu$ m, respectively, than the nebulae around the other candidate LBVs."
 Therefore. we suspect either that the mid-infrared flux near aand II2 represcuts interstellar dust. or that it is the nuresolved Cluission from several massive. dusty stars.," Therefore, we suspect either that the mid-infrared flux near and H2 represents interstellar dust, or that it is the unresolved emission from several massive, dusty stars."
 Observations with theTelescope should resolve this issuc., Observations with the should resolve this issue.
 The detection of aud H2 with js particularly intriguing because Of aud LDV stars are not always bright X-ray sources., The detection of and H2 with is particularly intriguing because Of and LBV stars are not always bright X-ray sources.
 Iu Table L. we tabulate the values of Lx/Lp for several of the massive stars listed iu Morrisetal.(1996) and Figer(1999).," In Table \ref{tab:lxlbol}, we tabulate the values of $L_{\rm X}/L_{\rm Bol}$ for several of the massive stars listed in \citet{mor96} and \citet{fig99}."
. The measurements of Lx were taken with imstrumenuts sensitive to different cnerey ranges. so sole caution should be used in comparing the results. as described im —16 notes to the table.," The measurements of $L_{\rm X}$ were taken with instruments sensitive to different energy ranges, so some caution should be used in comparing the results, as described in the notes to the table."
 In general. we have extrapolated —ie fluxes using conservative assuniptious that serve to lessen the dispersion iu ινρω.," In general, we have extrapolated the fluxes using conservative assumptions that serve to lessen the dispersion in $L_{\rm X}/L_{\rm Bol}$."
 Therefore. we expect the range in these values is not caused by selection effects. but represents mtriusie differences iu the efficiency of X-rav production iu these sources.," Therefore, we expect the range in these values is not caused by selection effects, but represents intrinsic differences in the efficiency of X-ray production in these sources."
 We find that there is at least a factor of 3000 range in the values. aud that the scatter is not correlated with spectral type.," We find that there is at least a factor of 3000 range in the values, and that the scatter is not correlated with spectral type."
" The LBV 7 Car aud the O6 f£ star IID 108 lie near the N-rav-bright end of the distribution with Zx/Lp;=5.3 and 6.2. respectively,"," The LBV $\eta$ Car and the O6 f star HD 108 lie near the X-ray-bright end of the distribution with $L_{\rm X}/L_{\rm Bol} = -5.3$ and $-6.2$, respectively."
 The LBV P. Cve aud the O5 If star ΠΟ 190129 lie at the faint end with Zx/Lp;-«8h., The LBV P Cyg and the O4.5 If star HD 190429 lie at the faint end with $L_{\rm X}/L_{\rm Bol} < -8.5$.
 aand II2 lie at the high end of the scatter in Ly/Lpor in Table L. which raises the question of why theyare so bright in N-vavs.," and H2 lie at the high end of the scatter in $L_{\rm X}/L_{\rm Bol}$ in Table \ref{tab:lxlbol}, which raises the question of why theyare so bright in X-rays."
 There are several possible explanations., There are several possible explanations.
 Oue is that the N-rav. ciittine shocks form at varviug optical depths within the winds of massive stars. so that In sole cases the X-ray cussion is absorbed before it can," One is that the X-ray emitting shocks form at varying optical depths within the winds of massive stars, so that in some cases the X-ray emission is absorbed before it can"
The main finding of this work is the approxiniate ‘equivalence’ between Lian ancl Li.,The main finding of this work is the approximate `equivalence' between $\lk$ and $\li$.
 Intriguinely. a comparison of our results with those ound in previous works (e.g. RSOL) about the correlation )etween the Narrow Line luminosity and the power requirec ο feed the extended lobes. shows that the same behaviour seems to hold in an analogous way (when a dillerent covering factor is assumed) for the Luminosity both in broad and in narrow lines.," Intriguingly, a comparison of our results with those found in previous works (e.g. RS91) about the correlation between the Narrow Line luminosity and the power required to feed the extended lobes, shows that the same behaviour seems to hold in an analogous way (when a different covering factor is assumed) for the luminosity both in broad and in narrow lines."
 Also in the latter case. in fact. one obtains an approximate equivalence between the ionizing radiation and he kinetic power supplied to the lobes or to the extended radio structures.," Also in the latter case, in fact, one obtains an approximate equivalence between the ionizing radiation and the kinetic power supplied to the lobes or to the extended radio structures."
 The first implication of these findings is to reinforce and stress the fact that the channeling of material into jets ds indeed a powerful process in term of the total energy budget of radioloud AGN. and more precisely this is comparable with the amount of the radiative power emitted isotropically.," The first implication of these findings is to reinforce and stress the fact that the channeling of material into jets is indeed a powerful process in term of the total energy budget of radio–loud AGN, and more precisely this is comparable with the amount of the radiative power emitted isotropically."
 In this sense. it is relevant to point out that. in the most accepted scenario on the physical processes operating in the central regions of AGN. Lip is related toa (semi)isotropic and. unbeamed radiation.," In this sense, it is relevant to point out that, in the most accepted scenario on the physical processes operating in the central regions of AGN, $\ll$ is related to a (semi)isotropic and unbeamed radiation."
 Both Iuminosities would. be therefore unallected by beaming and. represen ‘intrinsic’ powers., Both luminosities would be therefore unaffected by beaming and represent `intrinsic' powers.
 As shown in Fig., As shown in Fig.
 1. the energy outpu can reach up to 107 erg ," 1, the energy output can reach up to $^{48}$ erg $^{-1}$."
Furthermore. if Li; ds indeed related. to the power dissipated: during the acerction phase. this finding woutle suggest that there is a close link between the accretion process anc the phenomenon of jets. or. alternatively. tha a common ‘clement’ regulates the amount of. luminosity raciatively clissipatect during the accretion process ancl the power channeled into the jet in kinetic form.," Furthermore, if $\li$ is indeed related to the power dissipated during the accretion phase, this finding would suggest that there is a close link between the accretion process and the phenomenon of jets, or, alternatively, that a common `element' regulates the amount of luminosity radiatively dissipated during the accretion process and the power channeled into the jet in kinetic form."
" This ""element should be indeed tightly linked to the above quantities if it can produce a significant. correlation between the observed variables.", This `element' should be indeed tightly linked to the above quantities if it can produce a significant correlation between the observed variables.
 This deep link has also been suggested. by Saunders (1991) on the basis of the observational similarity found between the luminosity in narrow line and the kinetic power. estimated on the large radio scales (15591: see also Faleke et al.," This deep link has also been suggested by Saunders (1991) on the basis of the observational similarity found between the luminosity in narrow line and the kinetic power, estimated on the large radio scales (RS91; see also Falcke et al."
 1995. Wills Brotherton 1995).," 1995, Wills Brotherton 1995)."
 Xn interesting possibility is that the kev role is plaved by magnetic fields., An interesting possibility is that the key role is played by magnetic fields.
 On one side. recent theories of dise accretion ane observations at high energies favor the idea that most of we accreted power is clissipatecl in a corona above the lisc. due e.g. to magnetic reconnection.," On one side, recent theories of disc accretion and observations at high energies favor the idea that most of the accreted power is dissipated in a corona above the disc, due e.g. to magnetic reconnection."
 On the other hand. )0 most. accredited: possibility for. the collimation anc veceleration of jets involves the presence of magnetic fields.," On the other hand, the most accredited possibility for the collimation and acceleration of jets involves the presence of magnetic fields."
 In fact this hypothesis accounts Lor cdilferent aspects of 10 physics of jets. as well as dillerent. environments. for wir formation.," In fact this hypothesis accounts for different aspects of the physics of jets, as well as different environments for their formation."
" Phe power would be ""extracted from the rotational/gravitational energy of an accretion disc or the lack hole itself. ancl then channeled as kinetic power and/or Povnting Iux (see e.g. the reviews by Dlandford 1993. Sprui 1996 and references therein)."," The power would be `extracted' from the rotational/gravitational energy of an accretion disc or the black hole itself, and then channeled as kinetic power and/or Poynting flux (see e.g. the reviews by Blandford 1993, Spruit 1996 and references therein)."
 Η the magnetic field indeed. plavs an important. par in both processes. then one could expect some sort. of correlation between the energy. dissipated and the kinetic power.," If the magnetic field indeed plays an important part in both processes, then one could expect some sort of correlation between the energy dissipated and the kinetic power."
 In other words. the correlation found. here coulc be the signature of the equivalence of the power generate through accretion and the one extracted from the disk/black hole rotational energv and converted into kinetic form.," In other words, the correlation found here could be the signature of the equivalence of the power generated through accretion and the one extracted from the disk/black hole rotational energy and converted into kinetic form."
 The direct. comparison of the predictions of this models with observational data is still quite limited. ancl substantially concerns jets associated with star-size objects.," The direct comparison of the predictions of this models with observational data is still quite limited, and substantially concerns jets associated with star-size objects."
 In the case of AGN. we can only perform a rough estimate. assuming that the magnetic field is in equipartition with the radiation field. which plausibly dominates the pressure in the inner region of an accretion disk. and that a similar eLlliciency regulates the conversion of the produced energy into radiation and. bull energv.," In the case of AGN, we can only perform a rough estimate, assuming that the magnetic field is in equipartition with the radiation field, which plausibly dominates the pressure in the inner region of an accretion disk, and that a similar efficiency regulates the conversion of the produced energy into radiation and bulk energy."
 This estimate is consistent with obtaining similar values for the kinetic and ionizing power., This estimate is consistent with obtaining similar values for the kinetic and ionizing power.
 For example. for an isolated. Kerr black hole. the spin energy which can be extracted in electromagnetic form is approximately given by (e.g. Blandford 1990) where ec is the specific angular momentum and m=CALfc the gravitational radius (e« m).," For example, for an isolated Kerr black hole, the spin energy which can be extracted in electromagnetic form is approximately given by (e.g. Blandford 1990) where $ac$ is the specific angular momentum and $m=GM/c^2$ the gravitational radius $a < m$ )."
 Phen for a rapidly rotating black hole In this estimate equipartition between the magnetic (Cg) and the radiation energy density (CucLac απο} has been assumed. with a tvpical size for the region of the order of a Schwarzschild radius. A2..," Then for a rapidly rotating black hole In this estimate equipartition between the magnetic $U_B$ ) and the radiation energy density $U_{\rm rad}\simeq L_{\rm acc}/4\pi c R_{\rm
s}^2$ ) has been assumed, with a typical size for the region of the order of a Schwarzschild radius, $R_s$."
" Lo these luminosities are converted. with similar elliciency into kinetic ancl ionizing luminosities. respectively. then Lig,~Lia"," If these luminosities are converted with similar efficiency into kinetic and ionizing luminosities, respectively, then $\lk \sim \li$."
 Clearly. the above evaluation is based on gross approximations and physical hypothesis. which currently. are still at the level of speculation (e.g. Spruit 1996 and references therein).," Clearly, the above evaluation is based on gross approximations and physical hypothesis, which currently are still at the level of speculation (e.g. Spruit 1996 and references therein)."
 ]t is worth noticing that if energy is extracted from a spinning objects. it can propagate in electromagnetic and/or kinetic form.," It is worth noticing that if energy is extracted from a spinning objects, it can propagate in electromagnetic and/or kinetic form."
 Here we are indeed. assuming that this power ultimately. and on scales smaller than a parsec. is largely converted into kinetic power (from a theoretical. point of view. this balance depends on the boundary. conclitions of the hydromagnetie problem. sec c.g. Beeclman 1994. and references therein)," Here we are indeed assuming that this power ultimately, and on scales smaller than a parsec, is largely converted into kinetic power (from a theoretical point of view, this balance depends on the boundary conditions of the hydromagnetic problem, see e.g. Begelman 1994 and references therein)."
" lt is then conserved as such up to much larger scales. as shown by the comparison with the results of RSOL and the fact that radiative clissipation within the jet itself. Canbeamed. radiation) is a negligible fraction of the estimated Ly, (CE93)."," It is then conserved as such up to much larger scales, as shown by the comparison with the results of RS91 and the fact that radiative dissipation within the jet itself (unbeamed radiation) is a negligible fraction of the estimated $\lk$ (CF93)."
 Further support. to this possibility comes from the recent results by Bowman. Leahy Ixomissaroy (1996). who show that even dissipation caused by entrainment of material in the jet and. the consequent deceleration. cause only a relatively small loss in kinetic power.," Further support to this possibility comes from the recent results by Bowman, Leahy Komissarov (1996), who show that even dissipation caused by entrainment of material in the jet and the consequent deceleration, cause only a relatively small loss in kinetic power."
 Alternatively. one could argue that the found similarity," Alternatively, one could argue that the found similarity"
norm) regularization on the solution (e.g.. Press et 11992).,"norm) regularization on the solution (e.g., Press et 1992)."
 Other linear regularization constraints could also be enforced on he solution., Other linear regularization constraints could also be enforced on the solution.
 A direct non-regularized solution of the linear system is extremely ill-conditioned and will produce a result containing elements £L; of very large absolute value and different sign., A direct non-regularized solution of the linear system is extremely ill-conditioned and will produce a result containing elements $L_j$ of very large absolute value and different sign.
 An IGE model obtained in this way presents the risk of cancellation of signiticant digits when deriving any physical quantity., An MGE model obtained in this way presents the risk of cancellation of significant digits when deriving any physical quantity.
 Fig., Fig.
 4 shows that the maximum relative error in an MGE tit to multiple power-law profiles decreases almost exponentially with increasing number of Gaussians. approximately as fast as οκ€ου....," \ref{fig:mge_error} shows that the maximum relative error in an MGE fit to multiple power-law profiles decreases almost exponentially with increasing number of Gaussians, approximately as fast as $\epsilon\propto e^{-0.4 N}$."
" This ""nis not surprising.ea given: that every Gaussian: diverges exponentially from the power-law starting from the point of closest approach."," This is not surprising, given that every Gaussian diverges exponentially from the power-law starting from the point of closest approach."
 Although the actual numbers will depend on the profiles and on the adopted radial ranges. these examples provide an estimate of the number of Gaussians that are required to reach a given accuracy in the fit. and explain why 10-15 Gaussians generally produce a very good MGE fit to realistic Galaxy profiles.," Although the actual numbers will depend on the profiles and on the adopted radial ranges, these examples provide an estimate of the number of Gaussians that are required to reach a given accuracy in the fit, and explain why 10–15 Gaussians generally produce a very good MGE fit to realistic Galaxy profiles."
 We finally mention that an asymptotic power-law nuclear cusp can also be imposed on the MGE protile as described in Emsellem et (1999)., We finally mention that an asymptotic power-law nuclear cusp can also be imposed on the MGE profile as described in Emsellem et (1999).
 Although this choice breaks the elegance of the MGE parametrization somewhat. this detail does not change in any way the essence of our fitting algorithm.," Although this choice breaks the elegance of the MGE parametrization somewhat, this detail does not change in any way the essence of our fitting algorithm."
 We now extend the 1D algorithm described above to the fit of actual 2D images., We now extend the 1D algorithm described above to the fit of actual 2D images.
 We perform this extension by starting from the obvious requirement that the 2D algorithm provides the same results of the ID case (to be precise the same Lj; and σι} when it is used to fit a 2D image that presents the same major axis profile of the ID case and has perfectly elliptical isophotes with constant axial ratio q' and PA., We perform this extension by starting from the obvious requirement that the 2D algorithm provides the same results of the 1D case (to be precise the same $L_j$ and $\sigma_j$ ) when it is used to fit a 2D image that presents the same major axis profile of the 1D case and has perfectly elliptical isophotes with constant axial ratio $q'$ and PA.
" This goal is achieved by fitting ""in. parallel"" the MGE model of equation (19. in polar coordinates. to a certain number Nas Of photometric profiles. measured along sectors uniformly spaced in angle from the major axis to the minor axis (sectors in the four quadrants are averaged together before fitting)."," This goal is achieved by fitting “in parallel” the MGE model of equation \ref{eq:surf_twist}) ), in polar coordinates, to a certain number $N_{\rm sec}$ of photometric profiles, measured along sectors uniformly spaced in angle from the major axis to the minor axis (sectors in the four quadrants are averaged together before fitting)."
 All profiles include the central pixel and proper integration of the Gaussians on that pixel was performed for an accurate data-model comparison., All profiles include the central pixel and proper integration of the Gaussians on that pixel was performed for an accurate data-model comparison.
 The sampling of the photometry along the sectors is now logarithmically spaced in the radius mn-a7oyEfe. where q' is a representative axial ratio of the galaxy isophotes (see Fig. 59).," The sampling of the photometry along the sectors is now logarithmically spaced in the radius $m'^2=x'^2+y'^2/q'^2$, where $q'$ is a representative axial ratio of the galaxy isophotes (see Fig. \ref{fig:sectors}) )."
 This latter detail ensures that the requirement stated in the previous paragraph is precisely satistied also for flattened objects., This latter detail ensures that the requirement stated in the previous paragraph is precisely satisfied also for flattened objects.
 This way of sampling galaxy photometry is a time proven practice common to most photometry packages (e.g. Jedrzejewski 1987).," This way of sampling galaxy photometry is a time proven practice common to most photometry packages (e.g., Jedrzejewski 1987)."
 In the tit the same relative error is assigned to all measurements as in the 1D case. following equation (15).," In the fit the same relative error is assigned to all measurements as in the 1D case, following equation \ref{eq:chi2new}) )."
 The actual photometric measurements are obtained by selecting sets of pixels on the image. based on their angular position and logarithmic radius ranges.," The actual photometric measurements are obtained by selecting sets of pixels on the image, based on their angular position and logarithmic radius ranges."
 Once the pixels have been selected. 18 luminosity weighted average coordinates of each photometric »oint are evaluated from the pixels themselves and from the original sectors.," Once the pixels have been selected, the luminosity weighted average coordinates of each photometric point are evaluated from the pixels themselves and from the original sectors."
 This ensures that the coordinates of sectors at ye border. or containing masked pixels. are properly computed and. even more importantly. it permits an accurate modeling of qe galaxy surface brightness distribution close to the centre. where ye discrete nature of pixels cannot be neglected.," This ensures that the coordinates of sectors at the border, or containing masked pixels, are properly computed and, even more importantly, it permits an accurate modeling of the galaxy surface brightness distribution close to the centre, where the discrete nature of pixels cannot be neglected."
 In particular surface brightness interpolation is performed in the centre. but the observed pixel values are used in the fit.," In particular surface brightness interpolation is performed in the centre, but the observed pixel values are used in the fit."
 To guarantee a proper sampling of the image. both in the nucleus and at large radii. 24 sectors have been adopted for every decade in radius (this corresponds to the usual factor 1.1 separation between successive radial intervals).," To guarantee a proper sampling of the image, both in the nucleus and at large radii, 24 sectors have been adopted for every decade in radius (this corresponds to the usual factor 1.1 separation between successive radial intervals)."
250 days.,250 days.
 Also plotted in Fig.2 are the approximate aud sienificance levels. caleulated as described by Scarele (1982)) aud Press et al. (1992)):," Also plotted in \ref{lomb} are the approximate, and significance levels, calculated as described by Scargle \cite{scargle}) ) and Press et al. \cite{press}) ):"
 where M ds taken as twice the ΠΙΟ of data points., where M is taken as twice the number of data points.
" The cutire periodogram (Fig.2.. upper panel) shows highly sienificaut peaks at >3004 and approximately τοπ, 29! (this feature is split into two peaks at about 28 and 29°), and 151."," The entire periodogram \ref{lomb}, upper panel) shows highly significant peaks at $\rm > 300^d$ and approximately $\rm 195^d$, $\rm 29^d$ (this feature is split into two peaks at about $\rm 28^d$ and $\rm 29^d$ ), and $\rm 15^d$."
 The two peaks at Ost and 291 are split bv one dav. and the two low-siguificauce peaks on either side of the 15! oeak also correspond to 1 dav.," The two peaks at $\rm 28^d$ and $\rm 29^d$ are split by one day, and the two low-significance peaks on either side of the $\rm 15^d$ peak also correspond to $\pm$ 1 day."
 The lower-frequency region of the periodoeraum (Fig.2.. lower paucl) shows three significaut peaks (cousicerine the peaks around f 0.0028 days.+ to be part of the sane feature).," The lower-frequency region of the periodogram \ref{lomb}, lower panel) shows three significant peaks (considering the peaks around f = 0.0028 $\rm days^{-1}$ to be part of the same feature)."
" All three features are broadened. and ejve Py=66007211007. P5.=5208+30%) and P4=360°+307,"," All three features are broadened, and give $\rm P_1 ~ = ~ 6600^d \pm 1400^d$, $\rm P_2 ~ = ~ 520^d \pm
30^d$, and $\rm P_3 ~ = ~ 360^d \pm 30^d$."
 The approximate separation of the peaks near f£ = 0.0028 davs+ corresponds to about 35 davs., The approximate separation of the peaks near f = 0.0028 $\rm days^{-1}$ corresponds to about 35 days.
 The estimated uncertainties of Po aud Ps. aud the substructure iu P3. are all about 1 mouth.," The estimated uncertainties of $\rm P_2$ and $\rm P_3$, and the substructure in $\rm P_3$, are all about 1 month."
 Thus. the data show signatures of all the periods preseut in the data-takiug process: oue. vear. six mouths. oue month. two weeks aud one day.," Thus, the data show signatures of all the periods present in the data-taking process: one year, six months, one month, two weeks and one day."
 Even though the last is outside the frequency range sampled by the data. the regular samplue eusures partial samplue of this and evenshorter periods.," Even though the last is outside the frequency range sampled by the data, the irregular sampling ensures partial sampling of this and evenshorter periods."
" We conclude that none of these periods is iutriusic o V Ilva. and that the data support oulv two stellar periods: Py=6600"" and Py=530°."," We conclude that none of these periods is intrinsic to V Hya, and that the data support only two stellar periods: $\rm P_1 ~ = ~ 6600^d$ and $\rm P_2
~ = ~ 530^d$."
 The lareeOo uncertainty in the longerC» period is due both to the small umber of eveles which have been supled iu the data and to the presence of the shorter-period variations. both intrinsic aud observational.," The large uncertainty in the longer period is due both to the small number of cycles which have been sampled in the data and to the presence of the shorter-period variations, both intrinsic and observational."
 To investigate lis further. we analyzed the time series directly using the ata-foldiug technique οἱ Sclavarzenbere-Czermy (1989)) who describes the use of one-way analysis of variance for searching for aud ucasuriug the periods of variable stars.," To investigate this further, we analyzed the time series directly, using the data-folding technique of Schwarzenberg-Czerny \cite{schwarz}) ) who describes the use of one-way analysis of variance for searching for and measuring the periods of variable stars."
 The basis of this ucthod is to fok the observed light curve vi(t;i) with some val period P. i.e. to cut the data sequence mto portions P ong and examine the normalized point-to-point variance of the folded data.," The basis of this method is to fold the observed light curve $\rm y_i(t_i)$ with some trial period P, i.e. to cut the data sequence into portions P long and examine the normalized point-to-point variance of the folded data."
 The scatter around the average value of he data y will be close to random regardless of the value of P unless the data stream las some periodicity close to P. The scatter of the folded data about y is then computed or cach value of P aud compared with raudom scatter., The scatter around the average value of the data $\rm \bar{y}$ will be close to random regardless of the value of P unless the data stream has some periodicity close to P. The scatter of the folded data about $\rm \bar{y}$ is then computed for each value of P and compared with random scatter.
 Tf N is the umber of observatious vi(ti). The data are now folded aud binned ito M subsets with Ny poiuts in the biu.," If N is the number of observations $\rm y_i(t_i)$, The data are now folded and binned into M subsets with $\rm N_k$ points in the bin."
 Define vij as the poiut in the bin. aud where aud 0. the test statistic forthe analysis of variance metlod. equals 1 for pure gaussian noise and is ercater than 1 if there is any non- random componcut in the data.," Define $\rm y_{ij}$ as the point in the bin, and where and $\theta$, the test statistic forthe analysis of variance method, equals 1 for pure gaussian noise and is greater than 1 if there is any non- random component in the data."
 The stellar period is found by plotting 0. 1 versus P: values of P for which 0. lis larger than some threshold. correspouding to some confidence level. are considered to be periods of the star.," The stellar period is found by plotting $\theta$ –1 versus P; values of P for which $\theta$ –1 is larger than some threshold, corresponding to some confidence level, are considered to be periods of the star."
 The analvsis of variance code written by A. Udalski (Udalski et al., The analysis of variance code written by A. Udalski (Udalski et al.
 1991) was used to analyze the elt curve of V Uva shown in Fie.l.., 1994) was used to analyze the light curve of V Hya shown in \ref{lightcurve}. .
" Two periods at > significauce ave found: Py=6160+ 1007, and Py=529.164308 (these are final values - see below)."," Two periods at $>$ significance are found: $\rm P_1 = 6160^d \pm 400^d$ , and $\rm P_2
~ = ~ 529.4^d \pm 30^d$ (these are final values - see below)."
 Several spurious, Several spurious
different surveys.,different surveys.
 This paper describes measurements of global stellar mass deusity out to z=6 in the NICMOS Ultra Deep Field., This paper describes measurements of global stellar mass density out to $z=6$ in the NICMOS Ultra Deep Field.
 There are a number of ueasurements of the global stellar mass cleusity in the literature. hudnicketal.(2003).," There are a number of measurements of the global stellar mass density in the literature. \citet{rudnick},"
.. workine with the HDF-South aud the FIRES kx data (Labbéetal.2003).. ujeasured the rest-frame optical properties of galaxies out to z=3 aud from their average properties deduced. their masses.," working with the HDF-South and the FIRES K data \citep{labbe}, measured the rest-frame optical properties of galaxies out to $z=3$ and from their average properties deduced their masses."
 Dickinsonetal.(2003) used the HDF-North with NICMOS J and H observations., \citet{dickmass} used the HDF-North with NICMOS J and H observations.
 They computeL the global luminosity density out z=3. computed the average ratio of galaxies a ilese redshifts aud so deduced the global mass. Foutauaetal.(2001," They computed the global luminosity density out $z=3$, computed the average mass-to-light ratio of galaxies at these redshifts and so deduced the global mass. \citet{k20},"
).. on the other hand computed stellar masses for galaxies individually., on the other hand computed stellar masses for galaxies individually.
 They used the N20 survey. which consists of UBVIBIzJIx oJHEvations with a limiting magnitude of Ayo=20 and spectroscopic redshifts. out to 2=2.," They used the K20 survey, which consists of UBVRIzJK observations with a limiting magnitude of $K_{\rm Vega}=20$ and spectroscopic redshifts, out to $z=2$."
 Clazebrookοἱal.(2001) used the GDDS survey (Abrahametal.2001).. which is also based on muti-wavelength data aud spectroscopic redshilts but exteuds slightly deeper auc is specially tuued to lie waveleugth rauge 0.8<z« 2.," \citet{gddsmass} used the GDDS survey \citep{gddsdata}, which is also based on multi-wavelength data and spectroscopic redshifts but extends slightly deeper and is specially tuned to the wavelength range $0.8<z<2$ ."
 Droryetal.(2001). used the MUNICS survey (Droryetal.2001 which is shallower (Xv= 19.5) and extends only out to z=1.1.," \citet{drory2004} used the MUNICS survey \citep{munics1}
 which is shallower $K_{\rm Vega}=19.5$ ) and extends only out to $z=1.4$."
 The highest redshift survey so far is that of Droryetal.(2005).. who used the GOODS-South data aud Fors Deep Field (Heidtetal.2003) to go to z25 and Aap=25.1.," The highest redshift survey so far is that of \citet{drory2005}, who used the GOODS-South data and Fors Deep Field \citep{fors} to go to $z=5$ and $K_{AB}=25.4$."
 In Section 2. we describe how we resampled all the various unagiuig data ancl produced the catalog of photometry and photometric redshifts., In Section \ref{sec:data} we describe how we resampled all the various imaging data and produced the catalog of photometry and photometric redshifts.
 In Section 3. we derive masses aid mass functions for the galaxies in the sample aud examine the evolution of the global stellar mass density. with time., In Section \ref{sec:masses} we derive masses and mass functions for the galaxies in the sample and examine the evolution of the global stellar mass density with time.
 In Section | we compare our measurements of the global stellar mass density evolution with the Hartwick(2001) iuodel. the preclictious [rom the star formation rate deusity. auc previous work.," In Section \ref{sec:discussion} we compare our measurements of the global stellar mass density evolution with the \citet{hartwick2004} model, the predictions from the star formation rate density, and previous work."
" Throughout this paper we use the AB magnitude system (OkeL971) aud adopt the concordance cosinologv of Hy=70kins'Adpe1. Q,,=0.3. Q4=0.7."," Throughout this paper we use the AB magnitude system \citep{abmag} and adopt the concordance cosmology of $H_0=70~km s^{-1} Mpc^{-1}$, $\Omega_m=0.3$, $\Omega_\Lambda=0.7$."
 The NICMOS UDF field lies within the COODS South region of the sky aud las been imaged by alarge number of telescopes at a variety of wavelengtlis., The NICMOS UDF field lies within the GOODS South region of the sky and has been imaged by alarge number of telescopes at a variety of wavelengths.
 For this project four sources of imagine data for this field were considered., For this project four sources of imaging data for this field were considered.
Microlensing planet searches have been directed toward cliscovering planets in orbits whose size is comparable to the size of the Einstein radius. Aj. of the central star.,"Microlensing planet searches have been directed toward discovering planets in orbits whose size is comparable to the size of the Einstein radius, $R_E$, of the central star."
 Here we study (he detectability of planets in much closer orbits., Here we study the detectability of planets in much closer orbits.
 This is necessary. because we now know that many planets are in such closeorbits!.," This is necessary, because we now know that many planets are in such close."
". We demonstrate that ground based survevs for lensing events can detect a wide range of close-orbit planets. including. ""hot Jupiters"" orbiting sun-like stars. aud even Earth-mass planets in the habitable zones of M clwarls."," We demonstrate that ground based surveys for lensing events can detect a wide range of close-orbit planets, including “hot Jupiters” orbiting sun-like stars, and even Earth-mass planets in the habitable zones of M dwarfs."
 We also discuss the role of space nissions., We also discuss the role of space missions.
 We were led (to the study of close-orbit planets by work to determine whether or not planets in the habitable zones of nearby stars could produce detectable lensing signatures Night 2005)., We were led to the study of close-orbit planets by work to determine whether or not planets in the habitable zones of nearby stars could produce detectable lensing signatures Night 2008).
 The standard. planetarv-lensine scenario is most sensitive to planets. wilh, The standard planetary-lensing scenario is most sensitive to planets with
frequency and k is wavenumber.,frequency and $k$ is wavenumber.
" For the slow mode (cy,=ος). this relationship is indicated by the orange dotted line (6=0) in cach plot of Figure 3.."," 	For the slow mode $c_{ \mathrm{ph} } = c_S$ ), this relationship is indicated by the orange dotted line $\theta = 0$ ) in each plot of Figure \ref{wmst}."
" Considering the properties of the four MHD modes discussed above, we can interpret Figure 3 ina straightforward manner (see Table 1))."," Considering the properties of the four MHD modes discussed above, we can interpret Figure \ref{wmst} 	in a straightforward manner (see Table \ref{modes}) )."
" Red triangles (significant oscillations in intensity and line width) can be interpreted as fast sausage mode, and blue crosses (significant oscillations only in Doppler velocity) can be interpreted as torsional Alfvénn mode or fast kink mode."," Red triangles (significant oscillations in intensity and line width) can be interpreted as fast sausage mode, and blue crosses (significant oscillations only in Doppler velocity) can be interpreted as torsional Alfvénn mode or fast kink mode."
" Green diamonds (significant oscillations in intensity and Doppler velocity) correspond to slow mode, red dots (significant oscillations only in intensity) may be interpreted as fast sausage mode when line width oscillations are obscured by turbulent plasma motion, or slow mode when Doppler velocity signals are reduced due to the inclination of the magnetic field to the line of sight."," Green diamonds (significant oscillations in intensity and Doppler velocity) correspond to slow mode, 	and red dots (significant oscillations only in intensity) may be interpreted as 	fast sausage mode when line width oscillations are obscured by turbulent plasma motion, 	or slow mode when Doppler velocity signals are reduced due to the inclination of the magnetic field to the line of sight."
" Finally, black dots represent non-significant oscillations, oscillations of a shorter period than the time cadence, and oscillations of sub-resolution spatial scale."," Finally, black dots represent non-significant oscillations, 	oscillations of a shorter period than the time cadence, and oscillations of sub-resolution spatial scale."
" The above interpretations therefore suggest that fast sausage mode and slow mode were observed at the moss regions, while fewer oscillations were observed at the apexes of loops."," The above interpretations therefore suggest that fast sausage mode and slow mode were observed at the moss regions, while fewer oscillations were observed at the apexes of loops."
" Distinguishing between fast kink mode and torsional Alfvénn mode is possible by analyzing the phase difference at opposite edges of coronal structures, but this requires spatial resolution high enough to resolve the elemental magnetic strands."," 	Distinguishing between fast kink mode and torsional Alfvénn mode is possible by analyzing the phase difference at opposite edges of 	coronal structures, but this requires spatial resolution high enough to resolve the elemental magnetic strands."
 Phase delays between intensity and Doppler velocity were calculated for data points where both intensity and Doppler velocity amplitudes exceeded the sienificance level., Phase delays between intensity and Doppler velocity were calculated for data points where both intensity and Doppler velocity amplitudes exceeded the significance level.
 The results are shown in Figure 4.., The results are shown in Figure \ref{phst}.
 Panel (a) shows frequency in the horizontal axis and phase delay in the vertical axis., Panel (a) shows frequency in the horizontal axis and phase delay in the vertical axis.
 Panel (b) shows the position from the bottom of the slit in the horizontal axis and phase delay in the vertical axis., Panel (b) shows the position from the bottom of the slit in the horizontal axis and phase delay in the vertical axis.
 Both panels suggest the existence of upwardly propagating slow mode waves (phase delay: +77) and standing waves (phase delay: 7/2) with frequencies of 1—4mHz., Both panels suggest the existence of upwardly propagating slow mode waves (phase delay: $\pm \pi$ ) and standing waves (phase delay: $\pm \pi/2$ ) with frequencies of $1-4 \mspace{3mu} \mathrm{mHz}$.
" theory suggests that phase delays of 0 or xr between intensity and Doppler velocity oscillations are indicative of propagating waves, while phase delays of +7/2 suggest the presence"," 	Adiabatic theory suggests that phase delays of $0$ or $\pm \pi$ between intensity and Doppler velocity oscillations are indicative of propagating waves, 	while phase delays of $\pm \pi/2$ suggest the presence"
Our model solutions depend upon several physical parameters as. for instance. the jet and ambient temperature and density. the jet velocity and its radius.,"Our model solutions depend upon several physical parameters as, for instance, the jet and ambient temperature and density, the jet velocity and its radius."
 With the aim to reduce the number of free parameters m our exploration of the parameter space. we fixed the jet radius to r;=30 AU. according to ? who found this characteristic linear scale. from the X-ray thermal fit. and according to?) who showed HST images of the internal knots of HH 154 with dimension ¢z30 AU at the base of thejet'.," With the aim to reduce the number of free parameters in our exploration of the parameter space, we fixed the jet radius to $r_{j} \approx 30$ AU, according to \citet{ffm02} who found this characteristic linear scale, from the X-ray thermal fit, and according to \citet{fld05} who showed HST images of the internal knots of HH 154 with dimension $r\approx30$ AU at the base of the."
. However detailed simulations with different rj; values are not necessary since we can predict the effects of varying the jet radius from the model results we obtained so far., However detailed simulations with different $r_{j}$ values are not necessary since we can predict the effects of varying the jet radius from the model results we obtained so far.
 In fact we expect the X-ray emitting region to grow in size às rj grows., In fact we expect the X-ray emitting region to grow in size as $r_{j}$ grows.
 Since the X-ray luminosity is defined as Lx=nVPT). it depends on the cube of the radius.," Since the X-ray luminosity is defined as $L_{\rm X} = n^{2} V P(T)$, it depends on the cube of the radius."
 This means that. Lx being constrained from observations. a Jet with a greater radius needs a lower density in order to reproduce observations.," This means that, $L_{\rm X}$ being constrained from observations, a jet with a greater radius needs a lower density in order to reproduce observations."
 We impose ar initial jet length :;=300 AU to avoid the ejected plasma that travels back inside the boundary during the jet evolution., We impose an initial jet length $z_{j} = 300$ AU to avoid the ejected plasma that travels back inside the boundary during the jet evolution.
 This choice of a non-zero initial jet length allows us to obtain an unperturbed boundary surface at >=0., This choice of a non-zero initial jet length allows us to obtain an unperturbed boundary surface at $z = 0$.
 Inall our simulations. we model a jet with initial density and temperature οἱ=500 em * and Ti-10! K respectively. according to the values derived from observations (?. and ?)).," In all our simulations, we model a jet with initial density and temperature $n_{\rm j} = 500$ cm $^{-3}$ and $T_{\rm j} = 10^{4}$ K respectively, according to the values derived from observations \citealt{fl98} and \citealt{ffm02}) )."
" The density and temperature of the ambient medium. 74 and T, respectively. are derived from the choice of the ambient-to-jet density contrast. 7 and from the hypothesis of initial pressure balance between the ambient medium and the jet."," The density and temperature of the ambient medium, $n_{\rm a}$ and $T_{\rm
 a}$ respectively, are derived from the choice of the ambient-to-jet density contrast, $\nu$ and from the hypothesis of initial pressure balance between the ambient medium and the jet."
 We are left. therefore. with two non-dimensional control parameters: the jet Mach number. 17. and the ambient-to-jet density contrast. v.," We are left, therefore, with two non-dimensional control parameters: the jet Mach number, $M$, and the ambient-to-jet density contrast, $\nu$."
 For a more extended exploration of the parameter space. see Sect. 3.5.. ," For a more extended exploration of the parameter space, see Sect. \ref{Varying the jet density parameter},"
concerning the variation of the initial jet density. Hy ," concerning the variation of the initial jet density, $n_{\rm j}$."
In our simulations we account for à wide jet/ambient parameters range shown in Tab. 2.., In our simulations we account for a wide jet/ambient parameters range shown in Tab. \ref{tab:range}.
 In Sect. 3..," In Sect. \ref{Results},"
 we discuss the results derived from. the exploration of the parameters space defined by 1/ and v., we discuss the results derived from the exploration of the parameters space defined by $M$ and $\nu$.
" We performed a wide exploration of the parameter space defined by two free parameters: the jet Mach number. A/= θα. and the ambient-to-jet density ratio. ν=(4/7, (see Sect. 2.3))."," We performed a wide exploration of the parameter space defined by two free parameters: the jet Mach number, $M = v_{\rm j}/c_{\rm a}$ , and the ambient-to-jet density ratio, $\nu = n_{\rm a}/n_{\rm j}$ (see Sect. \ref{Parameters}) )."
 The aim is to determine the range of parameters leading to X-ray emission from protostellar jets in agreement with the observations., The aim is to determine the range of parameters leading to X-ray emission from protostellar jets in agreement with the observations.
 We first analyzed adiabatic hydrodynamic models. ie. without thermal conduction and radiative losses.," We first analyzed adiabatic hydrodynamic models, i.e. without thermal conduction and radiative losses."
 Then. for the most promising cases (Le. those which reproduce the values of jet velocity. temperature and luminosity of the X-ray source derived from the observations). we performed more realistic simulations in. which we have taken into account thermal conduction and radiative losses effects.," Then, for the most promising cases (i.e. those which reproduce the values of jet velocity, temperature and luminosity of the X-ray source derived from the observations), we performed more realistic simulations in which we have taken into account thermal conduction and radiative losses effects."
 By comparing these models with those without thermal conduction. anc radiative cooling. we explored how the presence of these physical processes affects the jet/ambient system evolution.," By comparing these models with those without thermal conduction and radiative cooling, we explored how the presence of these physical processes affects the jet/ambient system evolution."
 We found that. in general. models with thermal conductiot and radiation reach lower temperatures (up to 5 times lower than those achieved in the adiabatic cases).," We found that, in general, models with thermal conduction and radiation reach lower temperatures (up to $5$ times lower than those achieved in the adiabatic cases)."
 We also found that thermal conduction smooths the structures (well visible i the pure hydrodynamic cases) in the density and temperature distributions., We also found that thermal conduction smooths the structures (well visible in the pure hydrodynamic cases) in the density and temperature distributions.
 In the following subsections. we discuss the models (shown in Fig. 2) ," In the following subsections, we discuss the models (shown in Fig. \ref{fig:griglia}) )"
in which both radiative losses and thermal conduction are taken into account., in which both radiative losses and thermal conduction are taken into account.
 In Fig., In Fig.
 2 green and red dots refer to those cases with X-ray luminosity Ly1077 erg J|. shock front velocity e;>100 km ! and fitting temperature T<10* K. consistent with observations.," \ref{fig:griglia} green and red dots refer to those cases with X-ray luminosity $L_{X} > 10^{28}$ erg $^{-1}$, shock front velocity $v_{sh} > 100$ km $^{-1}$ and fitting temperature $T < 10^{7}$ K, consistent with observations."
 We have chosen Lx one order of magnitude lower than the minimum value observed (see Tab. 1)), We have chosen $L_{X}$ one order of magnitude lower than the minimum value observed (see Tab. \ref{tab:obs}) )
 to take into account fainter sources ot detected so far; the red dot refers to the representative case of 1154 discussed in ?.., to take into account fainter sources not detected so far; the red dot refers to the representative case of 154 discussed in \citet{bop04}.
" Squares show cases with velocity in the range of values observed. but with Ly<1075 ergs ο,"," Squares show cases with velocity in the range of values observed, but with $L_{X} < 10^{28}$ erg $^{-1}$."
 Diamonds mark the cases with velocity and X-ray luminosity ot consistent with observations., Diamonds mark the cases with velocity and X-ray luminosity not consistent with observations.
 Triangles mark cases with temperatures higher than 10° K. The lower panel of Fig., Triangles mark cases with temperatures higher than $10^{7}$ K. The lower panel of Fig.
" 2 show the initial velocity assumed in our simulations vs. the ""Sensity contrast.", \ref{fig:griglia} show the initial velocity assumed in our simulations vs. the density contrast.
 From our exploration of the parameter space. we derived that the models in agreement with observations are included in a well constrained region.," From our exploration of the parameter space, we derived that the models in agreement with observations are included in a well constrained region."
 In the following sections. we discuss in details the “best-fit” models. t. e. those models which osreproduce X-ray luminosity and shock front speed values as lose as possible to those observed. in the cases of light. ensity and heavy jets (see Tab. 3)).," In the following sections, we discuss in details the “best-fit” models, i. e. those models which reproduce X-ray luminosity and shock front speed values as close as possible to those observed, in the cases of light, equal-density and heavy jets (see Tab. \ref{tab_mod}) )."
 In Fig. 3..," In Fig. \ref{fig:mappe-20yr},"
 we show the mass density and temperature distributions 20 years since the beginning of the jet/ambient medium interaction for the three best-fit models in Tab. 3.., we show the mass density and temperature distributions $20$ years since the beginning of the jet/ambient medium interaction for the three best-fit models in Tab. \ref{tab_mod}.
 The light jet case is the one which better reproduce the physical parameters derived from observations by ? and ? for the 1154 protostellar Jet: its properties have been discussed in 9, The light jet case is the one which better reproduce the physical parameters derived from observations by \citet{fl98} and \citet{ffm02} for the 154 protostellar jet; its properties have been discussed in \citet{bop04}.
 In all the cases. at the head of the Jet there is clear evidence of a shock front due to the plasma propagating supersonically along the jet axis.," In all the cases, at the head of the jet there is clear evidence of a shock front due to the plasma propagating supersonically along the jet axis."
 Just behind the shock front there Is a localized hot and dense blob (see. for instance. the enlargement in Fig.," Just behind the shock front there is a localized hot and dense blob (see, for instance, the enlargement in Fig."
 4+ for the light jet case).," \ref{fig:zoom-light}
 for the light jet case)."
 The light jet is enveloped by a cocoon with temperature T=7<10° K. almost uniform due to the thermal conduction diffusive effect: nevertheless the cocoon temperature is not constant in time but decreases as the evolution goes on. leading to the formation of a cool and dense external envelope.," The light jet is enveloped by a cocoon with temperature $T \approx
7\times10^{5}$ K, almost uniform due to the thermal conduction diffusive effect; nevertheless the cocoon temperature is not constant in time but decreases as the evolution goes on, leading to the formation of a cool and dense external envelope."
 Fig., Fig.
statistical uncertainty radii).,statistical uncertainty radii).
 The limiting R-band magnitude for the selected optical objects. for example. is 17.5.," The limiting R-band magnitude for the selected optical objects, for example, is 17.5."
 From the residual oplical/N-rav position offsets of these associations. we inler that the positional inaceuracy after (he correction. has been made should be less than ~073.," From the residual optical/X-ray position offsets of these associations, we infer that the positional inaccuracy after the correction has been made should be less than $\sim 0\farcs3$."
 The location of the Quintuplet cluster is also indicated in 11., The location of the Quintuplet cluster is also indicated in 1.
 This GC cluster contains numerous massive stars including the huminous Pistol Star (Figer οἱ al., This GC cluster contains numerous massive stars including the luminous Pistol Star (Figer et al.
 1999)., 1999).
 However. we detect only one possible X-ray counterpart of a cluster member.," However, we detect only one possible X-ray counterpart of a cluster member."
 We [ind no counterpart for the bright N-ray source just north of GO.13-0.11. which is not considered," We find no counterpart for the bright X-ray source just north of G0.13-0.11, which is not considered"
the past. which could have been the case due to core collapse (Noyola Gebhardt 2006. de Marchi et al.,"the past, which could have been the case due to core collapse (Noyola Gebhardt 2006, de Marchi et al."
 2007) or tidal stripping (e.g. Lee et al., 2007) or tidal stripping (e.g. Lee et al.
 2007)., 2007).
 To properly perform this comparison. deeper UV imaging data will be required that allow detection of UV intermediate-bright to faint GCs down to My=~ —10mag (V=21.5 mag at the Fornax distance).," To properly perform this comparison, deeper UV imaging data will be required that allow detection of UV intermediate-bright to faint GCs down to $M_V\simeq-$ 10mag $V\simeq 21.5$ mag at the Fornax distance)."
 In this respect. the outcomes of the HST observations in Cycle 15 (GO10901. PI O'Connell) of GCs belonging to NGC 1399 are highly anticipated.," In this respect, the outcomes of the HST observations in Cycle 15 (GO10901, PI O'Connell) of GCs belonging to NGC 1399 are highly anticipated."
 We are grateful to Caleb Scharf for providing us with the source catalog of the Chandra Fornax survey., We are grateful to Caleb Scharf for providing us with the source catalog of the Chandra Fornax survey.
 The work of S.-C. R. was supported in part by KOSEF through the Astrophysical Research Center for the Structure and Evolution of the Cosmos (ARCSEC)., The work of S.-C. R. was supported in part by KOSEF through the Astrophysical Research Center for the Structure and Evolution of the Cosmos (ARCSEC).
Schneider et al. (,Schneider et al. (
2005) have presented a scenario for the ormation of galaxies in a concordance ACDAL cosmological which includes a self-consistcnt treatment of two kev feedback: processes: (i)feedback. suppressing star formation in Lls-cooling halos and the formation of ow-mass galaxies due to the cllects of UM. background racliation approaching the reionization epoch. and (ii)feedback. which controls the transition from meta stars (Poplll) to ordinary stars (PoX) through the orogressive enrichment of star forming eas with heavy elements released by supernova explosions (Schneider et al.,"2005) have presented a scenario for the formation of galaxies in a concordance $\Lambda$ CDM cosmological which includes a self-consistent treatment of two key feedback processes: (i), suppressing star formation in $_2$ -cooling halos and the formation of low-mass galaxies due to the effects of UV background radiation approaching the reionization epoch, and (ii), which controls the transition from metal-free stars (PopIII) to ordinary stars (PopII) through the progressive enrichment of star forming gas with heavy elements released by supernova explosions (Schneider et al."
 2002. 2004. 2005: Bromm ct al.," 2002, 2004, 2005; Bromm et al."
 2001)., 2001).
 €rwermical feedback o»opagateso throughe the hierarchy. of agalaxy mergers.oO [rom orogenitors to their descendants so that. at each redshift. existing halos which are allowed to form stars are classified as Popll (Poplll) galaxies depending on whether the halo itself or any of its progenitors have (have not) alreaciy experienced an episode of star formation.," Chemical feedback propagates through the hierarchy of galaxy mergers from progenitors to their descendants so that, at each redshift, existing halos which are allowed to form stars are classified as PopII (PopIII) galaxies depending on whether the halo itself or any of its progenitors have (have not) already experienced an episode of star formation."
" Within this model we can compute the comoving specific emissivity. co. which is given by where τηΑα: is the formation rate of halos οἱ total (clark|barvonic) mass Ad, with corresponding stellar mass Ad. and Ad,,i,(2) is the minimum mass of halos (corresponding to a virial temperature of 101 Ko ie. Abin(sz)~WOAL.(OL|2/10) 27) allowed to form stars ab redshift 2: Adj is the maximum mass of halos which depends on the details of their merging history. (sce Schneider οἱ al."," Within this model we can compute the comoving specific emissivity, $\epsilon_\nu$, which is given by where $d^2n/dM_h dz$ is the formation rate of halos of total (dark+baryonic) mass $M_h$ with corresponding stellar mass $M_\star$, and $M_{min}(z)$ is the minimum mass of halos (corresponding to a virial temperature of $10^4$ K, i.e. $M_{min}(z)\sim 10^8\;\Msun\;(1+z/10)^{-3/2}$ ) allowed to form stars at redshift $z$; $M_{max}(z)$ is the maximum mass of halos which depends on the details of their merging history (see Schneider et al."
 2005): fC.κ)) is the template specific luminosity for a stellar population of age £. (time elapsed between redshift z and z).," 2005); $l_\nu(t_{z,z^\prime})$ is the template specific luminosity for a stellar population of age $t_{z,z^\prime}$ (time elapsed between redshift $z^\prime$ and $z$ )."
 Following the results of Schneider et al., Following the results of Schneider et al.
" 2005. both Popll and Poplll stars are assumed to form. according to a Salpeter Initial Mass Function (ALE) with an cllicieney f,= 0.1."," 2005, both PopII and PopIII stars are assumed to form according to a Salpeter Initial Mass Function (IMF) with an efficiency $f_\star = 0.1$ ."
 The emission properties of Popll stars are taken from the GALANEY library (masses in the range 0.1-100. AZ... metallicity Z=10?Z.: Muzual Charlot 2003) and those of Poplll stars are based on the synthesis model for metal-free stars of Schaerer (1-100 AJ... Z=0: Schacrer 2002).," The emission properties of PopII stars are taken from the GALAXEV library (masses in the range 0.1-100 $\msun$, metallicity $Z=10^{-2}Z_\odot$; Bruzual Charlot 2003) and those of PopIII stars are based on the synthesis model for metal-free stars of Schaerer (1-100 $\msun$, $Z=0$; Schaerer 2002)."
 A set of independent. observational constraints can be accommodated: within this model (Schneider et al., A set of independent observational constraints can be accommodated within this model (Schneider et al.
 2005): he number counts of dropout galaxies at 2o6 (i-dropouts: Douwens et al., 2005): the number counts of dropout galaxies at $z\sim 6$ $i$ -dropouts; Bouwens et al.
 2005b) and at z~10 (J-dropouts: Bouwens et al., 2005b) and at $z\sim 10$ $J$ -dropouts; Bouwens et al.
" 2005a). and the value of the optical depth for Thomson scattering 7,.=0.16-E0.04 measured by the WALAD satellite (Ixogut et al."," 2005a), and the value of the optical depth for Thomson scattering $\tau_e=0.16\pm0.04$ measured by the WMAP satellite (Kogut et al."
 2003)., 2003).
 Indeed. the predicted surface density of i-dropouts at a limiting magnitude of7=28 is ~5 arcmin2 (observed value ~4.7 7: Bouwens et al.," Indeed, the predicted surface density of $i$ -dropouts at a limiting magnitude of $i=28$ is $\sim5$ $^{-2}$ (observed value $\sim 4.7$ $^{-2}$; Bouwens et al."
 2005b)., 2005b).
" ""he model is also consistent with the three candidate dropouts ab zcLO reported by Bouwens et al. (", The model is also consistent with the three candidate dropouts at $z\sim 10$ reported by Bouwens et al. (
2005a).,2005a).
" Moreover. the large value of 7. can be matched with standard values of the ionizingSl photon escape1 fraction [rom &galaxies. (ος,»<0.2) and eas clumping factor Co10. resulting in a reionization redshift of 13.2 (Schneider et al."," Moreover, the large value of $\tau_e$ can be matched with standard values of the ionizing photon escape fraction from galaxies $f_{esc} \leq 0.2$ ) and gas clumping factor $C \sim 10$, resulting in a reionization redshift of 13.2 (Schneider et al."
 2005)., 2005).
 The model presented in Section 2 (referred to as model C in Schneider et al., The model presented in Section 2 (referred to as model C in Schneider et al.
" 2005) allows to compute the NIIUB intensityJi, from z=5 galaxies seen at frequency vy by an observer at redshift, z,=0.", 2005) allows to compute the NIRB intensity$J_{\nu_{0}}$ from $z\ge 5$ galaxies seen at frequency $\nu_{0}$ by an observer at redshift $z_{0}=0$.
 This can be done in a fashion similar to Salvaterra Ferrara (2003) as: where v=voll|z)/(Gzo). dlfdz is he proper line element. rqrit.zo.2) is the elfective optical depth at f of the intergalactic medium between redshift zi and z (see Section 2.2 of Salvaterra Ferrara 2003 for a full description of the IGM modelling). and ος} is provided. by equation (1).," This can be done in a fashion similar to Salvaterra Ferrara (2003) as: where $\nu=\nu_0(1+z)/(1+z_0)$, $dl/dz$ is the proper line element, $\tau_{eff}(\nu_{0},z_{0},z)$ is the effective optical depth at $\nu_0$ of the intergalactic medium between redshift $z_0$ and $z$ (see Section 2.2 of Salvaterra Ferrara 2003 for a full description of the IGM modelling), and $\epsilon_\nu(z)$ is provided by equation (1)."
 Our results are shown in Figure 1., Our results are shown in Figure 1.
" The background intensity.: J),d−2 nW⇁ 2 losC. ds almost constant in the Q.8-8 jim range."," The background intensity, $\nu J_\nu \sim 1-2$ nW $^{-2}$ $^{-1}$, is almost constant in the 0.8-8 $\mu$ m range."
 Popll galaxies dominate the NIIUD in the entire wavelength range. while Poplll galaxies contribute at most of the total intensity (at Ac1.5 yam). via their strong Lya line emission.," PopII galaxies dominate the NIRB in the entire wavelength range, while PopIII galaxies contribute at most of the total intensity (at $\lambda\sim 1.5\;\mu$ m), via their strong $\alpha$ line emission."
 We have checked that the sources responsible for such emission. are too faint to be resolved ancl identified at the Spitzer magnitude limit in all four channels., We have checked that the sources responsible for such emission are too faint to be resolved and identified at the Spitzer magnitude limit in all four channels.
 The total contribution of all galaxies at zc5 is 1.55. 1.45. 1.07 and 0.74 nW m> in the 3.6. 4.5. 5.8 ancl 8.0 LRAC bands. respectively. representing a substantial paction ( YA) of the otal NIRB intensity estimated rom integration of Spitzer galaxy number counts (Fazio et al.," The total contribution of all galaxies at $z\ge 5$ is 1.55, 1.45, 1.07 and 0.74 nW $^{-2}$ $^{-1}$ in the 3.6, 4.5, 5.8 and 8.0 IRAC bands, respectively, representing a substantial fraction $\sim$ ) of the total NIRB intensity estimated from integration of Spitzer galaxy number counts (Fazio et al."
 2004b)., 2004b).
 ‘The angular correlation function ol intensity luctuations DNI due. to. inhomogencitics in the space clistribution of unresolved: sources (i.e. with Uuxes fainter han some thresyold S4) is ¢clined as: where (6'o) and (67LO) HaskelentilyDp (wo positionsMN on the sky separated. by an angle 6.," The angular correlation function of intensity fluctuations $\delta J$ due to inhomogeneities in the space distribution of unresolved sources (i.e. with fluxes fainter than some threshold $S_d$ ) is defined as: where $(\theta^\prime, \phi^\prime)$ and $(\theta'', \phi'')$ identify two positions on the sky separated by an angle $\theta$."
 The above expression can be written as the sum of two terms. Cp and Ce. the first one due to Poisson noise (Lo. Huctuations given. by randomly distributed. objects). and the second one owing to source clustering.," The above expression can be written as the sum of two terms, $C_P$ and $C_C$, the first one due to Poisson noise (i.e. fluctuations given by randomly distributed objects), and the second one owing to source clustering."
 Lt can be shown that the shot noisecontribution originating [ron z5 galaxies is negligible. so we only concentrate on Lluctuations which stem from the clustering of these sources. Le. we assume C'(8)= Cee.," It can be shown that the shot noisecontribution originating from $z\geq 5$ galaxies is negligible, so we only concentrate on fluctuations which stem from the clustering of these sources, i.e. we assume $C(\theta)\equiv C_C$ ."
 The method adopted here is similar to that. presented in Magliocchetti et al. (, The method adopted here is similar to that presented in Magliocchetti et al. (
2003). whereby angular Iuctuations are obtained by means of the expression:,"2003), whereby angular fluctuations are obtained by means of the expression:"
telescope to cover the entire far-infrared waveband (from 55 to ym) and looks likely to become one of the greatest astronomical achievements of this decade.,telescope to cover the entire far-infrared waveband (from 55 to $\mu$ m) and looks likely to become one of the greatest astronomical achievements of this decade.
 The Astrophysical Terahertz Large Area Survey (H-ATLAS—?) is the largest Open Time Key Project hhr). covering ddeg? of sky in regions selected on the basis of existing multi-wavelength coverage tthe Galaxy Evolution Explorer - GALEX. the Galaxy and Mass Assembly spectroscopic survey — GAMA. the 2dF Galaxy Redshift Survey — 2DFGRS. the Sloan Digital Sky Survey — SDSS. and the Dark Energy Survey - DES).," The Astrophysical Terahertz Large Area Survey \citep[\hatlas\ --][]{eales10} is the largest Open Time Key Project hr), covering $^2$ of sky in regions selected on the basis of existing multi-wavelength coverage the Galaxy Evolution Explorer – GALEX, the Galaxy and Mass Assembly spectroscopic survey – GAMA, the 2dF Galaxy Redshift Survey – 2DFGRS, the Sloan Digital Sky Survey – SDSS, and the Dark Energy Survey – DES)."
 wwill detect hundreds of thousands of galaxies (2). and provide an extensive census of dust-obscured activity in the local (2<0.3) Universe (2)).," will detect hundreds of thousands of galaxies \citealt{clements10}) ), and provide an extensive census of dust-obscured activity in the local $z<0.3$ ) Universe \citealt{amblard10}) )."
 The areal coverage of aalso makes it well suited to the identification of sources citealtgonzalez-nuevo10)). lensed galaxies at high redshift citealtnegrelloO7.| swinbank10)) and. local dust. clouds at high Galactic latitudes: there is also enormous potential —for serendipitous discovery.," The areal coverage of also makes it well suited to the identification of sources \\citealt{gonzalez-nuevo10}) ), lensed galaxies at high redshift \\citealt{negrello07, swinbank10}) ) and local dust clouds at high Galactic latitudes; there is also enormous potential for serendipitous discovery."
 iis exploiting the fast-scan ). parallel mode. ofHerschel. using two of the on-board instruments to. provide an efficient way to map large areas of sky in five wavebands simultaneously.," is exploiting the fast-scan $^{-1}$ ) parallel mode of, using two of the on-board instruments to provide an efficient way to map large areas of sky in five wavebands simultaneously."
" We are using the Photodetector Array Camera and Spectrometer (PACS—?) to observe at 100 and jim (its ""green and ‘red’ channels) whilst the Spectral and Photometric Imaging Receiver (SPIRE—?) is taking data at 250. 350 and 500 jim. This paper is one of a series of four technical papers describing our approach to the PACS (this paper) and SPIRE (Pascale et al.."," We are using the Photodetector Array Camera and Spectrometer \citep[PACS --][]{poglitsch10} to observe at 100 and $\mu$ m (its `green' and `red' channels) whilst the Spectral and Photometric Imaging Receiver \citep[SPIRE --][]{griffin10} is taking data at 250, 350 and $\,\mu$ m. This paper is one of a series of four technical papers describing our approach to the PACS (this paper) and SPIRE (Pascale et al.,"
 in preparation) data products. to source extraction (Rigby et al..," in preparation) data products, to source extraction (Rigby et al.,"
 in preparation) and to cross-identification (Smith et al..," in preparation) and to cross-identification (Smith et al.,"
 in preparation) for the Science Demonstration Phase (SDP) of the ssurvey., in preparation) for the Science Demonstration Phase (SDP) of the survey.
 These data are public and available atazlas., These data are public and available at.
orgq/. Here. we describess the pipeline— used to reduce data obtained with PACS. and the quality of data products. to give an idea of its scientific potential.," Here, we describe the pipeline used to reduce data obtained with PACS, and the quality of data products, to give an idea of its scientific potential."
 In refSinstrument.. we provide a brief description of the instrument: in refSthedata.. we present the SSDP observations: in refSpipeline.. we deseribe the customised procedures we have developed within theHersche! Interactive Processing Environment — ?)) to reduce data from PACS.," In \\ref{Sinstrument}, , we provide a brief description of the instrument; in \\ref{Sthedata}, we present the SDP observations; in \\ref{Spipeline}, we describe the customised procedures we have developed within the Interactive Processing Environment – \citealt{ott10}) ) to reduce data from PACS."
 In refSanalyses we describe tests of the resulting images and we state sme concluding remarks in refSconclusions.., In \\ref{Sanalyses} we describe tests of the resulting images and we state some concluding remarks in \\ref{Sconclusions}.
 PACS is a multi-colour camera and low- and medium-resolution spectrometer covering the S55—210-;7m wavelength range (see Fig. 19., PACS is a multi-colour camera and low- and medium-resolution spectrometer covering the $\mu$ m wavelength range (see Fig. \ref{filters}) ).
 It comprises four large-format detector arrays: two filled silicon bolometer arrays optimised for imaging in high-photon-background conditions and two Ge:Ga photo-conductor arrays for spectroscopy., It comprises four large-format detector arrays: two filled silicon bolometer arrays optimised for imaging in high-photon-background conditions and two Ge:Ga photo-conductor arrays for spectroscopy.
 Here we concentrate on the bolometer detectors used by the ssurvey — a more complete description of the instrument and its modes can be found at ?.., Here we concentrate on the bolometer detectors used by the survey – a more complete description of the instrument and its modes can be found at \citet{poglitsch10}.
 A dichroic beam splitter enables shotometry in two bands simultaneously — 70 or j/m Cblue? or ‘green’. selected by a filter wheel) and j/m Cred”) — over he same 3.5-aremin?. field of view.," A dichroic beam splitter enables photometry in two bands simultaneously – 70 or $\mu$ m (`blue' or `green', selected by a filter wheel) and $\mu$ m (`red') – over the same $\times$ $^2$ field of view."
 The bolometer arrays comprise 64.32 and 16 pixels. with 3.2 and aarcsec + on-sky. respectively. providing close-to Nyquist: beam sampling for the blue/green and red filters.," The bolometer arrays comprise $\times$ 32 and $\times$ 16 pixels, with 3.2 and arcsec $^{-1}$ on-sky, respectively, providing close-to Nyquist beam sampling for the blue/green and red filters."
 The arrays each comprise sub-arravs of 16. [6 pixels. tiled together to form the Όσα plane (see ? and references therein).," The arrays each comprise sub-arrays of $\times$ 16 pixels, tiled together to form the focal plane (see \citealt{billot09} and references therein)."
 Working in ‘scan mode’. PACS modulates the sky signal by making use of the motion of the spacecraft (10. 20 or sslj with no chopping.," Working in `scan mode', PACS modulates the sky signal by making use of the motion of the spacecraft (10, 20 or $^{-1}$ ), with no chopping."
 The sky signal is stored in units of mV by the analogue-to-digital (ADU) converter. depending on the user-detined V/ADU gain value ¢high or low: see 23).," The sky signal is stored in units of mV by the analogue-to-digital (ADU) converter, depending on the user-defined V/ADU gain value (high or low; see \citealt{poglitsch10}) )."
 The signal from each bolometer pixel is sampled at a rate of HHz. although due to satellite telemetry limitations. in scan mode the signal is averaged into packages of four consecutive frames — rresulting in an effective rate of HHz.," The signal from each bolometer pixel is sampled at a rate of Hz, although due to satellite telemetry limitations, in scan mode the signal is averaged into packages of four consecutive frames – resulting in an effective rate of Hz."
 Data is also bit-rounded by the Signal-Processing Unit (SPU) (25) which results in a stronger quantisation of the signal than would be expected by the ADU converter., Data is also bit-rounded by the Signal-Processing Unit (SPU) \citealt{ottensamer08}) ) which results in a stronger quantisation of the signal than would be expected by the ADU converter.
 When using the SPIRE/PACS ‘fast-scan parallel mode’. as employed byH-ATLAS.. in particular for the blue/green filters. data suffer additional in-flight averaging (eight frames). resulting in an effective read-out frequency of HHz.," When using the SPIRE/PACS `fast-scan parallel mode', as employed by, in particular for the blue/green filters, data suffer additional in-flight averaging (eight frames), resulting in an effective read-out frequency of Hz."
 Due to the limited signal bandwidth of the detection chain. the on-board data compression produces significant degradation of the observed point spread function (PSF: Instrument Control Centre - [CC report».," Due to the limited signal bandwidth of the detection chain, the on-board data compression produces significant degradation of the observed point spread function (PSF; Instrument Control Centre – ICC )."
 Simulated parallel mode data based on PACS prime fast-scan observations show that point-source peak fluxes are reduced by 750 and 70 per cent. at 100 and j/m respectively. in comparison to nominal 1) sean (29).," Simulated parallel mode data based on PACS prime fast-scan observations show that point-source peak fluxes are reduced by $\sim$ 50 and 70 per cent, at 100 and $\mu$ m respectively, in comparison to nominal $^{-1}$ ) scan \citealt{poglitsch10}) )."
 Prior to assessment of the in-flight performance. the predicted 5-7 point-source ssensitivities based on the Herschel-Spot (ASpot’)) observation planning tool were 67 and 94+mmJy for observations with one pair of eross-seans. at [O0 and /m. respectively (23).," Prior to assessment of the in-flight performance, the predicted $\sigma$ point-source sensitivities based on the -Spot ) observation planning tool were 67 and mJy for observations with one pair of cross-scans, at 100 and $\mu$ m, respectively \citealt{eales10}) )."
 The performance is mostly dependent on the optical efficiency. the thermal and telescope background. the effects of Cosmic-Ray glitches — in particular high-energy protons — and the photon noise from the detector and multiplexer electronics which was found to introduce a 1/f excess below HHz prior to launch (see refSprojection)).," The performance is mostly dependent on the optical efficiency, the thermal and telescope background, the effects of Cosmic-Ray glitches – in particular high-energy protons – and the photon noise from the detector and multiplexer electronics which was found to introduce a $1/f$ excess below Hz prior to launch (see \\ref{Sprojection}) )."
 The PACS focal plane is offset with respect to SPIRE by a fixed separation of ~2] aaremin along the spacecraft z-axis. implying different instantaneous PACS and SPIRE coverages.," The PACS focal plane is offset with respect to SPIRE by a fixed separation of $\sim$ arcmin along the spacecraft $z$ -axis, implying different instantaneous PACS and SPIRE coverages."
 Thecoverage obtained by SPIRE in the SDP area is presented in, Thecoverage obtained by SPIRE in the SDP area is presented in
Tt is known that cosmic ravs arrival direction has a- energv dependent large angular scale anisotropy with a- amplitude of order 10110%.,It is known that cosmic rays arrival direction has an energy dependent large angular scale anisotropy with an amplitude of order $10^{-4}-10^{-3}$.
 The first comprehensive observation of this anisotropy was provided by a network of immon telescopes scusitive to συΤον cucreies aud located at different latitudes (Nagashimactal., The first comprehensive observation of this anisotropy was provided by a network of muon telescopes sensitive to sub-TeV energies and located at different latitudes \citep{nagashima}.
1998).. More recently an anisotropy was also observed in the inultiTeV energy range bv the Tibet AS* array (Amenomorietal.2006)... Super-Iiuniokaude (Ciuilliauetal.2007) aud by MILACBRO (Abdoetal.2009).. and the first hieh statistics observation im the southern hemisphere in the 10 TeV region. is being reported bv TeceCube (Abbasictal.2010).," More recently, an anisotropy was also observed in the multi-TeV energy range by the Tibet $\gamma$ array \citep{amenomori}, Super-Kamiokande \citep{guillian} and by MILAGRO \citep{abdo}, and the first high statistics observation in the southern hemisphere in the 10 TeV region, is being reported by IceCube \citep{abbasi}."
. The origin of the large augular scale anisotropy iu the cosmic ravs arrival direction is still uuknown., The origin of the large angular scale anisotropy in the cosmic rays arrival direction is still unknown.
 The structure of the local iuterstellar inagnetic field is likely to have an important role., The structure of the local interstellar magnetic field is likely to have an important role.
 However the combined study of the anisotropy energv and anenlar cepeudency. its time modulation and angular scale structure secur to sugecst that the observation imieht be a combination of multiple superimposed effects. caused by plenomenologics at dciffereut distances from Earth.," However the combined study of the anisotropy energy and angular dependency, its time modulation and angular scale structure seem to suggest that the observation might be a combination of multiple superimposed effects, caused by phenomenologies at different distances from Earth."
 Iu this context. particular interest is derived from the observation of a broad excess of sub-TeV cosmic ravs in a portion of the sky compatible with the direction of the heliospherie tail (or heliotzil) (Nagashimactal.1995:Talletal.1999). (see 82)).," In this context, particular interest is derived from the observation of a broad excess of sub-TeV cosmic rays in a portion of the sky compatible with the direction of the heliospheric tail (or heliotail) \citep{nagashima,hall} (see \ref{sec:obs}) )."
 The heliotzil is the reeion of the heliosphere dowustream the imterstellar uatter wind dehnüted within the heliopause. ie. the )oundary that separates the solar wind aud iuterstellar udasimas (Izmodenov&Wallenbach2006).," The heliotail is the region of the heliosphere downstream the interstellar matter wind delimited within the heliopause, i.e. the boundary that separates the solar wind and interstellar plasmas \citep{izmodenov}."
. The observed excess Was attributed to some unknown anisotropic Xxocess connected with the heliotail (thus. called. tail-in excess)., The observed excess was attributed to some unknown anisotropic process connected with the heliotail (thus called tail-in excess).
 The evro-radius of sub-TeV cosmidc protons is less than about 200 AU (in a —1 pC interstellar naenetic field). which is approximately the size of the wchosplere aud. most likely. smaller than the width aud leugth of the heliotail.," The gyro-radius of sub-TeV cosmic protons is less than about 200 AU (in a $\sim$ 1 $\mu$ G interstellar magnetic field), which is approximately the size of the heliosphere and, most likely, smaller than the width and length of the heliotail."
 The persistence of the cosmic rav alisotropy structure in the multi-TeV energy rauge makes it challengiug to link this observation to the heliosphere., The persistence of the cosmic ray anisotropy structure in the multi-TeV energy range makes it challenging to link this observation to the heliosphere.
 Although the unknown size and extension of the hehotail contributes to the uucertaiutv on the enerev scale at which heliospheric mfiueuce on cosmic ravs starts to be negligible., Although the unknown size and extension of the heliotail contributes to the uncertainty on the energy scale at which heliospheric influence on cosmic rays starts to be negligible.
 However. we kuow that the observations of multiΤον cosuic rays auisotropy show sinall angular scale patterus superimposed to the sincoth broad structure of the taibiu excess. which is sugeestive of a local origin. ic. within the leliotail.," However, we know that the observations of multi-TeV cosmic rays anisotropy show small angular scale patterns superimposed to the smooth broad structure of the tail-in excess, which is suggestive of a local origin, i.e. within the heliotail."
 With the same technique used im eanuua rav detection to estimate the backeround aud search for sources of eanuna rays. the MILAGRO collaboration discovered two localized excess regions in the cosmüc ravs arrival direction distribution (Abdooetal.2008).," With the same technique used in gamma ray detection to estimate the background and search for sources of gamma rays, the MILAGRO collaboration discovered two localized excess regions in the cosmic rays arrival direction distribution \citep{abdo2}."
. The same excess regions were reported by the ARGO-YDJ air shower array (Vernettoetal.2009)., The same excess regions were reported by the ARGO-YBJ air shower array \citep{vernetto}.
. The strongest and amore localized of them (with au augular size of about 107) coincides with the direction of the heliotail., The strongest and more localized of them (with an angular size of about $^{\circ}$ ) coincides with the direction of the heliotail.
 The pecularity of such an observation triggered au astroplivsical interpretation based ou the possibility that cosuic ravs accelerated by the supernova that produced Ceminga pulsar are focussed by an ad-hoc interstellar magnetic field structure (Salvati&Sacco2008:Aharonian2008) (sec 83)).," The peculiarity of such an observation triggered an astrophysical interpretation based on the possibility that cosmic rays accelerated by the supernova that produced Geminga pulsar are focussed by an ad-hoc interstellar magnetic field structure \citep{salvati,drury} (see \ref{sec:interp}) )."
 The localized regions lie in the same portion of the sky that is dominated by the broad tailin excess at lower cuerey aud. although it might be coincidental. we interpret this as manifestations of the sale phenomenoloey at different cucreics.," The localized regions lie in the same portion of the sky that is dominated by the broad tail-in excess at lower energy, and, although it might be coincidental, we interpret this as manifestations of the same phenomenology at different energies."
 It is proposed that both sub-Te tail-in excess aud the maulti-TeV localized excess of cosinicV. rays might be caused by magnetic reconnection iu the heliosphere aud. in particular in the hehotail. where the distance scale nueght be long enough to induce sufficient acceleration at high enerev.," It is proposed that both sub-TeV tail-in excess and the multi-TeV localized excess of cosmic rays might be caused by magnetic reconnection in the heliosphere and, in particular in the heliotail, where the distance scale might be long enough to induce sufficient acceleration at high energy."
 The very idea of appealing to maguetic reconnection for the acceleration of chergctic particles, The very idea of appealing to magnetic reconnection for the acceleration of energetic particles
We estimate the distaice as follows.,We estimate the distance as follows.
 First. we use secondary stars contribution. in Combination with the orbital period. to estimate a distauce spectrophotometrically.," First, we use secondary star's contribution, in combination with the orbital period, to estimate a distance spectrophotometrically."
 Next. we use a Bayesian ormalisin described by TIorsteusen(2003) to combine this estimate with the parallax aud proper notion.," Next, we use a Bayesian formalism described by \citet{thorparallax} to combine this estimate with the parallax and proper motion."
 This gives a best yossible cdistauce., This gives a best possible distance.
 Beuermanuetal.(1999) tabulate radii and other properties of a variety of late-type stars., \citet{beuermann99} tabulate radii and other properties of a variety of late-type stars.
 From his one can derive the suTace brightuess as a functiou of spectral type. which can be formulated convenieutly as Ady(LAG. ). he absolute V. inaguitudee of a I-solar-radius star of the same surface wightness.," From this one can derive the surface brightness as a function of spectral type, which can be formulated conveniently as $M_V(1 R_{\odot})$ , the absolute $V$ magnitude of a 1-solar-radius star of the same surface brightness."
 Among the ex:uuples given in their table. this quantity varies (roi a mininuun of 5.61 Oa παν of 9.97 in the rauge M3 — ΔΕΣ.," Among the examples given in their table, this quantity varies from a minimum of 8.61 to a maximum of 9.97 in the range M3 – M4.5."
 From this. we adopt Ad)(LR.)=9.30.7 for the 13.75£0.75 secondary of V105 Pee.," From this, we adopt $M_V(1 R_{\odot}) = 9.3 \pm 0.7$ for the $3.75 \pm 0.75$ secondary of V405 Peg."
 Beuermaunetal.(1998). give a convenient formula for the Roche lobe radius as a function of orbital period and secondary mass: at the period of V105 Peg. this vields ο.=0.616/(y)AloΛΙ.ytoE where f(q) is close to unity.," \citet{beuermann98} give a convenient formula for the Roche lobe radius as a function of orbital period and secondary mass; at the period of V405 Peg, this yields $R_2 / R_{\odot} = 0.616 f(q) (M_2 /
M_{\odot})^{1/3}$, where $f(q)$ is close to unity."
" The secondary mass AM» cau ouly be guessed at. but fortunately the radius depeuds ouly weakly ou M»: as a guide. the evolutionary. models calculated by span a rauge from 0.172 to 0.175 A£. at Pa,=E hr. implying 0.31xο.<Q.L8."," The secondary mass $M_2$ can only be guessed at, but fortunately the radius depends only weakly on $M_2$; as a guide, the evolutionary models calculated by \citet{bk00} span a range from 0.172 to 0.475 $M_{\odot}$ at $P_{\rm orb} = 4 $ hr, implying $0.34 \le R_2 / R_{\odot} \le 0.48$."
 Applying this to the surface brightuess rauge uxc above gives the secondarys absolute magnitude. Af=11.30.8.," Applying this to the surface brightness range found above gives the secondary's absolute magnitude, $M_V = 11.3 \pm 0.8$."
 The spectral decomposition gave V.—17.1x0.3 for the secondary. alone. which eives a distauce modulus η(—AL=5.8+0.9.," The spectral decomposition gave $V = 17.1 \pm
0.3$ for the secondary alone, which gives a distance modulus $m - M =
5.8 \pm 0.9$."
" Combining these gives a best distauce of 1LO(+80.—50) pc Note that this computation does not assume a ""normal. mass-racdits relation for the secondary — it is based iusteacd on tlie secoudarys surface brightness and the requirement that it fill the Roche lobe."," Combining these gives a best distance of $140 (+80,-50)$ pc Note that this computation does not assume a `normal' mass-radius relation for the secondary – it is based instead on the secondary's surface brightness and the requirement that it fill the Roche lobe."
 Our Bayesian analysis of the parallax and proper motion follows the procedure described by Thorstensen(2003)..., Our Bayesian analysis of the parallax and proper motion follows the procedure described by \citet{thorparallax}.
 As noted earlier. we measure Taps=7-21] mas.," As noted earlier, we measure $\pi_{\rm abs} = 7.2 \pm 1.1$ mas."
 We assume the samepriori distribution of space velocities as Thorsteuseu(2003).., We assume the same distribution of space velocities as \citet{thorparallax}.
 Using the parallax aud proper motion alone. the Bayesian analysis gives 120(—-28.—21) pc. consistent. with the spectrophotometric estimate.," Using the parallax and proper motion alone, the Bayesian analysis gives $150(+28,-21)$ pc, consistent with the spectrophotometric estimate."
 Foldiug the spectropliotoijetric estimate into the Bayesian analysis vields 119(2-26.—20) pc.," Folding the spectrophotometric estimate into the Bayesian analysis yields $149(+26,-20)$ pc."
 The final estimate is slightly ereater than L/z4544; because a correction for the Lutz-IJxelker bias increases distance estimates when tie relative parallaxerror is uot small., The final estimate is slightly greater than $1/\pi_{\rm abs}$ because a correction for the Lutz-Kelker bias increases distance estimates when the relative parallaxerror is not small.
 The rather small final uncertainty reflects the concordauce o all the information — the distance evidence is all nicely consistent., The rather small final uncertainty reflects the concordance of all the information – the distance evidence is all nicely consistent.
with L and M in solar units.,with $L$ and $M$ in solar units.
" This calibration is based on template-fitted masses (Bruzual&Charlot2003,ChabrierIMF) and luminosities derived for galaxies in the COSMOS field (Ilbertetal.2009b).", This calibration is based on template-fitted masses \citep[Chabrier IMF]{bruz03} and luminosities derived for galaxies in the COSMOS field \citep{ilbe09b}.
. We convert luminosities toVega zeropoint and apply a —0.124 dex mass offset to transform to the mass scale of the models by Charlot Bruzual (2007/2009 in prep.), We convert luminosities toVega zeropoint and apply a $-$ 0.124 dex mass offset to transform to the mass scale of the models by Charlot Bruzual (2007/2009 in prep.)
 that include contributions from TP-AGB stars., that include contributions from TP-AGB stars.
 The linear relation of eqn., The linear relation of eqn.
 2 is a fit to galaxies in redshift- (1.0<z 1.6) and color-range (0.38«(B—V)< 0.74) of our 7 host galaxies resolved in both HST bands., \ref{eqn:mass} is a fit to galaxies in redshift- $1.0<z<1.6$ ) and color-range $0.38<(B-V)<0.74$ ) of our 7 host galaxies resolved in both HST bands.
 The scatter of the fit corresponds to an RMS uncertainty in resulting stellar mass of +0.3 dex., The scatter of the fit corresponds to an RMS uncertainty in resulting stellar mass of $\pm$ 0.3 dex.
" We convert our measured (F814W-F160W) colors to restframe (B—V)post by applying both K-correction and a color term, individually for each object and its redshift."," We convert our measured (F814W–F160W) colors to restframe $(B-V)_\mathrm{host}$ by applying both $K$ -correction and a color term, individually for each object and its redshift."
" For this purpose we identify for each galaxy the single stellar population model again from Charlot Bruzual IMF, solar metallicity) with the closest (Chabrier at a given z."," For this purpose we identify for each galaxy the single stellar population model again from Charlot Bruzual (Chabrier IMF, solar metallicity) with the closest (F814W–F160W) at a given $z$ ."
 Since the interpolation intervals are (F814W-F160W)rather small and we are, Since the interpolation intervals are rather small and we are
the results obtained within this approximation do not show any appreciable difference with respect to those obtained by carrying out the complete calculations.,the results obtained within this approximation do not show any appreciable difference with respect to those obtained by carrying out the complete calculations.
 To emphasize the atomic aspects involved in the problem. avoiding complications due to radiative. transfer effects. we consider an optically thin slab of Se tons located 1000 km above the solar surface (as defined in Sect.," To emphasize the atomic aspects involved in the problem, avoiding complications due to radiative transfer effects, we consider an optically thin slab of Sc ions located 1000 km above the solar surface (as defined in Sect."
 12.3 of LLO4). and we assume that the slab is illuminated from below by the photospheric continuum (see Fig. 3)).," 12.3 of LL04), and we assume that the slab is illuminated from below by the photospheric continuum (see Fig. \ref{fig:slab}) )."
 Under these hypotheses. the atomic polarization can be calculated by solving the SEE directly for the given incident continuum radiation field.," Under these hypotheses, the atomic polarization can be calculated by solving the SEE directly for the given incident continuum radiation field."
 We remark that this model is basically the same as that employed by Belluzzietal.(2007) to model the Q// profile of the Bau D» line., We remark that this model is basically the same as that employed by \citet{Bel07} to model the $Q/I$ profile of the Ba $_2$ line.
 We assune that the continuum 1s unpolarized and cylindrically symnjetric. about the local vertical., We assume that the continuum is unpolarized and cylindrically symmetric about the local vertical.
 Under these assumptions. takiig a reference system with the z-axis (the quantization. axis) directed along the vertical. it can. be shown that only two conponents of the radiation field tensor. in terms of which we describe the incident continuum. are vanishing where µ is the cosine of the heliocentric angle.," Under these assumptions, taking a reference system with the $z$ -axis (the quantization axis) directed along the vertical, it can be shown that only two components of the radiation field tensor, in terms of which we describe the incident continuum, are non-vanishing where $\mu$ is the cosine of the heliocentric angle."
 The former quantity is the mean intensity of the radiation field (averaged over all directions). the second quantifies the anisotropy of the radiation field (imbalance between vertical and horizontal illumination).," The former quantity is the mean intensity of the radiation field (averaged over all directions), the second quantifies the anisotropy of the radiation field (imbalance between vertical and horizontal illumination)."
 The mean intensity of the radiation field can also be expressed in terms of the average number of photons per mode. 5. while the anisotropy degree is often quantified through the so-called anisotropy factor. w.," The mean intensity of the radiation field can also be expressed in terms of the average number of photons per mode, $\bar{n}$, while the anisotropy degree is often quantified through the so-called anisotropy factor, $w$."
 These new. non-dimensional quantities arerelated to Ju and J by the equations We calculate the values of n(v) and w(v) of the photospheric continuum following Sect.," These new, non-dimensional quantities arerelated to $J^0_0$ and $J^2_0$ by the equations We calculate the values of $\bar{n}(\nu)$ and $w(\nu)$ of the photospheric continuum following Sect."
 12.3 of LLO4 taking /;21000 km. and using the values of the disk center intensities. and of the limb-darkening coefficients given by Pierce(2000).," 12.3 of LL04 taking $h=1000$ km, and using the values of the disk center intensities, and of the limb-darkening coefficients given by \citet{Pie00}."
. At the height of the slab. and at the frequency of the spectral line under investigation (4247 A) we find @=0.158x107. and w=0.189.," At the height of the slab, and at the frequency of the spectral line under investigation (4247 ) we find $\bar{n}=0.158 \times 10^{-2}$, and $w=0.189$."
 Taking into account the atomie weight of scandium. assuming a temperature of 6000 K. and neglecting microturbulent velocity. we obtain for this line a Doppler width. Avly. given by where wy 1s the thermal velocity. 7 the kinetic temperature. M the mass of the ton. and Kp the Boltzmann constant.," Taking into account the atomic weight of scandium, assuming a temperature of 6000 K, and neglecting microturbulent velocity, we obtain for this line a Doppler width, $\Delta \lambda_D$, given by where $w_T$ is the thermal velocity, $T$ the kinetic temperature, $M$ the mass of the ion, and $k_B$ the Boltzmann constant."
 Once the SEE have been written down and solved numerically. we can calculate the radiative transfer coefficients according to the equations of Sect.," Once the SEE have been written down and solved numerically, we can calculate the radiative transfer coefficients according to the equations of Sect."
 7.9 of LLO4., 7.9 of LL04.
 We consider the radiation scattered by the slab at 907. and we take the reference direction for positive Q parallel to the slab (see Fig. 3»).," We consider the radiation scattered by the slab at $90^{\circ}$, and we take the reference direction for positive $Q$ parallel to the slab (see Fig. \ref{fig:slab}) )."
 In the case of a tangential observation. under the approximation ofa weakly polarizing atmosphere (2;>&g.&p.eyiniqoNeedy Po-pu-py). it can be shown that the emergent fractional polarization is given by (seeTrujilloBueno.2003) The first and second—go. terms in the right-hand side of Eq. (9))," In the case of a tangential observation, under the approximation ofa weakly polarizing atmosphere $\varepsilon_I \gg \varepsilon_Q, \varepsilon_U, 
\varepsilon_V; \eta_I \gg \eta_Q, \eta_U, \eta_V, \rho_Q, \rho_U, \rho_V$ ), it can be shown that the emergent fractional polarization is given by \citep[see][]{JTB03}
 The first and second terms in the right-hand side of Eq. \ref{eq:dichroism}) )"
 represent the contribution to the emergent radiation due to processes of selective emission and selective absorption (dichroism) of polarization components. respectively.," represent the contribution to the emergent radiation due to processes of selective emission and selective absorption (dichroism) of polarization components, respectively."
 We note that while the lower level of Bau D». with /=1/2. can carry alignment only because of the presence of HFS. the lower level of this scandium line. with 7= 2. can be polarized also neglecting HPS.," We note that while the lower level of Ba $_2$, with $J=1/2$, can carry alignment only because of the presence of HFS, the lower level of this scandium line, with $J=2$ , can be polarized also neglecting HFS."
 As a consequence. while in the case of the Ba D» line the lower level results to be significantly less," As a consequence, while in the case of the Ba $_2$ line the lower level results to be significantly less"
onboard l8 hours before the eruption with svuthetic nuages for the same filters from our siuulatiou in Figure 3.,onboard 18 hours before the eruption with synthetic images for the same filters from our simulation in Figure 3.
 Synthetic and real tages are plotted using the same scale., Synthetic and real images are plotted using the same scale.
 The filter respouse function peaks around L.l MIN. while the peaks around | MUN. illustrating different heights in the corona.," The filter response function peaks around 1.4 MK, while the peaks around 1 MK, illustrating different heights in the corona."
— Iu svuthetic aud real images. the three active regions are clearly visible as regions of cuhanced euission (hot aud dense): the large AR 10800 on the casteru side of the Sun near disk ceuter. AR 10797 near the western limb and AR LO7T9S near W50.," In synthetic and real images, the three active regions are clearly visible as regions of enhanced emission (hot and dense): the large AR 10800 on the eastern side of the Sun near disk center, AR 10797 near the western limb and AR 10798 near W50."
 Tn addition to the northern polar coronal hole. there are a number of equatorial coronal holes. includiusg two on the south-western side of ARs 10798 aud 10800.," In addition to the northern polar coronal hole, there are a number of equatorial coronal holes, including two on the south-western side of ARs 10798 and 10800."
 The southern polar coronal hole is alinost absent iu the simulated aud real images., The southern polar coronal hole is almost absent in the simulated and real images.
 Overall. there is good agreement between svuthetic aud real nuages. which gives us confidence that our model of the the corona is a relatively realistic represcutation of the actual corona at the time of the 2005 Aneust 22 eruption.," Overall, there is good agreement between synthetic and real images, which gives us confidence that our model of the the corona is a relatively realistic representation of the actual corona at the time of the 2005 August 22 eruption."
 The most important features for our study that the model reproduces are the appearance of AR 10795 and the presence and aspect of open field regions (dark) around it., The most important features for our study that the model reproduces are the appearance of AR 10798 and the presence and aspect of open field regions (dark) around it.
 The main difference between sinuulated aud real images is the eastern Linh of the Sun where the modeled enmüssiou is too weak as compared to the real ouc., The main difference between simulated and real images is the eastern limb of the Sun where the modeled emission is too weak as compared to the real one.
 This is relatively wuimportaut because these regions are far from the source region of the CME (nore than 120? separation) aud are not involved iu the eruption process., This is relatively unimportant because these regions are far from the source region of the CME (more than $^\circ$ separation) and are not involved in the eruption process.
 As soon as we superpose the flux rope onto the steady-state coronal maeguctic field. it crupts due to force inbaliuce.," As soon as we superpose the flux rope onto the steady-state coronal magnetic field, it erupts due to force imbalance."
 As in other simulations with the same CALE model. the exact lducmatics of the CME early on in the corona are not realistic as the CAIE reaches its κμπα speed (~1500 kin +) about 1.5 minutes after the superposition of the fiux rope.," As in other simulations with the same CME model, the exact kinematics of the CME early on in the corona are not realistic as the CME reaches its maximum speed $\sim 1500$ km $^{-1}$ ) about 1.5 minutes after the superposition of the flux rope."
 However. the simulated CALE kinematics past 2 .. (after 15 minutes) are In good agreement with the height-time profile as observed in LASCO/C2.," However, the simulated CME kinematics past 3 $_\odot$ (after 15 minutes) are in good agreement with the height-time profile as observed in LASCO/C2."
 Additionally. the CATE speed after L hour js about 1200 lau |. similar to what is ucasured with LASCO (~1250 kus 13.," Additionally, the CME speed after 1 hour is about 1200 km $^{-1}$, similar to what is measured with LASCO $\sim 1250$ km $^{-1}$ )."
 It confirms that he total cnerey of the simulated CALE is comparable to hat of the real CALE., It confirms that the total energy of the simulated CME is comparable to that of the real CME.
 The model does not iuteud to capture the slow rise phase before the loss of equilibrium jor the acceleration plase., The model does not intend to capture the slow rise phase before the loss of equilibrium nor the acceleration phase.
 This is why when comparing svuthetic and real tages. it is best not to use tle ouset iue of the flare as the starting time of the nuucerical simulation but a later time when the CATE has already sienificautly accelerated.," This is why when comparing synthetic and real images, it is best not to use the onset time of the flare as the starting time of the numerical simulation but a later time when the CME has already significantly accelerated."
 Previous studies have fouud hat the CATE acceleration happens duriug the N-Rax rise phase (Oliviuua&Shibata1998:Forbes2000:Tenueretal. 2008).," Previous studies have found that the CME acceleration happens during the X-Ray rise phase \citep[]{Ohyama:1998, Forbes:2000, Temmer:2008}."
. We use GOELES-12 (Illetal.2005) and RIIESSI (Linetal.2002). data to investigate the flare time in soft N-rav (SNR) and hard N-rav (ITXR respectively.," We use GOES-12 \citep[]{Hill:2005} and RHESSI \citep[]{Lin:2002} data to investigate the flare time in soft X-ray (SXR) and hard X-ray (HXR), respectively."
 For the ejection. the flare ousct in SNR was). 0:LEUTT. the TINR dare started at 01:02UT aud the fare veaked at 01:22UT.," For the ejection, the flare onset in SXR was 00:44UT, the HXR flare started at 01:02UT and the flare peaked at 01:22UT."
 We believe the onset of the IEXR fare is the best time to use for the start time of our simulation suce the CALE was already observed by LASCO at LR. 10 aninutes after the fare peak., We believe the onset of the HXR flare is the best time to use for the start time of our simulation since the CME was already observed by LASCO at 4 $_\odot$ 10 minutes after the flare peak.
 Figure 1 shows a line-of-sight image of the CME observed by LASCO/C2 aud processed with the method of Moreauetal(2006) and a svuthetic line-of-sight inuaese from our simulation., Figure 4 shows a line-of-sight image of the CME observed by LASCO/C2 and processed with the method of \citet{Morgan:2006} and a synthetic line-of-sight image from our simulation.
 The LASCO image is taken at 01:51 UT. 52 niuutes after the onset of the ITNR flare. while the svuthetic inuage is made 15 münutes after the superposition of the flux rope onto the solar surface.," The LASCO image is taken at 01:54 UT, 52 minutes after the onset of the HXR flare, while the synthetic image is made 45 minutes after the superposition of the flux rope onto the solar surface."
 The svuthetic inage is processed as explained in Lugazctal.(2009) usine a svuthetic 27-day nini dnage created from the steady-state simulation., The synthetic image is processed as explained in \citet{Lugaz:2009b} using a synthetic 27-day minimum image created from the steady-state simulation.
 The main difference between the svuthetic aud real images is the latitudinal direction of propagation of the CME., The main difference between the synthetic and real images is the latitudinal direction of propagation of the CME.
 Iu the svuthetic inaese. the fastest moving part of the CALE is at a PA," In the synthetic image, the fastest moving part of the CME is at a PA"
The Tull-Fisher relation (Tully&Fisher1977.here-afterTER) is au eiipirical relation between the absolute magnitude AY and the iiaxiumun rotation velocity Viuas of spiral galaxies. which is typically expressed as where the values of slope a and zero-point ./ are dependent on the photometric baud of the AL beine micasured.,"The Tully-Fisher relation \citep[][hereafter TFR]{Tully77} is an empirical relation between the absolute magnitude $M$ and the maximum rotation velocity $\Vmax$ of spiral galaxies, which is typically expressed as where the values of slope $\alpha$ and zero-point $\beta$ are dependent on the photometric band of the $M$ being measured."
 It las long been known that the TER has a morphological type depeudenuce. with earlier type spirals having systematically lower huuinositv at fixed Vias (Roberts1978:Rubinetal.1980).," It has long been known that the TFR has a morphological type dependence, with earlier type spirals having systematically lower luminosity at fixed $\Vmax$ \citep{Roberts78,Rubin80}."
. This luminosity difference is also waveband dependent. with larger offsets in shorter wavelengths.," This luminosity difference is also waveband dependent, with larger offsets in shorter wavelengths."
 For example. in 7 baud. Cüovanellictal.(1997). found a 0.32 mae lower zero-point of the TER for Sa/Sab galaxies aud 0.10 mae lower for Sb ealaxies than the She aud later type spirals.," For example, in $I$ band, \citet{Giovanelli97} found a 0.32 mag lower zero-point of the TFR for Sa/Sab galaxies and 0.10 mag lower for Sb galaxies than the Sbc and later type spirals."
 Iu D baud. Russell(2001). found. that Sb ealaxies have a zero-point of 0.57 mae lower than Sc galaxies.," In $B$ band, \citet{Russell04} found that Sb galaxies have a zero-point of 0.57 mag lower than Sc galaxies."
 In a recent study. Russell(2008) compared the TER-derived distances of nearby eroups aud clusters aud fod a mean difference of 0.19 mae in Z7 baud between the Sb aud Sc spiral galaxies.," In a recent study, \cite{Russell08} compared the TFR-derived distances of nearby groups and clusters and found a mean difference of 0.19 mag in $H$ band between the Sb and Sc spiral galaxies."
 Moreover. using a large sample of spiral ealaxies. Mastersctal.(2006.2008) found that the morphological dependence of the TER is uot only a shift of the zero-point. but even dependent on the Iunuinositv in the wav that the differences are more pronounced for more hIuninous galaxies.," Moreover, using a large sample of spiral galaxies, \cite{Masters06,Masters08} found that the morphological dependence of the TFR is not only a shift of the zero-point, but even dependent on the luminosity in the way that the differences are more pronounced for more luminous galaxies."
 The morphological dependence of the TER originates from either the differences of the stellar population(AL) or the disk dynamics (πας) or both., The morphological dependence of the TFR originates from either the differences of the stellar $M$ ) or the disk dynamics $\Vmax$ ) or both.
 ~Ju one hand. it is well known that the colors of earlier type spirals ave redder.," On one hand, it is well known that the colors of earlier type spirals are redder."
 Devereux&Young(1991) interpreted this color difference as originating from the bulee disk composition effect., \citet{Devereux91} interpreted this color difference as originating from the bulge disk composition effect.
 The stellar population of bulges are typically older aud more metal rich than that of the disks., The stellar population of bulges are typically older and more metal rich than that of the disks.
 A laveer fraction of the bulee component iu earlier type spirals naturally results iu au redder color ou average for the whole galaxy., A larger fraction of the bulge component in earlier type spirals naturally results in an redder color on average for the whole galaxy.
 However. with a bulge disk decomposed sample. I&eunicuttetal.(1991) studied aud compared the star formation histories of oulv the disk conrponeut of different type spirals and found that the stellar population of the disks of later type spiral is also on average vounger.," However, with a bulge disk decomposed sample, \citet{Kennicutt94} studied and compared the star formation histories of only the disk component of different type spirals and found that the stellar population of the disks of later type spiral is also on average younger."
 Ou the other haud. although the rotation curves of spiral galaxies are proposed to follow a universal shape Porsic&Salueci1991:Persicetal.1996).. there is evidence showing that the rotation curves of carly type spirals rise more rapidly in the iuner region than that of late types (Corradi&Capaccioli1990:Noordenueeretal. ," On the other hand, although the rotation curves of spiral galaxies are proposed to follow a universal shape \citep{Persic91,Persic96}, there is evidence showing that the rotation curves of early type spirals rise more rapidly in the inner region than that of late types \citep{Corradi90,Noordermeer07a}. ."
2007).. Noordermecr&Verheijen(2007) show that the massive Sa galaxies lie better in the well-defined TER wheu using the asviuptotie rotation velocity Visa instead of Vyas. naplviug a dependence of Via; ou the morphological type.," \citet{Noordermeer07b} show that the massive Sa galaxies lie better in the well-defined TFR when using the asymptotic rotation velocity $V_{\rm{asymp}}$ instead of $\Vmax$, implying a dependence of $\Vmax$ on the morphological type."
 Ou the theoretical side. the zero-poiut. slope and scatter of the observed TER cau all be well accommodated by the current disk formation iuodel in the framework of the cold dark matter hierarchical cosinogomes (Dalcautouctal.1997:Moet1998:Mo&Mao2000:Pizaguoetal. 2005).," On the theoretical side, the zero-point, slope and scatter of the observed TFR can all be well accommodated by the current disk formation model in the framework of the cold dark matter hierarchical cosmogonies \citep{Dalcanton97,MMW,Mo00,Pizagno05}."
. ILowever. the morphological dependence of the TFRs has not becu probed iu these studies because the bulee component imn these models is typically neglected.," However, the morphological dependence of the TFRs has not been probed in these studies because the bulge component in these models is typically neglected."
 Another lnunitatiou of these models is that the stellar populations have uot been tackled with a plivsical prescription. but with determined mass-to-lielt ratios.," Another limitation of these models is that the stellar populations have not been tackled with a physical prescription, but with pre-determined mass-to-light ratios."
 Tn this study. we annued to model the dynamics anc stellar population of different type spiral sealaxies in combination aud try to find out which factor is the main contributor to the morphological depeudence of the TER.," In this study, we aimed to model the dynamics and stellar population of different type spiral galaxies in combination and try to find out which factor is the main contributor to the morphological dependence of the TFR."
 Tn specific. we will follow the dynamical model of Mo. Mao White (1998. hereafter AIAIW) aud extend it to iuclude a bulge couponcut.," In specific, we will follow the dynamical model of Mo, Mao White (1998, hereafter MMW) and extend it to include a bulge component."
 For the stellar population. we will parameterize the starformation histories of the disks and bulges separately aud then derive thei mass-to- ratios iu differcut bands using the stellar population," For the stellar population, we will parameterize the starformation histories of the disks and bulges separately and then derive their mass-to-light ratios in different bands using the stellar population"
members HD4909| and HD49068 respectively.,members HD49091 and HD49068 respectively.
 Cameron (1985) used (51> photometry of 14+ members to derive [Fe/H]=+0.065 while from Strómmgren photometry of F star members. Nissen (1988) determined [Fe/H]=-0.10.," Cameron (1985) used $UBV$ photometry of 14 members to derive [Fe/H]=+0.065 while from Strömmgren photometry of F star members, Nissen (1988) determined [Fe/H]=-0.10."
 Collectively these results point towards the metalicity of NGC2287 being close to the solar value. perhaps marginally less.," Collectively these results point towards the metalicity of NGC2287 being close to the solar value, perhaps marginally less."
 The youthful age. the relative proximity and the low line of sight reddening of NGC2287 make it a particularly suitable target for the study of the IFMR. as was recognised several decades ago.," The youthful age, the relative proximity and the low line of sight reddening of NGC2287 make it a particularly suitable target for the study of the IFMR, as was recognised several decades ago."
 Indeed. Romanishin Angel (1980) undertook a search of the cluster using photographic plates and this led to the identification of tive candidate white dwarf members.," Indeed, Romanishin Angel (1980) undertook a search of the cluster using photographic plates and this led to the identification of five candidate white dwarf members."
 Follow-up spectroscopic observations obtained with the 3.6m ESO telescope and the Image Dissector Scanner contirmed at least three of these objects to be white dwarfs. two with Lar25000K and one with Puree l3000K (Koester Reimers 1981).," Follow-up spectroscopic observations obtained with the 3.6m ESO telescope and the Image Dissector Scanner confirmed at least three of these objects to be white dwarfs, two with $T$$_{\rm eff}$$\approx$ 25000K and one with $T$$_{\rm eff}$$\approx$ 13000K (Koester Reimers 1981)."
 Tt was concluded by these authors that the two hotter stars were likely to be cluster members while the cooler white dwarf was probably a foreground object., It was concluded by these authors that the two hotter stars were likely to be cluster members while the cooler white dwarf was probably a foreground object.
 However. as the quality of the existing spectral data is rather poor. no detailed analyses of the Balmer line profiles in the spectral energy distributions of these stars has ever been undertaken.," However, as the quality of the existing spectral data is rather poor, no detailed analyses of the Balmer line profiles in the spectral energy distributions of these stars has ever been undertaken."
 Thus robust estimates of the effective temperatures and surface gravities of these objects are unavailable to contirm or otherwise these conclusions and to allow them to be fully exploited in the context of the IFMR., Thus robust estimates of the effective temperatures and surface gravities of these objects are unavailable to confirm or otherwise these conclusions and to allow them to be fully exploited in the context of the IFMR.
 Low resolution. high signal-to-noise optical spectroscopy of the nine white dwarf candidate members of the clusters NGC3532 and NGC2287 was obtained in service mode with the European Southern Observatory (ESO) Very Large Telescope and the Focal Reducer and low dispersion Spectrograph (FORSI) within the periods 2007/04/24-27 and 2007/10/06-11/21.," Low resolution, high signal-to-noise optical spectroscopy of the nine white dwarf candidate members of the clusters NGC3532 and NGC2287 was obtained in service mode with the European Southern Observatory (ESO) Very Large Telescope and the Focal Reducer and low dispersion Spectrograph (FORS1) within the periods 2007/04/24-27 and 2007/10/06-11/21."
 A full description of the FORSI instrument may be found on the ESO As these targets are comparatively bright. we specitied fairly relaxed constraints on the sky conditions and thus the observations were generally undertaken in poorer seeing and/or with some cloud oesent.," A full description of the FORS1 instrument may be found on the ESO As these targets are comparatively bright, we specified fairly relaxed constraints on the sky conditions and thus the observations were generally undertaken in poorer seeing and/or with some cloud present."
 All data were acquired using the 2 binning mode of he E2V CCD. the 600B-12 grism and a 1.67 slit which gives a nominal resolution of A/NXA-- 500.," All data were acquired using the $\times$ 2 binning mode of the E2V CCD, the 600B+12 grism and a 1.6"" slit which gives a nominal resolution of $\lambda$ $\Delta$$\lambda$$\sim$ 500."
 Flat and are exposures were obtained within a few hours of the acquisition of each of the science Tames., Flat and arc exposures were obtained within a few hours of the acquisition of each of the science frames.
 The CCD data were debiased and flat fielded using the IRAF procedure CCDPROC., The CCD data were debiased and flat fielded using the IRAF procedure CCDPROC.
 Cosmic ray hits were removed using he routine LACOS SPEC (van Dokkum 2001)., Cosmic ray hits were removed using the routine LACOS SPEC (van Dokkum 2001).
 Subsequently he spectra were extracted using the APEXTRACT package and wavelength calibrated by comparison with the He+HgCd are spectra., Subsequently the spectra were extracted using the APEXTRACT package and wavelength calibrated by comparison with the He+HgCd arc spectra.
 Remaining instrument signature was removed using a spectrum of the featureless DC white dwarf WD00004345 obtained with an identical set-up during this programme., Remaining instrument signature was removed using a spectrum of the featureless DC white dwarf WD0000+345 obtained with an identical set-up during this programme.
 We have used recent versions of the plane-parallel. hydrostatic. non-local thermodynamic equilibrium (non-LTE) atmosphere and spectral synthesis codes TLUSTY (v200: Hubeny 1988. Hubeny Lanz 1995) and SYNSPEC (v48: Hubeny. I. and Lanz. T. 2001. http://novacastro.umd.edu/) to generate a grid of pure-H synthetic spectra covering the Zr and surface gravity ranges 14000-38000K and log g=7.25-8.75 respectively.," We have used recent versions of the plane-parallel, hydrostatic, non-local thermodynamic equilibrium (non-LTE) atmosphere and spectral synthesis codes TLUSTY (v200; Hubeny 1988, Hubeny Lanz 1995) and SYNSPEC (v48; Hubeny, I. and Lanz, T. 2001, http://nova.astro.umd.edu/) to generate a grid of pure-H synthetic spectra covering the $T$$_{\rm eff}$ and surface gravity ranges 14000-35000K and log $g$ =7.25-8.75 respectively."
 We have employed a model H atom incorporating the 8 lowest energy levels and one superlevel extending from n=9 to nz80. where the dissolution of the high lying levels was treated by means of the occupation probability formalism of Hummer Mihalas (1988). generalised," We have employed a model H atom incorporating the 8 lowest energy levels and one superlevel extending from n=9 to n=80, where the dissolution of the high lying levels was treated by means of the occupation probability formalism of Hummer Mihalas (1988), generalised"
Camunie 2001. Jcyhoson CGaiinuie 2003.) and Duriseu 2006 for recent discussions).,"Gammie 2001, Johnson Gammie 2003, and Durisen 2006 for recent discussions)."
 This is au especially attractive possibility in the earliest stages of stellar evolutiou. as it is quite likely that inmost of the 1dass Of stars is accreted through the circuustellar disk. due to the finite (aud Large) angular iionuenta of protostellar clot(ls.," This is an especially attractive possibility in the earliest stages of stellar evolution, as it is quite likely that most of the mass of stars is accreted through the circumstellar disk, due to the finite (and large) angular momenta of protostellar clouds."
 On the other haud. T Tauri stars have clearly already accreted most of their 1lass. and typical disk mass estityates are ai order of magnitude below the values recuired lor gravitational instaility (e.g... Beckwithet 1990: Andrews Williams 2005).," On the other hand, T Tauri stars have clearly already accreted most of their mass, and typical disk mass estimates are an order of magnitude below the values required for gravitational instability (e.g., Beckwith 1990; Andrews Williams 2005)."
 In addition. lie surface densities reculred for gravitational instability appear to be implausibly large iu the i1nernost ¢isk: some other nechanisin of tansport must operate there.," In addition, the surface densities required for gravitational instability appear to be implausibly large in the innermost disk; some other mechanism of transport must operate there."
 In tlis paper we ¢iscuss the curreu evidence corcerniug T Tauri accretion rates and disk llasses ald jelr üinplications for tje mechaUstus of mass and angular momentum trausport., In this paper we discuss the current evidence concerning T Tauri accretion rates and disk masses and their implications for the mechanisms of mass and angular momentum transport.
 We present a revision of tle stinple laverec model of Cauunie to include irradiatiou by the central star. which itrocduces some depeuxence προi stellar mass through the depeudenuce of disk heating ou the stellar luminosity.," We present a revision of the simple layered model of Gammie to include irradiation by the central star, which introduces some dependence upon stellar mass through the dependence of disk heating on the stellar luminosity."
 TIe inclusionM. «of ir‘acliatiou does not appear to vield a sUliciently steep depeuclence of nass accretion rate OL stellar uass to explain all t1 observations. alhough it might ye consistent. with the ipper enveope of estimated accretion raes.," The inclusion of irradiation does not appear to yield a sufficiently steep dependence of mass accretion rate on stellar mass to explain all the observations, although it might be consistent with the upper envelope of estimated accretion rates."
 A fully viscous disk might be able to explain the observations if tle initial disk size is 5rougly correlated with tje. stellar mass. ut it is uot clear why clisks aroun al least the most assive stars should be ionizec| down through o thetr midplanes at al radii.," A fully viscous disk might be able to explain the observations if the initial disk size is strongly correlated with the stellar mass, but it is not clear why disks around at least the most massive stars should be ionized down through to their midplanes at all radii."
 We sweest that the observed behavior of mass acc'ellon in young stars nay be a ¢'omplicated mix o ‘layered accretion witl possible gravitational iustability iu the nost massive stars aud ull viscous accretion iu the lowesπο objects., We suggest that the observed behavior of mass accretion in young stars may be a complicated mix of layered accretion with possible gravitational instability in the most massive stars and full viscous accretion in the lowest-mass objects.
 Observations spauniug the range from 2-3 M. [or interimecdiate-1nass T Tauri stars to brown cdwarfs (Calvet 2001: Mwerolle 2003a. 2005) suggest a strong dependeuce of mass accretion rate oi stellar mass. roughly {ixAl? (Figure 1).," Observations spanning the range from 2-3 $\msun$ for intermediate-mass T Tauri stars to brown dwarfs (Calvet 2004; Muzerolle 2003a, 2005) suggest a strong dependence of mass accretion rate on stellar mass, roughly $\mdot \propto M_*^2$ (Figure 1)."
 While observational selection effects tend to produce some correlatio1. because low accretion rates are uot detectable in higlier-DIumniuosity stars. the overal trend of decreasiug acc'etiou rate with decreasing mass is clear (ef," While observational selection effects tend to produce some correlation, because low accretion rates are not detectable in higher-luminosity stars, the overall trend of decreasing accretion rate with decreasing mass is clear (cf."
 Muzerolle 2005)., Muzerolle 2005).
" As shewht in Figure 1. there is a large scatter in AZ at agiven stellar mass AL, among populations of siular age. so that mass is uot the only. parameter controlling accretion."," As shown in Figure 1, there is a large scatter in $\mdot$ at agiven stellar mass $M_*$ among populations of similar age, so that mass is not the only parameter controlling accretion."
 Oue well-laowlmechaulsi1 which p'oduces an Azο dependence of the accretion rate on stellar inass is Boucli-Hovle accretion:," One well-knownmechanism which produces an $M_*^2$ dependence of the accretion rate on stellar mass is Bondi-Hoyle accretion: = ,"
systems is high enough out to impact parameter of 40 kpe in the galactic disk that continuum racliation passing hrough the plane of the disk will always be subjected to Lyman limit absorption.,systems is high enough out to impact parameter of $\sim 40$ kpc in the galactic disk that continuum radiation passing through the plane of the disk will always be subjected to Lyman limit absorption.
 Thus. if GRBs occur at olfsets 40 kpe from the centres of their host galaxy. it is expected hat roughly 50 percent (half in front of the clisk. half behind he disk) of CRB alterelow will have a Lyman limit. break in the spectrum: this Lyman limit will correspond. to the orecise redshift of the burst since the limit. svsten will be ocal to the GRB.," Thus, if GRBs occur at offsets $<< 40$ kpc from the centres of their host galaxy, it is expected that roughly 50 percent (half in front of the disk, half behind the disk) of GRB afterglow will have a Lyman limit break in the spectrum; this Lyman limit will correspond to the precise redshift of the burst since the limit system will be local to the GRB."
 Intervening galaxies. not associated with the GRB bu in the lineofsieht of the afterglow. may also absorb the continuum: thus a redshift inferred from a Lyman limit wil not necessarily be the redshift of the GRB.," Intervening galaxies, not associated with the GRB but in the line–of–sight of the afterglow, may also absorb the continuum; thus a redshift inferred from a Lyman limit will not necessarily be the redshift of the GRB."
 What cllect wil this have on the determination of GRB recdshilts?, What effect will this have on the determination of GRB redshifts?
 Storric-Lombardi et al. (, Storrie-Lombardi et al. (
1994) survey QSO absorption spectra and find that for any random lineofsight. the density of intervening Lyman limit svstems is NV(2)ο0.38(1|Qj for redshifts z«3.0.,"1994) survey QSO absorption spectra and find that for any random line–of–sight, the density of intervening Lyman limit systems is $N(z) \simeq 0.38 (1 +
z)^{1.04}$ for redshifts $z < 3.0$."
 hus at redshifts 21.5 it is expecte that most GRB afterglows will be subjected to at least one Lyman limit in their continuum that does not necessarily correspond to the intrinsic redshift of the GIU., Thus at redshifts $z \age 1.5$ it is expected that most GRB afterglows will be subjected to at least one Lyman limit in their continuum that does not necessarily correspond to the intrinsic redshift of the GRB.
 IC GIBs are ejected to distances comparable to the scale length of the Lyman limit absorption svstems in the disk. then the probability that the host galaxy will absorb the spectrum shortward of the Lymer limit is reduced.," If GRBs are ejected to distances comparable to the scale length of the Lyman limit absorption systems in the disk, then the probability that the host galaxy will absorb the spectrum shortward of the Lyman limit is reduced."
 The top portion of figure (1)) shows the expected. probability of the existence of a Lyman limit (7> 1) in the spectrum ofa GRB alfterglow as a function of redshift and the ollset scale length of GRBs from their host galaxy., The top portion of figure \ref{fig:lls}) ) shows the expected probability of the existence of a Lyman limit $\tau \ge 1$ ) in the spectrum of a GRB afterglow as a function of redshift and the offset scale length of GRBs from their host galaxy.
 The relationship between the frequency. of absorption from the host galaxy ancl the olfset scale is computed by assuming that absorption only occurs ifthe CRB is seen through the 40 kpe absorbing disk and that the disk has random viewing inclination., The relationship between the frequency of absorption from the host galaxy and the offset scale is computed by assuming that absorption only occurs if the GRB is seen through the 40 kpc absorbing disk and that the disk has random viewing inclination.
 As seen. the frequeney of Lyman limit absorption in the spectrum of GRBs at low redshifts (2< 1) may be used to determine the intrinsic ollset of GRBs from their host galaxies since most Lyman limit absorption at low redshifts is expected to come the host galaxy.," As seen, the frequency of Lyman limit absorption in the spectrum of GRBs at low redshifts $z \ale 1$ ) may be used to determine the intrinsic offset of GRBs from their host galaxies since most Lyman limit absorption at low redshifts is expected to come the host galaxy."
 Phe bottom half of figure (1)) shows. as à function of offset scale and redshift. the probability that the inferred. redshift is within 20 percent. of the redshift of the GIU. afterglow.," The bottom half of figure \ref{fig:lls}) ) shows, as a function of offset scale and redshift, the probability that the inferred redshift is within 20 percent of the redshift of the GRB afterglow."
 LW GRBs occur with about 60 kpe of their host galaxy (solid lino). then more than 60 percent of the redshifts inferred from the Lyman limit in the spectrum will be a mocerately accurate («20 percent) measure of the redshift of the CRB.," If GRBs occur with about 60 kpc of their host galaxy (solid line), then more than 60 percent of the redshifts inferred from the Lyman limit in the spectrum will be a moderately accurate $< 20$ percent) measure of the redshift of the GRB."
 Absorption [rom Ale Η. with a galactic impact parameter of ~50 kpc (Bergeron. ct al.," Absorption from Mg II, with a galactic impact parameter of $\sim 50$ kpc (Bergeron et al."
 L994). and € LV is expected. but not required. to accompany a Lyman limit svstem. (Storrie-Lombardi ct al.," 1994), and C IV is expected, but not required, to accompany a Lyman limit system (Storrie-Lombardi et al."
 L994)., 1994).
 Thus. figure (1)) could also be scen as a prediction. of the frequency of absorption lines in GRB afterglows.," Thus, figure \ref{fig:lls}) ) could also be seen as a prediction of the frequency of absorption lines in GRB afterglows."
 Indeed. lor CRB 970508. both Ale LE and Fe LE absorption was detected using Ixeck (Metzger et al.," Indeed for GRB 970508, both Mg II and Fe II absorption was detected using Keck (Metzger et al."
 1997): Arav οσο (1997) have found that the absence of detectable € IV. in the spectrum of the afterglow. limits the redshift of the GRB to ς<Ls (95 percent confidence), 1997); Arav Hogg (1997) have found that the absence of detectable C IV in the spectrum of the afterglow limits the redshift of the GRB to $z \le 1.8$ (95 percent confidence).
 The STIS CCD instrument onboard UST with a G230LB lowresolution grating gives complete spectral coverage over the wavelength range 16853065A., The STIS CCD instrument onboard HST with a G230LB low–resolution grating gives complete spectral coverage over the wavelength range 1685–3065.
. Phis would be sullicient to detect the Lyman limit over the redshift range (0.85<2< 24)., This would be sufficient to detect the Lyman limit over the redshift range $0.85 \ale z \ale 2.4$ ).
 The preferred instrument for obtaining redshifts from UW spectra. however. are the Near UV. (NUV) and Far UV (PUY) Aulti-Anode Microchannel Array (ALAALA) detectors. whieh is more sensitive anc has much wider spectral range (FU: L1501736X: NUV: 15709150 A) corresponding to a broad redshift range (0.26<2zz.9 and 0.12zizz;2.5. respectively).," The preferred instrument for obtaining redshifts from UV spectra, however, are the Near UV (NUV) and Far UV (FUV) Multi-Anode Microchannel Array (MAMA) detectors which is more sensitive and has much wider spectral range (FUV: 1150–1736; NUV: 1570–3180 ) corresponding to a broad redshift range $0.26 \ale z \ale .9$ and $0.72 \ale z \ale 2.5$, respectively)."
 ‘Temporal fits to the broadband: spectra. of. the optical transicnts GRB 970228 (Van Paraclijs et al., Temporal fits to the broadband spectra of the optical transients GRB 970228 (Van Paradijs et al.
 1997) and GRB 970508 (Bond 1997) indicate that the GRB afterglow may, 1997) and GRB 970508 (Bond 1997) indicate that the GRB afterglow may
ος 250 Lucvs (0971) iterative decovolution scheme. as adapter and developed by Waricls(1988). was used to derive he HI surtace nass doeusitv.,"KK98 250, Lucy's (1974) iterative decovolution scheme, as adapted and developed by \cite{warmel88}, was used to derive the HI surface mass density."
 Fig. 3|[, Fig. \ref{fig:smd}[ [
À] shows the best fit ITI racial profie obtaiuec from the 26421 resolutiou III image usiug Lucvs scheuie.,A] shows the best fit HI radial profile obtained from the $''\times21''$ resolution HI image using Lucy's scheme.
" Note that the deconvoved TT surfac' desity oxofile shows a steep rise of ~L0% oein the coutral ~25, ic. wihin oue svuthesized yal."," Note that the deconvolved HI surface density profile shows a steep rise of $\sim$ in the central $\sim 25^{''}$, i.e. within one synthesized beam."
 This is likely o be an artifact produced by the deconvolujon: simular artifacts xodnueed 1w this method rave been seen earlier. by for e.g.Yo Swaters (1999).," This is likely to be an artifact produced by the deconvolution; similar artifacts produced by this method have been seen earlier, by for e.g. Swaters (1999)."
 We herefore fit the surface denusitv Mute) by a. double Cassia (shown as dotted line in Fig. 3] , We therefore fit the surface density $\Sigma_{\rm{HI}}(r)$ by a double Gaussian (shown as dotted line in Fig. \ref{fig:smd}[ [
A|]. aud use only he broad compoucut for further anavais.,"A]), and use only the broad component for further analysis."
 For reference he full double Canes fitis elven by: with ryy rodlLI£087. 5.0dOAM.pe iud NooX=L24x02M.pe7.," For reference the full double Gaussian fit is given by: with $r_{01}=73.9^{''}\pm1.7^{''}$, $r_{02}=14.4^{''}\pm0.8^{''}$, $\Sigma_{01}=
5.0\pm0.1 \rm{M_\odot}~pc^{-2}$ and $\Sigma_{02}=4.2\pm0.2 \rm{M_\odot}~pc^{-2}$."
 For KI98 251. the III surface chsity profile was obtained iun the usual wav from fitting elliptical aunuli.," For KK98 251, the HI surface density profile was obtained in the usual way from fitting elliptical annuli."
 The surface deusitv Mitre) (shown in Fie. 3|[, The surface density $\Sigma_{\rm{HI}}(r)$ (shown in Fig. \ref{fig:smd}[ [
"B]} was obtained from the integrated III cohnu density image at 26""421"" resolution nuage aud is well represeuted bv a Cassia: with ry312.607£0.77 .¢=19.2.£08 and Ny2τὸἙFig.264.TheLILvelocitvficlelsofIxIxX98250andWKAOS5QAAL.pe 7.","B]) was obtained from the integrated HI column density image at $''\times21''$ resolution image and is well represented by a Gaussian: with $r_0=34.2^{''}\pm0.7^{''}$, $19.2^{''}\pm0.8^{''}$ and $\Sigma_0=7.8\pm0.1 \rm{M_\odot}~pc^{-2}$ ."
 The velocity fields o ASTS98 250 and 1ος 251 derived frou the 26421 resolution III data cube are showm Fie. L.," The velocity fields of KK98 250 and KK98 251, derived from the $''\times21''$ resolution HI data cube are shown in Fig. \ref{fig:mom1}."
 The velocity fields are regular aud the isovelocity contours are approxinately parallel. which is a signanre of rigid body rotation.," The velocity fields are regular and the isovelocity contours are approximately parallel, which is a signature of rigid body rotation."
 The velocity field of IKI&98251 is slightly lopsided. isovelocity contours 1i the north-casteri half of the eal are more curved than the south-westeru half.," The velocity field of KK98 251 is slightly lopsided, the isovelocity contours in the north-eastern half of the galaxy are more curved than the south-western half."
 Some s scale kinks are also seen in the isovelocity coutours m outer regions of the galaxy. indicative of a weak war of the outer disk.," Some small scale kinks are also seen in the isovelocity contours in the outer regions of the galaxy, indicative of a weak warping of the outer disk."
 These effects are more prominent iu lugher spatial resolution velocity fields(not shown). w also show a auld twis tin the kinematical major axis iu outer region of the galaxy. consistent with warpine.," These effects are more prominent in the higher spatial resolution velocity fields (not shown), which also show a mild twist in the kinematical major axis in the outer region of the galaxy, consistent with warping."
 The rotation curves of galaxies were derived separately or cach of the 26-—<2 217.16the«11 and 11.57«10 velocity fields. using the tilted rius model (Warner et al.," The rotation curves of the galaxies were derived separately for each of the $26^{''}\times21^{''}$, $16^{''}\times14^{''}$ and $11.5^{''}\times10^{''}$ velocity fields, using the usual tilted ring model (Warner et al."
 L973)., 1973).
 For the edee «n ealaxy λος 250. as described iu nore detai below. the rotation curve was also derived by fitting to tje position-velocitv diagram.," For the edge on galaxy KK98 250, as described in more detail below, the rotation curve was also derived by fitting to the position-velocity diagram."
 For cach galaxy. the ceuter and systemic velocity obtained from a global fit to the various resolution velocity fields matered within the error bars: the svstemuc velocity," For each galaxy, the center and systemic velocity obtained from a global fit to the various resolution velocity fields matched within the error bars; the systemic velocity"
in between indicates models with dark energy densities that are monotonie functions of lime.,in between indicates models with dark energy densities that are monotonic functions of time.
 Note that the weak enerey condition requires that the total equation of slate w>—1., Note that the weak energy condition requires that the total equation of state $w \geq -1$.
" In the context of treating Cardassian models as a fluid (Gondolo&Freese2002).. this is w=py 1. which satisfies w>—1 for the parameter choices we are interested in: Q""5<Lon «]. and q>1."," In the context of treating Cardassian models as a fluid \citep{gondolo}, this is w = = -, which satisfies $w \geq -1$ for the parameter choices we are interested in: $\Omega_m^{obs} <1$, $n<1$, and $q>1$."
 Therelore. all viable MP Cardassian moclels satis[v the weak energy. condition but can have wy<—1.," Therefore, all viable MP Cardassian models satisfy the weak energy condition but can have $w_X < -1$."
 Note an effective ix«—] is consistei with recent CMD and Large scale structure data (Schueckeretal.2002:Melchiorri2002).," Note an effective $w_X <-1$ is consistent with recent CMB and large scale structure data \citep{Schue02,Mel02}."
. some scalar field dark οποιον models with wy<—1 iwe been studied. previously (Caldwell1999):: models which are stable despite violating variants of the weak energy conditions have been found to be diffieult to construct (Carrol.Hoffman.&Trodden2003)., Some scalar field dark energy models with $w_X < -1$ have been studied previously \citep{caldwell}; models which are stable despite violating variants of the weak energy conditions have been found to be difficult to construct \citep{Carroll03}.
. Our proposal here is a different alternative to the models previously studied., Our proposal here is a different alternative to the models previously studied.
 We will now assume that the data set is given. either [rom the simulated. data sets described above. or. in the future fromSNAP*.," We will now assume that the data set is given, either from the simulated data sets described above, or, in the future from."
. We will show that investigation of the data set can reproduce the sien of the time-dependence of the dark energy density. assuming one knows the matter density (o an accuracy of1054.," We will show that investigation of the data set can reproduce the sign of the time-dependence of the dark energy density, assuming one knows the matter density to an accuracy of."
.. Given our data set. we now proceed as though we have no information on where it comes from: Le. we proceed as though we did not know which model it came from.," Given our data set, we now proceed as though we have no information on where it comes from; i.e., we proceed as though we did not know which model it came from."
 We parametrize the dark energy density in order to allow us to compare it to the data set., We parametrize the dark energy density in order to allow us to compare it to the data set.
 We take px(z) to be an arbitrary function., We take $\rho_X(z)$ to be an arbitrary function.
" To approximate the function. we parametrize it by ils value at ijj, equally spaced redshilt values. z;. ο Ἡμτῃ=tmer"," To approximate the function, we parametrize it by its value at $n_{bin}$ equally spaced redshift values, $z_i$, $i=1$ $n_{bin}$, $z_{n_{bin}}=z_{max}$."
 Lhe value of px(z) al other redshifts are given by linear interpolation. ie. PealTN (---- Pj tems QE," The value of $\rho_X(z)$ at other redshifts are given by linear interpolation, i.e., ( ( _i, 1cm < z z_i,"
exhibit no significant fhictuations within a timescale of months to vears. while the behavior ol the X-ray continuum is much more active within much shorter timescales (IXrolik οἱ al.,"exhibit no significant fluctuations within a timescale of months to years, while the behavior of the X-ray continuum is much more active within much shorter timescales (Krolik et al."
 1993)., 1993).
 As a result. the observed ecuivalent width of the narrow line would be surely smaller while (he X-ray continuum in a higher state. ancl vise versa.," As a result, the observed equivalent width of the narrow line would be surely smaller while the X-ray continuum in a higher state, and vise versa."
 For instance. Figure 6 plots the observations of and with (the only two radio quiet sources with more than two LETC observations) and the best-fit lines.," For instance, Figure 6 plots the observations of and with (the only two radio quiet sources with more than two HETG observations) and the best-fit lines."
 Although there are onlv 6 points lor NGC 3783 and 3 points for NGC 4151. we performed a linear regression to give the best-fit slopes of them (a=—L.775240.8130. Rs = —0.657 for NGC 3783 and v=—1.524140.0732. Rs = —1.000 lor NGC 4151).," Although there are only 6 points for NGC 3783 and 3 points for NGC 4151, we performed a linear regression to give the best-fit slopes of them $\alpha=-1.7752
\pm0.8130$, Rs = $-0.657$ for NGC 3783 and $\alpha=-1.5244\pm0.0732$ Rs = $-1.000$ for NGC 4151)."
 As expected. its EW clearly decreases with increasing luninositv.," As expected, its EW clearly decreases with increasing luminosity."
 Such an effect may naturally lead to the observed anti-correlation slope between the EW and luminosity [or a sample of AGN., Such an effect may naturally lead to the observed anti-correlation slope between the EW and luminosity for a sample of AGN.
 In this section we run simulations to check this possibility., In this section we run Monte-Carlo simulations to check this possibility.
 An anti-correlation between the X-ray luminosity aud long term variation amplitude in Sevferl 1 galaxies was described by Markovitz Edelson (2004)., An anti-correlation between the X-ray luminosity and long term variation amplitude in Seyfert 1 galaxies was described by Markovitz Edelson (2004).
" Fractional variability amplitudes (£,,.,) were measured for each light curve to «quantify the intrinsic variability anmiplitucde.", Fractional variability amplitudes $F_{var}$ ) were measured for each light curve to quantify the intrinsic variability amplitude.
 where 5? is the total variance of the light curve. (στι)CrP> is the mean error squared and (X) is (he mean count rate of N total points.," where $S^2$ is the total variance of the light curve, $\langle
\sigma_{err}^2 \rangle$ is the mean error squared and $ \langle X \rangle
$ is the mean count rate of $N$ total points."
" They eave the anti-correlation between Fractional variability amplitude and Iuminositv: FL,xL,""0.1: for long (1296 days) timescale data.", They gave the anti-correlation between Fractional variability amplitude and luminosity: $F_{var} \propto L^{-0.135}_x$ for long (1296 days) timescale data.
 llere we adopt a tov model to simulate (he X-ray continuum. variation. while line flux is assumed. to be invariable.," Here we adopt a toy model to simulate the X-ray continuum variation, while line flux is assumed to be invariable."
 The observed X-ray luminosity is assumed. to be normally clistvibuted with the width of the Gaussian distribution calculated to match the observed excess variance al different luminosities (Markovitz Eclelson 2004)., The observed X-ray luminosity is assumed to be normally distributed with the width of the Gaussian distribution calculated to match the observed excess variance at different luminosities (Markovitz Edelson 2004).
 By normalizing the observed line EW for the combined RQ sample (with upper limits) to the best-fit line. we first construct a set of line EW which does not correlate with the A-ray luminosity.," By normalizing the observed line EW for the combined RQ sample (with upper limits) to the best-fit line, we first construct a set of line EW which does not correlate with the X-ray luminosity."
 Random conünuum variations are (hen added to the hunmüinositles. and the line EW are mocdified Corresponcinely since we assume no change to the line flix.," Random continuum variations are then added to the luminosities, and the line EW are modified correspondingly since we assume no change to the line flux."
 We repeated (his step to build 1000 artificial samples with different random seed for each time., We repeated this step to build 1000 artificial samples with different random seed for each time.
 We used ASURV {ο perform linear regression to the artificial samples., We used ASURV to perform linear regression to the artificial samples.
 The distribution of the best-fit slopes of the artificial samples was presented in Figure 7., The distribution of the best-fit power-law slopes of the artificial samples was presented in Figure 7.
 We can see that of the simulations produce anti-correlation slopes steeper (han the observed value. and the mean value is —0.04852E0.0536.which is compatible with the observed value (a=—0.10192:0.0524) within the errors.," We can see that of the simulations produce anti-correlation slopes steeper than the observed value, and the mean value is $-0.0485\pm0.0536$ ,which is compatible with the observed value $\alpha=-0.1019\pm0.0524$ ) within the errors."
we stack the cata from the eight halos together and then separate the satellite population into dillering mass bins we see little variation.,we stack the data from the eight halos together and then separate the satellite population into differing mass bins we see little variation.
 “Vhis common behaviour still. holes and. probably reflects the scale free force. of gravity. and hierarchical construction., This common behaviour still holds and probably reflects the scale free force of gravity and hierarchical construction.
 The disrupted satellites. however. have a dillerent distribution. almost mirroring the survivors about the e= O.5-axes. with the peak ceccentricitynear ey)d=0.34⋅ and a dispersion.. of Dui65?=0.16.," The disrupted satellites, however, have a different distribution, almost mirroring the survivors about the $e=0.5$ -axes, with the peak eccentricitynear $\langle{e^d_0}\rangle=0.34$ and a dispersion of $\langle{\sigma^d}\rangle=0.16$."
 .Thus. the destroved satellites were preferentially on. more. circular orbits.," Thus, the destroyed satellites were preferentially on more circular orbits."
 The pericentre clistributions provide acdcditional insight into the nature of the orbits of the disrupted satellites., The pericentre distributions provide additional insight into the nature of the orbits of the disrupted satellites.
 The result for redshift >=0 can be seen in where p has been normalised by the virial radius of the host., The result for redshift $z=0$ can be seen in where $p$ has been normalised by the virial radius of the host.
 Again. live satellites are represented as thick lines and disrupted ones as thin lines.," Again, live satellites are represented as thick lines and disrupted ones as thin lines."
 A striking characteristic of these distributions is again the similarity between the halos., A striking characteristic of these distributions is again the similarity between the halos.
 Moreover. we also observe a similarity in the distribution for the live and disrupted satellites.," Moreover, we also observe a similarity in the distribution for the live and disrupted satellites."
 This is emphasised. particularly by the best-fit Gaussians to the distributions., This is emphasised particularly by the best-fit Gaussians to the distributions.
 Phe maximum peak lies at of the virial radius for live satellites with a mean dispersion of (65=0.12 as opposed to of [for⋅⋅ the disrupted. ones with. a dispersion. of⋅ ἐσd=0.11., The maximum peak lies at of the virial radius for live satellites with a mean dispersion of $\langle\sigma\rangle=0.12$ as opposed to of for the disrupted ones with a dispersion of $\langle\sigma^d\rangle=0.11$.
 Unlike cecentricity. the pericentre distribution rises quickly towards the peak and falls olf moderately to the outer parts of the host halo.," Unlike eccentricity, the pericentre distribution rises quickly towards the peak and falls off moderately to the outer parts of the host halo."
 We already. noted. the lack of correlation with mass. age. environment and richness. but there is. however. à mild. dependence on the state (live or. dead) of the satellite as disrupted satellites appear to have had marginally nearer excursions towards the host centre.," We already noted the lack of correlation with mass, age, environment and richness, but there is, however, a mild dependence on the state (live or dead) of the satellite as disrupted satellites appear to have had marginally nearer excursions towards the host centre."
 When we stack the data from the eight halos ancl separate into dillering satellite mass bins we again see little variation., When we stack the data from the eight halos and separate into differing satellite mass bins we again see little variation.
 In summary the cillerence between live ancl disrupted satellites lies primarily in the eccentricity distribution., In summary the difference between live and disrupted satellites lies primarily in the eccentricity distribution.
 Disrupted satellites seem to be on more circular orbits., Disrupted satellites seem to be on more circular orbits.
 Since he disrupted: satellites have similar pericentres to. those which survive their circular nature means that they. spend more time in the deeper regions of the potential well., Since the disrupted satellites have similar pericentres to those which survive their circular nature means that they spend more time in the deeper regions of the potential well.
 Llence. hey experience stronger tidal forces Lor longer periods. and are thus disrupteddissolved more reaclily.," Hence, they experience stronger tidal forces for longer periods, and are thus disrupted/dissolved more readily."
 We close this subsection with an examination of the »ericentre-eccentricity relation. 9))., We close this subsection with an examination of the pericentre-eccentricity relation ).
 The crosses represent the satellites that survived. until redshift 2.= ). while the diamonds. represent. the disrupted: satellites.," The crosses represent the satellites that survived until redshift $z=0$ , while the diamonds represent the disrupted satellites."
vvalues than the high-a stars. there is considerable scatter inY/Fe|: several of the thick-disk and high-e stars fall among the low-alpha stars.,"values than the $\alpha$ stars, there is considerable scatter in; several of the thick-disk and $\alpha$ stars fall among the low-alpha stars."
 In addition. it should be noted that one star. 24-25.. has [Y/Fe]=0.82 and falls above the upper limit of the figure.," In addition, it should be noted that one star, , has $\yfe = 0.82$ and falls above the upper limit of the figure."
 This star is also very overabundant in Ba and other s-process elements with an abundance pattern that may be explained by mass transfer from an AGB companion (L. Shu et al..," This star is also very overabundant in Ba and other $s$ -process elements with an abundance pattern that may be explained by mass transfer from an AGB companion (L. Shu et al.,"
 in preparation)., in preparation).
 According to Latham et al. (2002)).," According to Latham et al. \cite{latham02}) ),"
 is an SBI spectroscopic binary with an orbital period of yyears., is an SB1 spectroscopic binary with an orbital period of years.
" Abundances of Ba are determined from the ,,05853.7.6141.7 llines."," Abundances of Ba are determined from the $\lambda \lambda \, 
5853.7, \, 6141.7$ lines."
 Odd numbered Ba tsotopes exhibit hyperfine splitting. and it was therefore first investigated if this has any significant effect on the derived abundances.," Odd numbered Ba isotopes exhibit hyperfine splitting, and it was therefore first investigated if this has any significant effect on the derived abundances."
 Adopting HFS data from MeWilliam (1998)). profiles and equivalent widths were calculated for some representative stars.," Adopting HFS data from McWilliam \cite{mcwilliam98}) ), profiles and equivalent widths were calculated for some representative stars."
 Contrary to the case of the Hine at ((e.g. Collet et al. 2009)).," Contrary to the case of the line at (e.g. Collet et al. \cite{collet09}) ),"
" the HFS effect on the j,05853.7.6141.7 lines is small and changes the derived Ba abundances by less than ddex even if one assumes a pure r-process distribution of the isotopes instead of a solar system distribution."," the HFS effect on the $\lambda \lambda \, 5853.7, \, 6141.7$ lines is small and changes the derived Ba abundances by less than dex even if one assumes a pure $r$ -process distribution of the isotopes instead of a solar system distribution."
 Thus. the abundance analysis was carried out neglecting hypertine splitting.," Thus, the abundance analysis was carried out neglecting hyperfine splitting."
 The (5853.7 lline is suitable for precise abundance determination by having equivalent widths from about 10 to iin the majority of our stars.," The $\lambda \, 5853.7$ line is suitable for precise abundance determination by having equivalent widths from about 10 to in the majority of our stars."
 The (6141.7 line has equivalent widths typically in the range 50 -mma. and is therefore more sensitive to the adopted microturbulence and damping constant.," The $\lambda \, 6141.7$ line has equivalent widths typically in the range 50 -, and is therefore more sensitive to the adopted microturbulence and damping constant."
 Nevertheless. the two sets of Ba abundances show a satisfactory agreement: the rms difference is ddex.," Nevertheless, the two sets of Ba abundances show a satisfactory agreement; the rms difference is dex."
 Fig., Fig.
 3 (middle panel) shows vversus[Fe/H]., \ref{fig:y.ba.bay-fe} (middle panel) shows versus.
 Like in the case ofY/Fe]. there is some mixing of high- and low-« stars.," Like in the case of, there is some mixing of high- and $\alpha$ stars."
 The s-process rich star.24-25.. has [Ba/Fe]=1.45 and falls far above the upper limit of the figure.," The $s$ -process rich star, has $\bafe = 1.45$ and falls far above the upper limit of the figure."
 The same is the case with96-20.. a mildly s-process rich star. which has |[Ba/Fe]=0.68 and |Y/Fe|=0.28.," The same is the case with, a mildly $s$ -process rich star, which has $\bafe = 0.68$ and $\yfe = 0.28$."
 Interestingly. there is a correlation between the scatter in aand aat a given metallicity.," Interestingly, there is a correlation between the scatter in and at a given metallicity."
 This is evident from the lower panel of Fig. 3..," This is evident from the lower panel of Fig. \ref{fig:y.ba.bay-fe},"
 where iis plotted vs.[Fe/H]. The high- and low-« stars show well-defined trends of wwith increasing separation as a function of increasing[Fe/H]., where is plotted vs. The high- and $\alpha$ stars show well-defined trends of with increasing separation as a function of increasing.
Still. a few stars deviatevery significantly from these trends.,"Still, a few stars deviatevery significantly from these trends."
 with |Ba/Y=0.64 falls above the limit of the figure. whereas the other s-process rich star.96-20.. has [Ba/Y]20.40.," with $\bay = 0.64$ falls above the limit of the figure, whereas the other $s$ -process rich star, has $\bay = 0.40$."
 Other deviating stars are (Na-rich according to NS10) and6810., Other deviating stars are (Na-rich according to NS10) and.
. An estimate of the statistical errors of the various abundance ratios can be obtained by comparing values derived from UVES and FIES spectra for six stars observed with both instruments., An estimate of the statistical errors of the various abundance ratios can be obtained by comparing values derived from UVES and FIES spectra for six stars observed with both instruments.
 As the atmosphericparameters are determined spectroscopically. the comparison will include the effect of the uncertainties in Ty. logg. and oon the abundance ratios.," As the atmosphericparameters are determined spectroscopically, the comparison will include the effect of the uncertainties in , , and on the abundance ratios."
Alore detailed. discussion on the CALCD model can be ound in Paper LH. in which we have applied the same model for a spherical. geometry to six. ULXs. observed withNALALNewtou.,"More detailed discussion on the CMCD model can be found in Paper II, in which we have applied the same model for a spherical geometry to six ULXs observed with."
" The fitted. ;, ( 0.050.3. keV) of hese sources are cdistinctlv dillerent from the values (1 keV) obtained. for known stellar-mass DlIs. as presented in this paper and in Paper L Indeed. the inferred. DII masses (Mpgg) of the ULXs are ~LOCAL... consistent with he intermeciate-mass DIE interpretation of these sources."," The fitted $T_{in}$ $\sim$ 0.05–0.3 keV) of these sources are distinctly different from the values $\sim 1$ keV) obtained for known stellar-mass BHs, as presented in this paper and in Paper I. Indeed, the inferred BH masses $M_{BH}$ ) of the ULXs are $\sim 10^3 M_\odot$, consistent with the intermediate-mass BH interpretation of these sources."
 We have also shown that the AICD|PL model gives an equivalent spectral description of the ULAs. although the CALCD moclel provides unique constrains on the corona properties and on the disk inclination angles. as well as on the DII masses.," We have also shown that the MCD+PL model gives an equivalent spectral description of the ULXs, although the CMCD model provides unique constrains on the corona properties and on the disk inclination angles, as well as on the BH masses."
 Because of the lower disk. temperatures. compared to those of the stellar mass DII systems. the nonphysical elfects of the MCD|PL model are. typically not significant in the observable photon energy range of the intermediate-mass DII candidates.," Because of the lower disk temperatures, compared to those of the stellar mass BH systems, the nonphysical effects of the MCD+PL model are typically not significant in the observable photon energy range of the intermediate-mass BH candidates."
 In the present work. we conduct a critical test. of the CAICD and. ACD|PL models by comparing parameters (Mew. 0. and the equivalent neutral hydrogen absorption Nay) inferred from the N-ray. spectra of LMC X.1 and X with the more direct. measurements based on optical and dispersed. X-ray. spectra: we also compare the results from the dilferent corona geometrical configurations of the CALCD model.," In the present work, we conduct a critical test of the CMCD and MCD+PL models by comparing parameters $M_{BH}$, $\theta$ , and the equivalent neutral hydrogen absorption $N_H$ ) inferred from the X-ray spectra of LMC X–1 and X--3 with the more direct measurements based on optical and dispersed X-ray spectra; we also compare the results from the different corona geometrical configurations of the CMCD model."
 We first. briclly describe these measurements. and the X-ray. observations in £2. and then present the spectral fitting results in 83.," We first briefly describe these measurements and the X-ray observations in 2, and then present the spectral fitting results in 3."
 We describe the specific comparisons in 84 and. present the discussion and our conclusions in 5., We describe the specific comparisons in 4 and present the discussion and our conclusions in 5.
 We select LMC X.1 and X3 for this studs chielly because of their location in our nearest neighboring galaxy. the Large Magellanic Cloud (LMC: 2=50 kpe is adopted throughout the work).," We select LMC X–1 and X–3 for this study chiefly because of their location in our nearest neighboring galaxy, the Large Magellanic Cloud (LMC; $D = 50$ kpc is adopted throughout the work)."
 Both the well-determined distance and the relatively low foreground soft. X-ray absorption are essential to our test., Both the well-determined distance and the relatively low foreground soft X-ray absorption are essential to our test.
 These two sources are also among the three well-known persistent. DIENDs and are usually. found in the high/soft state., These two sources are also among the three well-known persistent BHXBs and are usually found in the high/soft state.
 The lowhard state was occasionally reported. for LMC X3 (Bove&Smale2000:Homanetal. 2000).. but never for LAIC X.1.," The low/hard state was occasionally reported for LMC X–3 \citep{boy00, hom00}, but never for LMC X–1."
 The remaining known persistent. DIIND. Cvenus X.1. is in our Galaxy. ancl stays mostly in the lowhard state (e.g. Pottschmidt 2003).," The remaining known persistent BHXB, Cygnus X–1, is in our Galaxy and stays mostly in the low/hard state (e.g., Pottschmidt 2003)."
 The X-ray. spectra of this source also show a strong clisk reflection. component (Cilfanovetal.L999:Frontera.al.2001) a complication that is not included in the CALCD models.," The X-ray spectra of this source also show a strong disk reflection component \citep{gil99, fro01} — a complication that is not included in the CMCD models."
 Fable 1. summarizes the kev parameters of LMC X. Land X.3. which are used for the comparison with our spectrally inferred. values (84).," Table \ref{tab:keymeasure} summarizes the key parameters of LMC X–1 and X–3, which are used for the comparison with our spectrally inferred values 4)."
 Both LMC X.1 and X.3 were observed with (ασ. Cui 2002) and (c.g... Page 2003).," Both LMC X–1 and X–3 were observed with (e.g., Cui 2002) and (e.g., Page 2003)."
 The dispersed: X-ray spectra of LAIC X.3 have been used to measure the N-ray. absorption edges (mainly for Oxvgen). which tightly constrains the the absorbing matter column density Wy along the line of sight (Pageetal. 2003).," The dispersed X-ray spectra of LMC X–3 have been used to measure the X-ray absorption edges (mainly for Oxygen), which tightly constrains the the absorbing matter column density $N_H$ along the line of sight \citep{pag03}."
. The absorption towards LMC X.LE is. however. substantially higher.," The absorption towards LMC X–1 is, however, substantially higher."
 As a result. the photon flux at. the Oxveen edge is too low to allow for a useful constraint on [Ng based on the existing data.," As a result, the photon flux at the Oxygen edge is too low to allow for a useful constraint on $N_H$ based on the existing data."
 We here. utilize the cata from the observations. which were carried out on LOOT October 5 [or LMC X.Lane October 11 for XN3 ClErevesetal.2000).," We here utilize the data from the observations, which were carried out on 1997 October 5 for LMC X–1 and October 11 for X–3 \citep{tre00}."
.. The data do not have pile-up problems. which could be present for N-rav CCD imaging observations of bright. sources.," The data do not have pile-up problems, which could be present for X-ray CCD imaging observations of bright sources."
 Four types of narrow-field. instruments (NEIs). were on board: Low Energy. Concentrator System (LECS). Medium ποιον Concentrator Systems (AIECS). Leh Pressure Gas Scintillation Proportional counter (IIPGSPC). and Phoswich Detector System (PDS) (Boellaetal.1997).," Four types of narrow-field instruments (NFIs) were on board: Low Energy Concentrator System (LECS), Medium Energy Concentrator Systems (MECS), High Pressure Gas Scintillation Proportional counter (HPGSPC), and Phoswich Detector System (PDS) \citep{boella97}."
. The exposure for LMC X.1 and X3 were about 15. ks each for the LECS and about 40 ks each for the MECS., The exposure for LMC X–1 and X–3 were about 15 ks each for the LECS and about 40 ks each for the MECS.
 ‘These two instruments were sensitive to X-rays in the energy ranges of 0.1.10 and 1.3.10 keV respectively., These two instruments were sensitive to X-rays in the energy ranges of 0.1–10 and 1.3–10 keV respectively.
 Data from the IIPGSIC and the PDS. which were sensitive to. photons in 4120 keV and 15300 keV ranges respectively. were not included. because of poor counting statistics. ancl also because of possible source contamination from PSL. 69 (which is 25/ away from LMC X.1) (Sewardetal.1984:Llaareltetal. 2001).," Data from the HPGSPC and the PDS, which were sensitive to photons in 4–120 keV and 15–300 keV ranges respectively, were not included because of poor counting statistics, and also because of possible source contamination from PSR 0540--69 (which is $^\prime$ away from LMC X–1) \citep{sew84,haa01}."
. We extracted the spectra from a radius of S and SA. around. cach source from the LECS and the ALECS observations and used the energy ranges of 0.24 keV and 1.810 keV for these two instruments in this study., We extracted the spectra from a radius of $'$ and $'$ around each source from the LECS and the MECS observations and used the energy ranges of 0.2–4 keV and 1.8–10 keV for these two instruments in this study.
 The background. contributions to the spectra are small anc are estimated from a blank field., The background contributions to the spectra are small and are estimated from a blank field.
 Phe spectra from the LECS. the MIECS2. ancl the MECS3 were jointly fitted. for each source. using the software package055.," The spectra from the LECS, the MECS2, and the MECS3 were jointly fitted for each source, using the software package."
 We summarize the spectral fitting results and the inferred source fluxes in Table 2.., We summarize the spectral fitting results and the inferred source fluxes in Table \ref{tab:fit-parameters}.
 Phe quoted uncertainty ranges of the parameters are all at confidence level., The quoted uncertainty ranges of the parameters are all at confidence level.
 Pig., Fig.
 1 shows the spectral fits with the CALCD models., \ref{fig:cmcd} shows the spectral fits with the CMCD models.
 Phe systematic deviation of the data from the model at low energies (=1 keV) might be due to poor calibration of the instrument spectral response (Martinetal., The systematic deviation of the data from the model at low energies $\la 1$ keV) might be due to poor calibration of the instrument spectral response \citep{martin96}. .
 199, Fig.
G6).. Fig. 2 illustrates the elfects of the Comptonization in the spherical corona Ssyslens., \ref{fig:compton} illustrates the effects of the Comptonization in the spherical corona systems.
 Both CALCD and MCD|PL mocdels give acceptable fits., Both CMCD and MCD+PL models give acceptable fits.
 The model parameters are all well constrained. except for the CALCD parameter 6. for which only the upper or lower limit is constrained.," The model parameters are all well constrained except for the CMCD parameter $\theta$, for which only the upper or lower limit is constrained."
 Phe MCD|PL model parameters we obtained here are consistent with those reported by Llaareltetal. (2001)., The MCD+PL model parameters we obtained here are consistent with those reported by \citet{haa01}.
. The fitted. Ny values from MCD|PL are systematically higher than those from CALCDs., The fitted $N_H$ values from MCD+PL are systematically higher than those from CMCDs.
 The same is true for the inferred absorption-corrected Duxes. especially for LMC X.1 which is a factor of ~ 4 higher from MCD|PL than from CAICD (Yable 2)).," The same is true for the inferred absorption-corrected fluxes, especially for LMC X–1 which is a factor of $\sim$ 4 higher from MCD+PL than from CMCD (Table \ref{tab:fit-parameters}) )."
" For LAIC X3. £i, from AICD|PL is slightly higher than those from CMCDs. whereas the value o£ Avepbfcos(?) is consistent with those of {νοο."," For LMC X–3, $T_{in}$ from MCD+PL is slightly higher than those from CMCDs, whereas the value of $K_{MCD}/{\rm cos}(\theta)$ is consistent with those of $K_{CMCD}$."
 Ehe best fit parameters in the two different &cometric CMCDs are nearly identical., The best fit parameters in the two different geometric CMCDs are nearly identical.
 The small value of τι indicates that only a small portion of disk photons have been up-scatterec to high energies., The small value of $\tau_c$ indicates that only a small portion of disk photons have been up-scattered to high energies.
 For LAIC X.1. except for à consistent Ny value. the fitted parameters are significantly differentbetween the two cillerent geometric CALCDs (Table 2)).," For LMC X–1, except for a consistent $N_H$ value, the fitted parameters are significantly differentbetween the two different geometric CMCDs (Table \ref{tab:fit-parameters}))."
 We will see in £44. the results from the slab-like configuration are inconsistent with the independent. measurements.," We will see in 4, the results from the slab-like configuration are inconsistent with the independent measurements."
" Z;, values from the spherical CMCD and the MCD|PL are consistent with cach", $T_{in}$ values from the spherical CMCD and the MCD+PL are consistent with each
of the shell shape.,of the shell shape.
 Mass ancl energy loss from a star forming within the compressed gas ofan existing shell may blow open a secondary wind shell in the primary shell wall. leading to deformation of the primary shell. ancl ultimately. à departure from the signature shell shape.," Mass and energy loss from a star forming within the compressed gas of an existing shell may blow open a secondary wind shell in the primary shell wall, leading to deformation of the primary shell, and ultimately, a departure from the signature shell shape."
 A shell expanding into a low ambient density region. will not accumulate a high density rim as quickly as one embedded: in a higher density., A shell expanding into a low ambient density region will not accumulate a high density rim as quickly as one embedded in a higher density.
 As star formation usually occurs only after a threshold column density is reached. we might not expect shells that are expanding into a low ambient density medium to be as reaclily disturbed. by secondary. star formation.," As star formation usually occurs only after a threshold column density is reached, we might not expect shells that are expanding into a low ambient density medium to be as readily disturbed by secondary star formation."
 Secondary star formation within the shell wall has been. observed. in the SAIC (Stanimirovic.. 1999). while small sshells clustered within the wall of a larger sshell have also been observed in the ELMC (sim et al.," Secondary star formation within the shell wall has been observed in the SMC (Stanimirović,, 1999), while small shells clustered within the wall of a larger shell have also been observed in the LMC (Kim et al."
 1999)., 1999).
 Shell-like features. such as blow-outs. or chimnevs. that were not included in this catalogue. were found occasionally throughout the cube.," Shell-like features, such as blow-outs, or chimneys, that were not included in this catalogue, were found occasionally throughout the cube."
 X blow-out. or a tunnel can develop bv an expanding shell forming close to a region of much lower relative density.," A blow-out, or a tunnel can develop by an expanding shell forming close to a region of much lower relative density."
 The expanding eas can blow through the boundary separating the two densities. such as through the wall of a gas cloud. and into the low density. region.," The expanding gas can blow through the boundary separating the two densities, such as through the wall of a gas cloud, and into the low density region."
 Such structures can also form through the merging of two expanding shells. and have been observed in the Galaxy (eg MeClure-Crilliths et al. 2000). as well as other galaxies (eg.," Such structures can also form through the merging of two expanding shells, and have been observed in the Galaxy (eg McClure-Griffiths et al, 2000), as well as other galaxies (eg."
 Ott et aL.," Ott et al.,"
 2001)., 2001).
 Under these conditions. the calculation of the dynamic age. which is based on an assumption of constant and homogeneous ambient gas density. is incorrect.," Under these conditions, the calculation of the dynamic age, which is based on an assumption of constant and homogeneous ambient gas density, is incorrect."
 AX study of these shell-like structures will be included in a [uture project., A study of these shell-like structures will be included in a future project.
 Any constraints on shell radii imposed. by the extent. of the eas in the Magellanic Bridge are not considered to be significant: the height of the high dedensitv region in the Bridge. in. Declination. is almost four times the diameter of the largest shell found {rom this survey. although it is of 10 same order of the cliameter of the largest. supershell in in the SAIC (Stanimirovic.. 1999).," Any constraints on shell radii imposed by the extent of the gas in the Magellanic Bridge are not considered to be significant: the height of the high density region in the Bridge, in Declination, is almost four times the diameter of the largest shell found from this survey, although it is of the same order of the diameter of the largest supershell in in the SMC (Stanimirović,, 1999)."
 Phe largest shell radius found during this survey was11.77... equivalent to 72204 pe. and the radius of the largest supershell found in the SAIC was ~9910 pe.," The largest shell radius found during this survey was, equivalent to 204 pc, and the radius of the largest supershell found in the SMC was 910 pc."
 Given the mechanism olf formation of the Bridge. one possible mechanism. of deformation. 1s. tical stretching.," Given the mechanism of formation of the Bridge, one possible mechanism of deformation is tidal stretching."
 Llowever. of the shells surveved. there does not appear to be," However, of the shells surveyed, there does not appear to be"
During the past dozen years. a substantial number of circumstellar accretion disks around young stellar objects have been with a number of different techniques such as direct optical/infrared imaging. (sub-)mm or radio-mapping. radio-interferometry. and (more recently) infrared interferometry (see.e.g..Stapelfeldtetal.1998;Zinnecker1999:Wolfetal. 2010)..,"During the past dozen years, a substantial number of circumstellar accretion disks around young stellar objects have been with a number of different techniques such as direct optical/infrared imaging, (sub-)mm or radio-mapping, radio-interferometry, and (more recently) infrared interferometry \citep[see, e.g.,][]{Stapelfeldt98,Zinnecker99,Wolf03,Leinert04,Preibisch06,Kraus08,Ratzka09,Gramajo10,Quanz10}."
 Spectroscopic observations have also provided important kinematical evidence of accretion disks (see.e.g..Daviesetal. 2010).," Spectroscopic observations have also provided important kinematical evidence of accretion disks \citep[see, e.g.,][]{Chandler95,Bik04,Wheelwright10,Davies10}."
. These studies have greatly advanced our understanding of the physics of young stellar objects (YSOs). how the young stars gain their final mass. how the circumstellar matter. in which planets may be forming. evolves with time and is finally dispersed (seeDullemondetal.2007).," These studies have greatly advanced our understanding of the physics of young stellar objects (YSOs), how the young stars gain their final mass, how the circumstellar matter, in which planets may be forming, evolves with time and is finally dispersed \citep[see][]{Dullemond07}."
. So far. nearly all well characterized circumstellar disks surround low- or intermediate mass stellar objects. with less than about 5—8M..," So far, nearly all well characterized circumstellar disks surround low- or intermediate mass stellar objects, with less than about $5-8\,M_\odot$."
 Only a very limited number of disk detections have been reported around objects with higher masses in the range M.~8—20 (see.e.g..Chinietal.2010:Daviesetal. 2010).," Only a very limited number of disk detections have been reported around objects with higher masses in the range $M \sim 8 - 20\,M_\odot$ \citep[see, e.g.,][]{Chini04,Jiang05,Nielbock07,Kraus10,Davies10}."
. However. in many of these cases the mass estimates for the central star are quite uncertain; the true mass may in fact be considerably smaller than initially estimated (see.e.g..Sakoetal.2005).," However, in many of these cases the mass estimates for the central star are quite uncertain; the true mass may in fact be considerably smaller than initially estimated \citep[see, e.g.,][]{Sako05}."
. Another problem is that the observed flattened structure around these more massive young stellar objects are not necessarily accretion disks. from which material is directly transferred to the central star: they may rather represent a thick torus that could be the remnant of a flattened envelope (Allenetal.2003).," Another problem is that the observed flattened structure around these more massive young stellar objects are not necessarily accretion disks, from which material is directly transferred to the central star; they may rather represent a thick torus that could be the remnant of a flattened envelope \citep{Allen03}."
. Considering even higher stellar masses. M220M... no clear observational evidence for aceretion disks has yet been found (see.e.g..Cesarontetal. 2007).," Considering even higher stellar masses, $M \ga 20\,M_\odot$, no clear observational evidence for accretion disks has yet been found \citep[see, e.g.,][]{Cesaroni07}."
. This rarity and lack of convineing disk detections ts an important aspect in the long-standing discussion about how massive stars form., This rarity and lack of convincing disk detections is an important aspect in the long-standing discussion about how massive stars form.
 A fundamental problem ts that the strong radiation pressure resulting from the very high luminosity of à massive protostars tends to halt the accretion flow and should thus limit the final stellar mass. at least i the case of spherical accretion (Kahn 1974)..," A fundamental problem is that the strong radiation pressure resulting from the very high luminosity of a massive protostars tends to halt the accretion flow and should thus limit the final stellar mass, at least in the case of spherical accretion \citep{Kahn74}. ."
" More recent theoretical results alleviate this problem by mechanisms such às spherical accretion. the “flash-light effect”. and ""photon-bubbles” (see.e.g..Yorke&Sonnhalter2002:Krumholzetal. 2009)."," More recent theoretical results alleviate this problem by mechanisms such as non-spherical accretion, the “flash-light effect”, and ``photon-bubbles'' \citep[see, e.g.,][]{Yorke02,Krumholz09}."
. The more accurate treatment of frequency-dependent radiative feedback by Kuiperetal.(2010) suggests that stars with masses well beyond the upper mass limit of spherical aceretion can be formed by accretion.," The more accurate treatment of frequency-dependent radiative feedback by \citet{Kuiper10}
 suggests that stars with masses well beyond the upper mass limit of spherical accretion can be formed by accretion."
 However. due to the limited numerical resolution of these simulations (e.g..1.27AUinthestudyofKuiperetal. 2010).. the reliability of the numerical results 1s not fully clear.," However, due to the limited numerical resolution of these simulations \citep[e.g., 1.27~AU in the study of][]{Kuiper10}, the reliability of the numerical results is not fully clear."
 As discussed in Zinnecker&Yorke (2007).. there are still strong indications that the formation of massive stars is simply a sealed-up version of the low-mass star formation process.," As discussed in \citet{Zinnecker07}, there are still strong indications that the formation of massive stars is simply a scaled-up version of the low-mass star formation process."
 Alternative models for the formation of massive stars highlight the importance of the cluster environment during early star formation stages (see.e.g..Bonnelletal.2004:Krumholz 2010).. and processes such as protostellar collisions and mergers may be possible in the central regions of very dense clusters (see.e.g..Daviesetal. 2006).," Alternative models for the formation of massive stars highlight the importance of the cluster environment during early star formation stages \citep[see, e.g.,][]{Bonnell04,Krumholz10}, and processes such as protostellar collisions and mergers may be possible in the central regions of very dense clusters \citep[see, e.g.,][]{Bonnell98,Bonnell05,Bally05,Davies06}."
. In the context of these open questions. every new detection and good characterization of a circumstellar disk around a massive (M.2 8M.) young stellar object provides an importantnew mosaic stone that helps to solve the puzzle," In the context of these open questions, every new detection and good characterization of a circumstellar disk around a massive $M \ga 8\,M_\odot$ ) young stellar object provides an importantnew mosaic stone that helps to solve the puzzle"
juchides sharp-divided tages obtained by dividing the observed mniases by their filtered counterparts. difference nuages. obtained by fitting ellipses to the isophotes aud subtracting such imnodols to the observed images. aud colour JN iniges.,"includes sharp-divided images obtained by dividing the observed images by their filtered counterparts, difference images, obtained by fitting ellipses to the isophotes and subtracting such models to the observed images, and colour J/K' images."
 A bulee|disk model was fit to the nuage profiles. and the correspouding fit owanuieters are elven.," A bulge+disk model was fit to the image profiles, and the corresponding fit parameters are given."
 Four (one) out of five (one) of the optically classified uon-biurred active (control) galaxies resul to harbour a bar., Four (one) out of five (one) of the optically classified non-barred active (control) galaxies result to harbour a bar.
 Three of them had already been described as barred ealaxies. as derived either from NIR (UCC 1395. which also has a secondary bar. aud NGC 6890) or optical I (NGC 3281) analyses.," Three of them had already been described as barred galaxies, as derived either from NIR (UGC 1395, which also has a secondary bar, and NGC 6890) or optical I (NGC 3281) analyses."
 The other two (NGC 2639 and NGC 1162) ire classified as barred for the first time., The other two (NGC 2639 and NGC 4162) are classified as barred for the first time.
 For 15 (9 active. 6 coutrol) out of 21 (11 active. LO coutrol) of the optically classified barred galaxies (SB or SX) we find that a secondary bar (or a disk. a leuse or au elongated rine: see table 5. for those cases for which the reported ceutral elongatiou remains uncertaiu) is preseut.," For 15 (9 active, 6 control) out of 24 (14 active, 10 control) of the optically classified barred galaxies (SB or SX) we find that a secondary bar (or a disk, a lense or an elongated ring; see table \ref{bars} for those cases for which the reported central elongation remains uncertain) is present."
 À discussion on the physical properties of these ealaxies. together with the comparison of the properties of active versus non-active ealaxies derived from these datz will be presented in a forthcoming paper (Márquez et al.," A discussion on the physical properties of these galaxies, together with the comparison of the properties of active versus non-active galaxies derived from these data will be presented in a forthcoming paper (Márrquez et al.,"
 nu preparation)., in preparation).
Since also have la,Since we also have 1.
belro2 &We consider the cases when r >2and À =xor r >Q0and A €(0.x).Theore," We consider the cases when $r\ge 2$ and $\lambda=\infty$ or $r\ge0$ and $\lambda\in(0,\infty)$."
"m As in the case of indisunguishable balls. using and we get the following An urn contains WAL balls, AZ balls of each of .V colors."," As in the case of indistinguishable balls, using and we get the following An urn contains $NM$ balls, $M$ balls of each of $N$ colors."
 From the urn a simple random sample of 1 elements is drawn., From the urn a simple random sample of $n$ elements is drawn.
 We want to study theasymptoties of the number of colors with exactly + balls in the sample., We want to study theasymptotics of the number of colors with exactly $r$ balls in the sample.
" Moreprecisely, let £;—en denote the number of balls of color 7, 7=1..... κ,"," Moreprecisely, let $\xi_i=\xi_i^{(n)}$ denote the number of balls of color $i$ , $i=1,\ldots,N$ ."
 Then, Then
outflows/jets with the approach adopted in Cao(002).,outflows/jets with the approach adopted in \citet{c2002a}.
 We summarize the model as follows (seeCao2002.forthedetails).," We summarize the model as follows \citep*[see][for the
details]{c2002a}."
 For a relativistic jet accelerated by the magnetic tield of the disc. the Alfvénn velocity is (Michel1969:Camenzind1986) where and py are the poloidal field strength and the density of the noutflow/jet at Alfvénn point. and +; is the Lorentz factor of the bulk motion of the outflows/Jets.," For a relativistic jet accelerated by the magnetic field of the disc, the Alfvénn velocity is \citep{1969ApJ...158..727M,c1986}
 where $B_{\rm p}^{\rm A}$ and $\rho_{\rm A}$ are the poloidal field strength and the density of the outflow/jet at Alfvénn point, and $\gamma_{\rm j}$ is the Lorentz factor of the bulk motion of the outflows/jets."
 In this work. all our calculations of the outflows/Jets are in the special relativistic frame.," In this work, all our calculations of the outflows/jets are in the special relativistic frame."
 The Alfvénn velocity eI*4O(C(R4). where A is the radius of the Alfvénn point along the field line. and O(/24) is the angular velocity of the accretion flow at the field footpoint {η ," The Alfvénn velocity $v_{\rm A}\sim R_{\rm A}\Omega(R_{\rm d})$, where $R_{\rm
A}$ is the radius of the Alfvénn point along the field line, and $\Omega(R_{\rm d})$ is the angular velocity of the accretion flow at the field footpoint $R_{\rm d}$."
The mass and magnetic flux conservation along the field line requires where mn is the mass loss rate in the outflow/jet from unit surface area of the disc., The mass and magnetic flux conservation along the field line requires where $\dot{m}_{\rm w}$ is the mass loss rate in the outflow/jet from unit surface area of the disc.
 The final bulk velocity of the magnetically driven outflow/Jet is ey. so the Lorentz factor of the outflow/Jet is Combining equations (14))}-17)). the mass loss rate in the outflow/jet from the unit surface area of the dise is The origin of the ordered magnetic fields threading the dise is still unclear.," The final bulk velocity of the magnetically driven outflow/jet is $\sim v_{\rm A}$, so the Lorentz factor of the outflow/jet is Combining equations \ref{b_p}) \ref{gam_j}) ), the mass loss rate in the outflow/jet from the unit surface area of the disc is The origin of the ordered magnetic fields threading the disc is still unclear."
 Tt was suggested that the magnetic fields can be generated through dynamo processes in the dise (e.g.Shakura&al. 1998). or the large-scale external magnetic fields are transported inward by the accretion flow (e.g..Bisnovatyi-Kogan&Ruzmaikin 2005)..," It was suggested that the magnetic fields can be generated through dynamo processes in the disc \citep*[e.g.,][]{s1973,1996MNRAS.281..219T,1998ApJ...501L.189A,1998ApJ...500..703R}, or the large-scale external magnetic fields are transported inward by the accretion flow \citep*[e.g.,][]{1976Ap&SS..42..401B,1994MNRAS.267..235L,2005ApJ...629..960S}."
 For simplicity. we assume the strength of the large-scale magnetic fields threading the dise to be comparable with that of the fields in the accretion disc. {μμστD.," For simplicity, we assume the strength of the large-scale magnetic fields threading the disc to be comparable with that of the fields in the accretion disc, $B_{\rm pd}\simeq B$."
 The magnetic pressure is conventionally assumed to be proportional to the gas pressure in the accretion flow., The magnetic pressure is conventionally assumed to be proportional to the gas pressure in the accretion flow.
 Thus. we have where «71 the ratio of the gas pressure to the total pressure. /5 is the strength of the magnetic fields in the accretion flow.," Thus, we have where $\beta$ is the ratio of the gas pressure to the total pressure, $B$ is the strength of the magnetic fields in the accretion flow."
 The magnetically driven outflow is described by the terminal velocity of the outflow ¢(Alfvénn velocity ον) when the values of two parameters ¢ and 7 are specitied., The magnetically driven outflow is described by the terminal velocity of the outflow (Alfvénn velocity $v_{\rm A}$ ) when the values of two parameters $\zeta$ and $\beta$ are specified.
 We use the Runge-Kutta method to solve a set of five differential equations CD).(3.05)47) and (8) for five variables: p. ep. QO. T. and 7i with suitable boundary conditions at the outer radius Au.," We use the Runge-Kutta method to solve a set of five differential equations (1),(3),(5),(7) and (8) for five variables: $\rho$, $v_{\rm R}$, $\Omega$, $T_{\rm e}$ and $T_{\rm i}$ with suitable boundary conditions at the outer radius $R_{\rm out}$."
 In our calculations. we adopt the black hole mass AJ=10M. fora typical AGN.," In our calculations, we adopt the black hole mass $M=10^{8}M_\odot$ for a typical AGN."
 The conventional values of the disc parameters: a=0.1 and 9.=0.1. are adopted in all our calculations.," The conventional values of the disc parameters: $\alpha=0.1$ and $\delta=0.1$, are adopted in all our calculations."
 The temperature of the ADAF at the outer radius is adopted as described by the self-similar solution of Narayan&Yi(1995b)., The temperature of the ADAF at the outer radius is adopted as described by the self-similar solution of \citet{n1995b}.
. We use a shooting point method in our calculations., We use a shooting point method in our calculations.
" Integrating these five equations from the outer boundary. of the flow at /?—A, inwards toward the black hole.we can obtain the global structure of the accretion flow passing the sonic point smoothly to the black hole horizon by tuning the value of radial velocity at 2.4."," Integrating these five equations from the outer boundary of the flow at $R=R_{\rm out}$ inwards toward the black hole,we can obtain the global structure of the accretion flow passing the sonic point smoothly to the black hole horizon by tuning the value of radial velocity at $R_{\rm out}$ ."
 The outer radius of the accretion flow Raw is adopted., The outer radius of the accretion flow $R_{\rm out}=5000R_{\rm g}$ is adopted.
 We find that the structure of the ADAF is insensitive to the outer boundary conditions., We find that the structure of the ADAF is insensitive to the outer boundary conditions.
 As discussed in Sect. 990..," As discussed in Sect. \ref{equations},"
 the magnetically driven outflow is described by ey. ¢ and -7.," the magnetically driven outflow is described by $v_{\rm A}$, $\zeta$ and $\beta$."
 The terminal bulk velocity of the outflow is comparable with the Alfénn velocity ey., The terminal bulk velocity of the outflow is comparable with the Alfénn velocity $v_{\rm A}$.
 For any unbounded outflows. its bulk velocity should ey. which implies eek.," For any unbounded outflows, its bulk velocity should $\ga v_{\rm
K}$, which implies $v_{\rm A}\ga v_{\rm K}$."
 The mass accretion rate at the outer radius. AJ=10Mp (the Eddington rate is defined as Apad=15LOMSALSAL. gs ο. ἐς adopted for the calculations plotted in Figs. 1-— 3..," The mass accretion rate at the outer radius, $\dot{M}=10^{-5}\dot{M}_{\rm Edd}$ (the Eddington rate is defined as $\dot{M}_{\rm Edd}=1.5\times10^{18} M/{{\rm M}_\odot}$ g $^{-1}$ ), is adopted for the calculations plotted in Figs. \ref{indexs1}- – \ref{va1}."
 In Fig. l..," In Fig. \ref{indexs1},"
 the global solutions forthe ADAFs with outflows are shown with different magnetic field strengths and distributions (.e.. different values of απ ος in which ey= is adopted.," the global solutions for the ADAFs with outflows are shown with different magnetic field strengths and distributions (i.e., different values of $\beta$ and $\zeta$ ), in which $v_{\rm A}=v_{\rm K}$ is adopted."
 In the left panel of Fig. |.," In the left panel of Fig. \ref{indexs1},"
 the radial velocity and the sound speed as functions of radius with different values of .? are plotted for ¢=1., the radial velocity and the sound speed as functions of radius with different values of $\beta$ are plotted for $\zeta=1$.
 In the right panel of Fig. |..," In the right panel of Fig. \ref{indexs1},"
 the mass accretion rates as functions of radius are plotted for the global solutions with different values of :J and c., the mass accretion rates as functions of radius are plotted for the global solutions with different values of $\beta$ and $\zeta$ .
 In Fig. 2..," In Fig. \ref{zeta1},"
 we plot different quantities of ADAF solutions with ¢=1 and +. respectively. where ry= and 3=0.9 are adopted.," we plot different quantities of ADAF solutions with $\zeta=1$ and 4, respectively, where $v_{\rm
A}=v_{\rm K}$ and $\beta=0.9$ are adopted."
 As comparison. we also plot the ADAF solutions without outflows in the same figure.," As comparison, we also plot the ADAF solutions without outflows in the same figure."
 In the calculations. the terminal velocity of the outflow is a free parameter.," In the calculations, the terminal velocity of the outflow is a free parameter."
 We compare the global ADAF solutions with different values of ey (he. ey= and ey= 2ex)in Fig.," We compare the global ADAF solutions with different values of $v_{\rm A}$ (i.e., $v_{\rm A}=v_{\rm K}$ and $v_{\rm
A}=2v_{\rm K}$ ) in Fig."
 3. with ¢=0.9 and ¢=1., \ref{va1} with $\beta=0.9$ and $\zeta=1$.
 In Figs. 4-—6..," In Figs. \ref{indexs2}- \ref{va2},"
" we plot the results calcwated with a relatively high mass accretion rate. ji,=10.> at the outer radius of the disc."," we plot the results calculated with a relatively high mass accretion rate, $\dot{m}_0=10^{-2}$, at the outer radius of the disc."
 The ratio of the bolometric luminosity of accretion dise Livi to kinetic power of outflows {ον is calculated with We plot the ratio {μην as functions of magnetic field strength «21η Fig. 7..," The ratio of the bolometric luminosity of accretion disc $L_{\rm
bol}$ to kinetic power of outflows $P_{\rm k,w}$ is calculated with We plot the ratio $L_{\rm bol}$ $P_{\rm k,w}$ as functions of magnetic field strength $\beta$ in Fig. \ref{ratio},"
 where oy=cy is adopted.," where $v_{\rm
A}=v_{\rm K}$ is adopted."
 A fraction of gases in accretion flow is carried away by the magnetically driven outflow. which leads to mass accretion rate of the accretion flow decreasing towards the black hole.," A fraction of gases in accretion flow is carried away by the magnetically driven outflow, which leads to mass accretion rate of the accretion flow decreasing towards the black hole."
 In many previous works. the outflow is induced by assuming the mass accretion rate 7? to be a power-law dependence of radius rox? (e.g..Blandford&Begelman1999).," In many previous works, the outflow is induced by assuming the mass accretion rate $\dot{m}$ to be a power-law dependence of radius $\dot{m} {\propto} r^s$ \citep*[e.g.,][]{b1999}."
. In this work. we obtain global solutions of ADAFs with magnetically driven outflows.," In this work, we obtain global solutions of ADAFs with magnetically driven outflows."
 Our results show that the mass accretion rate 74? decreases towards the black hole. which is close to a power-law r-dependence at larger radii. while it deviates from a power-law ;-dependence inthe inner region theADAF close to the black hole (see Figs.," Our results show that the mass accretion rate $\dot{m}$ decreases towards the black hole, which is close to a power-law $r$ -dependence at larger radii, while it deviates from a power-law $r$ -dependence inthe inner region of the ADAF close to the black hole (see Figs."
 | and 4., \ref{indexs1} and \ref{indexs2}) ).
 The large-scaleof magnetic fields are assumed to thread the accretion flow. which are believed to accelerate the outflow from the accretion disc.," The large-scale magnetic fields are assumed to thread the accretion flow, which are believed to accelerate the outflow from the accretion disc."
 In this work. the magnetic strength atthe dise surface is deseribed by a parameter 2. which is limited by the gas pressure," In this work, the magnetic strength atthe disc surface is described by a parameter $\beta$ which is limited by the gas pressure"
(allowing the exomoon to become unbound).,(allowing the exomoon to become unbound).
" It is possible that the exomoon could have survived and remained bound to the planet in these scenarios, but that is beyond the fidelity of these simulations."," It is possible that the exomoon could have survived and remained bound to the planet in these scenarios, but that is beyond the fidelity of these simulations."
" To estimate the detectability of the resulting systems, we calculated root-mean-squared Transit Timing Variations (TTVs) and Duration Variations (TDVs) for each simulation."," To estimate the detectability of the resulting systems, we calculated root-mean-squared Transit Timing Variations (TTVs) and Duration Variations (TDVs) for each simulation."
 The TTV and TDV are due to the orbit of the exomoon causing the exoplanet to begin the transit either sooner or later than the barycenter of the planet-moon system., The TTV and TDV are due to the orbit of the exomoon causing the exoplanet to begin the transit either sooner or later than the barycenter of the planet-moon system.
 These effects are maximized for low mass ratios (e.g. Earth mass exomoon around a Neptune-mass planet) and longer period exomoon orbits., These effects are maximized for low mass ratios (e.g. Earth mass exomoon around a Neptune-mass planet) and longer period exomoon orbits.
" Since the majority of resulting systems were low-inclination with respect to the stellarcentric orbit, we used the zero-inclination equations from Kipping"," Since the majority of resulting systems were low-inclination with respect to the stellarcentric orbit, we used the zero-inclination equations from \citet{Kipping2009}."
" Assuming zero-inclination maximizes the TTV and (2009)..TDV, allowing an estimate of whether the system would be even detectable in the best case scenario."," Assuming zero-inclination maximizes the TTV and TDV, allowing an estimate of whether the system would be even detectable in the best case scenario."
 The general result of the simulations is that loosely-captured exomoons around giant planets in habitable zones can survive to stabilize into long-lived orbits., The general result of the simulations is that loosely-captured exomoons around giant planets in habitable zones can survive to stabilize into long-lived orbits.
 Table 1 shows the fraction that stabilize for each mass/distance scenario., Table \ref{tab:gen} shows the fraction that stabilize for each mass/distance scenario.
" The exomoons stabilized much easier in the FO case than the MO case, with the solar-mass case in-between."," The exomoons stabilized much easier in the F0 case than the M0 case, with the solar-mass case in-between."
" The reason for this is readily apparent from the observation that while Hill radius scales linearly with distance from the star, the equilibrium temperature scales as the inverse square of distance."," The reason for this is readily apparent from the observation that while Hill radius scales linearly with distance from the star, the equilibrium temperature scales as the inverse square of distance."
" Thus, the super-solar mass star allows a much wider Hill radius at a habitable distance than for a planet of equal mass around a sub-solar mass star."," Thus, the super-solar mass star allows a much wider Hill radius at a habitable distance than for a planet of equal mass around a sub-solar mass star."
" This larger Hill radius provides the post-capture orbit much more room to move around, allowing longer period exomoons to stabilize."," This larger Hill radius provides the post-capture orbit much more room to move around, allowing longer period exomoons to stabilize."
 The truncation of periods greater than two days for the MO case in Figure 3 is a clear result of this., The truncation of periods greater than two days for the M0 case in Figure \ref{fig:circ} is a clear result of this.
" However, it also apparent from Figure 3 that the fast stellarcentric orbits of the MO case allow for much more rapid exomoon circularization that for the larger stars."," However, it also apparent from Figure \ref{fig:circ} that the fast stellarcentric orbits of the M0 case allow for much more rapid exomoon circularization that for the larger stars."
" This clearly shows the effect of Kozai-pumped eccentricities accelerating tidal decay, as otherwise the timescales would be independent of the distance from the star."," This clearly shows the effect of Kozai-pumped eccentricities accelerating tidal decay, as otherwise the timescales would be independent of the distance from the star."
" The median circularization timescales for the MO cases are of order 103 years, 10? years for the solar-mass star, and 10° years for the FO star."," The median circularization timescales for the M0 cases are of order $10^4$ years, $10^5$ years for the solar-mass star, and $10^6$ years for the F0 star."
" All of these timescales are very short relative to the lifetime of the star, but 10° years may be long enough that the exomoon's orbit could be externally perturbed by a planet or planetesimal disk."," All of these timescales are very short relative to the lifetime of the star, but $10^6$ years may be long enough that the exomoon's orbit could be externally perturbed by a planet or planetesimal disk."
" On the other hand, as Figure 1 shows, the semimajor axis decay typically happens at least an order of magnitude faster than full circularization."," On the other hand, as Figure \ref{fig:tale} shows, the semimajor axis decay typically happens at least an order of magnitude faster than full circularization."
" Therefore, it is a reasonable assumption that nearly all these orbits would stabilize before any external perturbation would disrupt them."," Therefore, it is a reasonable assumption that nearly all these orbits would stabilize before any external perturbation would disrupt them."
" 'To gauge the amount of tidal heating on the satellite, we estimated the amount of energy lost from the system from the difference in the initial and final orbits and spins."," To gauge the amount of tidal heating on the satellite, we estimated the amount of energy lost from the system from the difference in the initial and final orbits and spins."
" For the orbits, we first found the energy difference between the initial and final states."," For the orbits, we first found the energy difference between the initial and final states."
" Then, we muliplied this an estimate of the fraction of energy that went into the satellite using the masses and tidal Q of the QpianetMsat/(QptanetMsat+ This term "," Then, we muliplied this an estimate of the fraction of energy that went into the satellite using the masses and tidal $Q$ of the objects, $Q_{planet}M_{sat}/(Q_{planet}M_{sat}+Q_{sat}M_{planet})$ ."
"objects,was usually near unity, as the QsatMpianet):planet was assumed to have a very large Q."," This term was usually near unity, as the planet was assumed to have a very large $Q$."
" We then estimated the change in rotational energy of the satellite, subtracted this from the change in orbital energy, and divided the result by the circularization time to produce a heating rate."," We then estimated the change in rotational energy of the satellite, subtracted this from the change in orbital energy, and divided the result by the circularization time to produce a heating rate."
" The rate was dominated by the orbital term, as the initial rotation periods were not too different from the final orbital periods."," The rate was dominated by the orbital term, as the initial rotation periods were not too different from the final orbital periods."
" Generally, this rate was higher per unit mass for larger satellites circularized and closer orbits."," Generally, this rate was higher per unit mass for larger satellites (which circularized faster) and closer orbits."
" For the Earth-mass/Jupiter-mass(which faster)case, the median rates were 0.002 mW/kg around a FO star, 0.02 mW/kgaround a solar-mass star, and 0.5 mW/kg around a MO star."," For the Earth-mass/Jupiter-mass case, the median rates were 0.002 mW/kg around a F0 star, 0.02 mW/kgaround a solar-mass star, and 0.5 mW/kg around a M0 star."
we find from Eq. (,we find from Eq. (
25) that cosd is always positive ie. 9 lies in the first ancl fourth quadrant.,25) that $\delta$ is always positive i.e. $\delta$ lies in the first and fourth quadrant.
 Fig., Fig.
 1., 1.
 eives the correlation plots for this texture., gives the correlation plots for this texture.
" There exists a lower bound on effective Majorana neutrino mass A,> 0.042 eV ancl 0644>0.35"" We get hiehlv constrained parameter space for (he (wo Majorana tvpeCP - violating phases o and2 (Fig.", There exists a lower bound on effective Majorana neutrino mass $M_{ee} >$ 0.042 eV and $\theta_{13}> 0.35^o$ .We get highly constrained parameter space for the two Majorana typeCP - violating phases $\alpha$ and$\beta$ (Fig.
 1(b))., 1(b)).
 Larger values of 86415 are allowed for 9 near 90°.270° as seen [rom Fie.," Larger values of $\theta_{13}$ are allowed for $\delta $ near $90^o, 270^o$ as seen from Fig."
 l(c)., 1(d).
 The Jarlskog rephasing invariant Jep is within the range (-0.45)-(0.45) (Fig., The Jarlskog rephasing invariant $J_{CP}$ is within the range (-0.45)-(0.45) (Fig.
 1(0))., 1(c)).
 This texture structure has zero 13 element ancl zero minor corresponding to 11 entry (C4= 0)., This texture structure has zero 13 element and zero minor corresponding to 11 entry $C_{11}=0$ ).
 Class 3A also gives clear inverted mass hierarchy of neutrino masses., Class 3A also gives clear inverted mass hierarchy of neutrino masses.
 This texture has same values of ον. As and Ay as for Class 2A since both have the same vanishing minor.," This texture has same values of $A_1$, $A_2$ and $A_3$ as for Class 2A since both have the same vanishing minor."
 The analytical approximations [for (the mass ratios in (he leading order of s(4 are given as Ii can be seen from Eq. (, The analytical approximations for the mass ratios in the leading order of $s_{13}$ are given as It can be seen from Eq. (
26) that p is always greater than 1. which gives inverted hierarchy of neutrino masses.,"26) that $\rho$ is always greater than 1, which gives inverted hierarchy of neutrino masses."
 Since. {he mass ratio σ is always smaller than one we lind from Eq. (," Since, the mass ratio $\sigma$ is always smaller than one we find from Eq. ("
27) that cosd is always negative ie. 9 lies in the second. ancl (hid quadrant.,27) that $\delta$ is always negative i.e. $\delta$ lies in the second and third quadrant.
 Fig., Fig.
 2., 2.
 gives some interesting correlation plots for this texture structure., gives some interesting correlation plots for this texture structure.
" A ower bound on A,> 0.044 eV and 1-3 mixing angle (044> 0.37) is obtained for this class.", A lower bound on $M_{ee}>$ 0.044 eV and 1-3 mixing angle $\theta_{13}>0.3^o$ ) is obtained for this class.
" Larger values of θ are allowed near 907.270""."," Larger values of $\theta_{13}$ are allowed near $90^o, 270^o$."
 The rephasing invariant. Jep ies in the range (-0.045)-(0.045).," The rephasing invariant, $J_{CP}$ lies in the range (-0.045)-(0.045)."
 lt is found that the limit of a vanishing mass eigenvalue ie. πως can be reached or both class 24 and 3A. As can be seen from Fig., It is found that the limit of a vanishing mass eigenvalue i.e. $m_3$ =0 can be reached for both class 2A and 3A. As can be seen from Fig.
 2(b) that very small range (0.37«044 1.17) is allowed for the limit of vanishing mass eigenvalue 1ης., 2(b) that very small range $0.3^o\le$$\theta_{13}$$\le1.1^o$ ) is allowed for the limit of vanishing mass eigenvalue $m_3$ =0.
" It is interesting to note that class 3A is transformed to class 2A andvice-versa by the transformation 0—4-3 . Ao,—(4 034)."," It is interesting to note that class 3A is transformed to class 2A andvice-versa by the transformation $\delta\rightarrow\delta+\pi$ , $\theta_{23}\rightarrow(\frac{\pi}{2}-\theta_{23})$ ."
" Therefore. the predictions lor neutrino mass matrices for these (wo classes will be identical for all neutrino parameters except 905, and/ or 9."," Therefore, the predictions for neutrino mass matrices for these two classes will be identical for all neutrino parameters except $\theta_{23}$ and/ or $\delta$ ."
1908).,.
. For a review see Backinan&DParesce (1993).., For a review see \citet{back:93}. .
 Receutly Plets&Vyvuckier(1999). have discussed these earlier results aud coucludd that a significant excess at GOfan is found in 13+10% «of all mmain sequence stars with spectral type A. EF. CG aud Is. Uufortunatelv all these studies based ou IRAS «ata olvowere affected by severe selection effects aud did not answer muportaut questions such as: Will a star loose its cisk when it erows older?," Recently \citet{plet:99} have discussed these earlier results and concluded that a significant excess at $60\mu$ m is found in $13\pm 10\%$ of all main sequence stars with spectral type A, F, G and K. Unfortunately all these studies based on IRAS data only were affected by severe selection effects and did not answer important questions such as: Will a star loose its disk when it grows older?"
 On what tine-scale?, On what time-scale?
 Does the pTOsClCO ο: plajets depend on the stellar imain-sequence Mass, Does the presence of planets depend on the stellar main-sequence mass?
 Do un]tiple stars have disks nore. or less frequently?," Do multiple stars have disks more, or less frequently?"
 Do sars that formed in clusters lave ¢isks less oftcn?, Do stars that formed in clusters have disks less often?
 With such questions uaanswered we clearly do not 1iderstaud t svstematies of the formation of solar svstenis., With such questions unanswered we clearly do not understand the systematics of the formation of solar systems.
 Tere we present results of a coutinuatiou with ISO (Isessleretal.1900) of the succesful search of IRAS., Here we present results of a continuation with ISO \citep{kess:96} of the succesful search of IRAS.
 Our aiu has been to obtain a better «chi1ο sanuple of stars., Our aim has been to obtain a better defined sample of stars.
 The major step forward in this paper is nof in the detection of more renimaut disks. mut iu reliable information about the presence of a disk.," The major step forward in this paper is not in the detection of more remnant disks, but in reliable information about the presence of a disk."
 Laer reports ou results from our prograunhave been eivon in Tabingetal.(1996). Donüuiketal. (1998).. JourdaindeMuizonetal.(1f100) and IHabimgetal. (1999).," Earlier reports on results from our programhave been given in \citet{habi:96}, \citet{domi:98}, \citet{jour:99} and \citet{habi:99}."
. Stars were selected so that heir photospheric flux was witlin our seusitivitv init., Stars were selected so that their photospheric flux was within our sensitivity limit.
 Any excess would then appear nunediatelv., Any excess would then appear immediately.
 We also wautec ο lake certain that auv excess flux should be attributed to a circumstellar disk and not to some other pro)ertv of the star. such as cireunistellar matter ejected dving the stellar evolution or to the presence of a red conrMOL.," We also wanted to make certain that any excess flux should be attributed to a circumstellar disk and not to some other property of the star, such as circumstellar matter ejected during the stellar evolution or to the presence of a red companion."
 Iu selecting our stars we usd the following criteria: To illustrate what stars are bright cnough to be included we use an equation that eives the stellar colour. (V[60:00]). as a faction o0 V).," In selecting our stars we used the following criteria: To illustrate what stars are bright enough to be included we use an equation that gives the stellar colour, $V-[60 \mu m]$ ), as a function of $(B-V)$."
 The equation has been derived ορήσαν from IRAS data by Watersetal.(LOST): we use a slightly different version given by IL. Plets (private communication): The zero point in this equation has a formal error of 0.01., The equation has been derived empirically from IRAS data by \citet{wate:87}; we use a slightly different version given by H. Plets (private communication): The zero point in this equation has a formal error of 0.01.
 Tutvinsic. recddemime-free (2BV) values must be used. but all our stars are nearby aud we assunie that the measured values are reddenine-free.," Intrinsic, reddening-free $(B-V)$ values must be used, but all our stars are nearby and we assume that the measured values are reddening-free."
 A posteriori we checked that we may safely ignore the reddening produced by the disks that we detected: only in the case of .} Pic is a πια] effect expected., A posteriori we checked that we may safely ignore the reddening produced by the disks that we detected; only in the case of $\beta$ Pic is a small effect expected.
" Adopting a flux density of 1.19 Jv for [60,0]=0 we find apparent-maenitude limits and distance limits of suitable iuain-sequeuce stars as stnunarized in Table 1..", Adopting a flux density of 1.19 Jy for $[60\mu m]=0$ we find apparent-magnitude limits and distance limits of suitable main-sequence stars as summarized in Table \ref{tab:distlimits}.
 The distance lit varies strongly with spectral type., The distance limit varies strongly with spectral type.
 Table 2 contains basic data on a] stars from the Eje for which we preseut ISO data., Table \ref{tab:stars} contains basic data on all stars from the sample for which we present ISO data.
 Columns 1 aud 2 continu the umber of the star in the ΠΟ aud in the Tipparcos Catalogue (PerrvinaLeal.1997) and COULun 23 the nune., Columns 1 and 2 contain the number of the star in the HD and in the Hipparcos Catalogue \citep{perr:97} and column 3 the name.
 | “and BV have been aken from the Geneva photometric catalogue {Ixuuzhletal.1997)., $V$ and $B-V$ have been taken from the Geneva photometric catalogue \citep{kunz:97}.
. Cohmus 6 and 7 coutain the distance and t1ο spectral type as elven in tie Tipparcos Catalog (Perrvinan1997 )., Columns 6 and 7 contain the distance and the spectral type as given in the Hipparcos Catalog \citep{perr:97}.
. The ageOo oeiven in cohuun 8 is frou Lacinecta. (1999).. where errors iu the age determilatious are discussed.," The age given in column 8 is from \citet{lach:99}, where errors in the age determinations are discussed."
" The effective. temperature in coluuu 11 has been derived by fitting INurucz model atmuospheres to the Geneva photometry: we will need this temperature to calculate the cdist nass from thefliux-deusitv excess at ""un.", The effective temperature in column 11 has been derived by fitting Kurucz' model atmospheres to the Geneva photometry; we will need this temperature to calculate the dust mass from theflux-density excess at .
. Pre-launch reconunecudations mace us start with chopped meastrelents (observing mode PITTOUS: see, Pre-launch recommendations made us start with chopped measurements (observing mode PHT03; see
"restricted to the galaxies with IHE A£g,711d10/9AL. and distances smaller than 100 Mpc. or which rof.","restricted to the galaxies with HI $M_{HI}>1.1\times 10^{10}\
M_\odot$ and distances smaller than 100 Mpc, for which ref."
 [7] found a maximal correlation., \cite{gh08} found a maximal correlation.
 From the results depicted it is seen that the data wave a distance smaller than what is obtained in of the isotropic simulations. explaining the aree anisotropy probability found.," From the results depicted it is seen that the data have a distance smaller than what is obtained in of the isotropic simulations, explaining the large anisotropy probability found."
" Ou the other haud. events simulated according to the same uodel (southern radio galaxies with Ay,1d«1019AL. aud ο9.1 Jv kun 1) have henmiselves μπαλα: distances (left histogram) than the isotropic simulations aud the distance of he data is quite consistent with then."," On the other hand, events simulated according to the same model (southern radio galaxies with $M_{HI}>1.1\times 10^{10}\
M_\odot$ and $S_{int}>9.4$ Jy km $^{-1}$ ) have themselves smaller distances (left histogram) than the isotropic simulations and the distance of the data is quite consistent with them."
 It is also iteresting to note that simulated data following the jiearby. SWIFT ACN have a distribution of distances (central histogram) somewhat intermediate )etween the TICAT auc isotropic ones., It is also interesting to note that simulated data following the nearby SWIFT AGN have a distribution of distances (central histogram) somewhat intermediate between the HICAT and isotropic ones.
 For the present uuuber of events aud in this particular colparison. the data are still quite compatible with both the WICAT and SWIFT models. aud is ouly mareinally compatible with isotropy.," For the present number of events and in this particular comparison, the data are still quite compatible with both the HICAT and SWIFT models, and is only marginally compatible with isotropy."
 The bottourright paucl iu fig., The bottom-right panel in fig.
 2 is for a reference catalog obtained restricting the IICAT sample to the galaxies with HI mass bigeer than 2«1019AL..., 2 is for a reference catalog obtained restricting the HICAT sample to the galaxies with HI mass bigger than $2\times 10^{10}\ M_\odot$.
 We see in this case that the auisotropy probability is about.. distavoring the isotropic lapotlesis. but however the same HICAT model would typically eive mach smaller distances than the data (only ofthe ΠΙΟΑΤsinimlatious have a lareer value of D than the data).," We see in this case that the anisotropy probability is about, disfavoring the isotropic hypothesis, but however the same HICAT model would typically give much smaller distances than the data (only of the HICATsimulations have a larger value of $D$ than the data)."
 This implies that the data from the Auger Observatory do not look like a typical realization of this particlar model. even if the quoted anisotropy probability is large. showing clearly why it is not convenient to call it a correlation probability.," This implies that the data from the Auger Observatory do not look like a typical realization of this particular model, even if the quoted anisotropy probability is large, showing clearly why it is not convenient to call it a correlation probability."
 Ou the other haud. iu this test the simulations according to the SWIFT AGN model considered give a distribution of distances which eucompasses quite comfortably the value found for the data.," On the other hand, in this test the simulations according to the SWIFT AGN model considered give a distribution of distances which encompasses quite comfortably the value found for the data."
 It is however important to keep in nud that the largest distance between the ACN based models aud the reference TIPASS catalog could be in directions (and quadrants) completely differcut from that eivine the largest ciffercuce between the data and the IHIPASS reference catalog. so that the compatibility found in this particular test does not imply necessarily that the AGN model considered is a good CR source model.," It is however important to keep in mind that the largest distance between the AGN based models and the reference HIPASS catalog could be in directions (and quadrants) completely different from that giving the largest difference between the data and the HIPASS reference catalog, so that the compatibility found in this particular test does not imply necessarily that the AGN model considered is a good CR source model."
 Tn this seuse it is important to simultaneously check which is the value of the distance between the data aud the SWIFT AGN reference catalog to better establish the compatibility of the data with this particular source model., In this sense it is important to simultaneously check which is the value of the distance between the data and the SWIFT AGN reference catalog to better establish the compatibility of the data with this particular source model.
 From the different plots in fig., From the different plots in fig.
 2 one hence concludes that the WIPASS based models cousidered in the first three panels are all consistent with the data results. while the one in the fourth paucl is not.," 2 one hence concludes that the HIPASS based models considered in the first three panels are all consistent with the data results, while the one in the fourth panel is not."
 It is also worth noting that the higher level of rejection of isotropy obtained iun the third panel doesut imply necessarily a preference towards sources with A£j;71.1«1019AJ..., It is also worth noting that the higher level of rejection of isotropy obtained in the third panel doesn't imply necessarily a preference towards sources with $M_{HI}>1.1\times 10^{10}M_\odot$.
 Let us also note that in the examples in fe., Let us also note that in the examples in fig.
 2 the data fractious are compared always to the correspouding fractions iu the reference catalog used (the flux weighted IIIPASS. ealaxics with different cuts in each panel). aud the same is doue for the different source models tested (isotropy. SWIFT AGN and HIPASS ones).," 2 the data fractions are compared always to the corresponding fractions in the reference catalog used (the flux weighted HIPASS galaxies with different cuts in each panel), and the same is done for the different source models tested (isotropy, SWIFT AGN and HIPASS ones)."
 In fig., In fig.
 LI. which considered the BAT ACN as reference catalog we didu't show the correspondiug histograms for the HIPASS based models just to facilitate the presentation. but the same could have clearly been douc.," 1, which considered the BAT AGN as reference catalog, we didn't show the corresponding histograms for the HIPASS based models just to facilitate the presentation, but the same could have clearly been done."
 These cross comparisons are quite useful to discriminate amoug possible source models. aud we note that the same can be repeated with other reference catalogs aud source scenarios.," These cross comparisons are quite useful to discriminate among possible source models, and we note that the same can be repeated with other reference catalogs and source scenarios."
 A possible additional iaprovement of the source models would be to weight the differeut sources also by a factor accounting for the suppression of the fluxes above the energy threshold considered due to the GZEK effect., A possible additional improvement of the source models would be to weight the different sources also by a factor accounting for the suppression of the fluxes above the energy threshold considered due to the GZK effect.
 One can then account for the expected suppression at tle actual distance of cach object instead of just climinating the sources bevoud a given specified distance (such as in the examples where only sources within LOO Mpc are kept)., One can then account for the expected suppression at the actual distance of each object instead of just eliminating the sources beyond a given specified distance (such as in the examples where only sources within 100 Mpc are kept).
 This was indeed cousidered in the past in alternative analyses of the CR arrival direction distributions [LL 0]..," This was indeed considered in the past in alternative analyses of the CR arrival direction distributions \cite{wax, cuoco}. ."
 The required weight, The required weight
average dust temperature 7;=25.9/. of our galaxy sample as a representative value to compute the dust mass.,average dust temperature $\overline{T}_d=25.9K$ of our galaxy sample as a representative value to compute the dust mass.
 Table 3. presents the resulting dust masses Άι 7;; and the IRAS flux densities for 10 galaxies in our sample.," Table \ref{tabMd} presents the resulting dust masses $M_d$, $T_d$ and the IRAS flux densities for 10 galaxies in our sample."
 For six galaxies 11374. 11375. 11381. 11419. 11427. 3358-G06) the IRAS Faint Source Catalog lists only upper limits. the three remaining galaxies 11373. 11380A. 11963) are not listed in the catalog.," For six galaxies 1374, 1375, 1381, 1419, 1427, 358-G06) the IRAS Faint Source Catalog lists only upper limits, the three remaining galaxies 1373, 1380A, 1963) are not listed in the catalog."
 Fig., Fig.
 9 presents the fractional difference of the infrared and the optical dispersions as a function of the relative amount of dust in a galaxy. which is estimated in? by the ratio of the IRAS dust mass to the 4-band luminosity.," \ref{figfracdiffMd} presents the fractional difference of the infrared and the optical dispersions as a function of the relative amount of dust in a galaxy, which is estimated in \cite{Silge2003} by the ratio of the IRAS dust mass to the $B$ -band luminosity."
 The estimates for the relative amount of dust are given in Table 3.., The estimates for the relative amount of dust are given in Table \ref{tabMd}.
 Note that the IRAS dust mass estimates are a lower limit for the true dust masses. because IRAS is not sensitive to cold dust Qwhich emits the bulk of its radiation longwards of 1007/m.).," Note that the IRAS dust mass estimates are a lower limit for the true dust masses, because IRAS is not sensitive to cold dust (which emits the bulk of its radiation longwards of $\mu$ m.)."
 A negligible trend is visible (confirmed by the Spearman rank-order correlation coefficient equal to 0.213., A negligible trend is visible (confirmed by the Spearman rank-order correlation coefficient equal to 0.21).
 The best-fitting solid line is given by equation with both the slope and the intercept consistent with O and a reduced 47 of 1.76., The best-fitting solid line is given by equation with both the slope and the intercept consistent with 0 and a reduced $\chi^2$ of 1.76.
 This implies that warm dust does not affect optical dispersions., This implies that warm dust does not affect optical dispersions.
 We cannot yet make the same conclusion for colder dust. but Fig.," We cannot yet make the same conclusion for colder dust, but Fig."
 7 indicates that the typical effect is very weak., \ref{corrCOoptical} indicates that the typical effect is very weak.
 In this study. we investigate a complete magnitude-Iimited and unbiased sample of 22 early-type galaxies in the Fornax cluster and are able to determine the kinematics based on the 2.29//m !CO(- feature for 19 of those galaxies.," In this study, we investigate a complete magnitude-limited and unbiased sample of 22 early-type galaxies in the Fornax cluster and are able to determine the kinematics based on the $\mu$ m $^{12}$ CO(2-0) feature for 19 of those galaxies."
 We related the NIR velocity dispersions with the optical dispersions of ? and found no evidence for a a@-discrepancy for the ellipticals nor for the lenticulars., We related the NIR velocity dispersions with the optical dispersions of \cite{Kuntschner2000} and found no evidence for a $\sigma$ -discrepancy for the ellipticals nor for the lenticulars.
 Our results agree with a previous smaller study of ?.. but not with ?..," Our results agree with a previous smaller study of \cite{Silva2008}, but not with \cite{Silge2003}."
 We investigated this disagreement by providing a variety of input templates to pPXF with a large range of EWs. by changing the spatial width of the extraction window and by introducing Hermite coefficients in the LOSVD. but we were not able to clarify this discrepancy.," We investigated this disagreement by providing a variety of input templates to pPXF with a large range of EWs, by changing the spatial width of the extraction window and by introducing Gauss-Hermite coefficients in the LOSVD, but we were not able to clarify this discrepancy."
 We have computed the dust masses based on IRAS flux densities for 10. galaxies of our sample and investigated the influence of diffuse dust on the observed kinematics. which turned out to be negligible.," We have computed the dust masses based on IRAS flux densities for 10 galaxies of our sample and investigated the influence of diffuse dust on the observed kinematics, which turned out to be negligible."
 The one-to-one correspondence between the optical and the NIR velocity dispersions found for this homogeneous set of early-type galaxies implies that velocity dispersions measured a optical wavelengths are reliable kinematic parameters for early-type galaxies and hence that no bias is introduced in statistica relations that build on such dispersions (such as the Myj-c relation or the Fundamental Plane)., The one-to-one correspondence between the optical and the NIR velocity dispersions found for this homogeneous set of early-type galaxies implies that velocity dispersions measured at optical wavelengths are reliable kinematic parameters for early-type galaxies and hence that no bias is introduced in statistical relations that build on such dispersions (such as the $\text{M}_{\text{BH}}$ $\sigma$ relation or the Fundamental Plane).
 Combined with the simulations by ??.. it also supports the traditional point of view on the dust conten of early-type galaxies. namely that they are virtually optically thin.," Combined with the simulations by \cite{Baes2000,Baes2002}, it also supports the traditional point of view on the dust content of early-type galaxies, namely that they are virtually optically thin."
 While some observational studies hinted towards the existence of a substantial diffuse dust component in early-type galaxies (e.g. 2??? y). the most recent results from the recently launched Hersche Space Observatory indicate a dearth of diffuse dust in the few elliptical galaxies studied so far (222).," While some observational studies hinted towards the existence of a substantial diffuse dust component in early-type galaxies (e.g. \citealt{Temi2004,Temi2007,Leeuw2004,Vlahakis2005}) ), the most recent results from the recently launched Herschel Space Observatory indicate a dearth of diffuse dust in the few elliptical galaxies studied so far \citealt{Clemens2010,Baes2010,Gomez2010}) )."
 Whereas our results support this scenario. we must be carefu for two caveats.," Whereas our results support this scenario, we must be careful for two caveats."
 On the one hand. our results only set limits on the presence of a smooth. diffusely distributed dust component. which one would expect if the dust has an internal origin.," On the one hand, our results only set limits on the presence of a smooth, diffusely distributed dust component, which one would expect if the dust has an internal origin."
 If the dust has an external origin. it is not necessarily coincident with the stellar body.," If the dust has an external origin, it is not necessarily coincident with the stellar body."
 An example is the nearby Virgo Cluster elliptical M86. where the dust is clearly related to stripping from a nearby spiral galaxy and is concentrated some 10 kpe to the south-east of the nucleus (?)..," An example is the nearby Virgo Cluster elliptical M86, where the dust is clearly related to stripping from a nearby spiral galaxy and is concentrated some 10 kpc to the south-east of the nucleus \citep{Gomez2010}."
 On the other hand. the comparison of optical and NIR velocity dispersions might not be the most sensitive way to measure the optical thickness of early-type galaxies.," On the other hand, the comparison of optical and NIR velocity dispersions might not be the most sensitive way to measure the optical thickness of early-type galaxies."
 The simulations of indicate an effect of a few percent only for optical depths of order unity.," The simulations of \cite{Baes2000,Baes2002} indicate an effect of a few percent only for optical depths of order unity."
 Combined with the measurement errors and other xossible effects such as different stellar populations dominating the kinematies at optical and NIR wavelengths. our results should be considered only in a statistical sense and one should take care not ο directly interpret results on individual galaxies in terms of optical hiekness.," Combined with the measurement errors and other possible effects such as different stellar populations dominating the kinematics at optical and NIR wavelengths, our results should be considered only in a statistical sense and one should take care not to directly interpret results on individual galaxies in terms of optical thickness."
 A clear example is the case of 11380: in spite of a clear dust lane in optical images (2). and a significant IRAS dust mass of 7.6.107AL.. the NIR velocity dispersion is lower than he optical dispersion.," A clear example is the case of 1380: in spite of a clear dust lane in optical images \citep{Jordan2007} and a significant IRAS dust mass of $7.6 \times 10^5~M_\odot$, the NIR velocity dispersion is lower than the optical dispersion."
 A clearer picture on the optical thickness of early-type galaxies will hopefully emerge in the near future when substantial numbers of nearby early-type galaxies will be observed as part of several Herschel Key Programs. including the Herschel Virgo Cluster Survey (2). and the Herschel Reference Survey (?)..," A clearer picture on the optical thickness of early-type galaxies will hopefully emerge in the near future when substantial numbers of nearby early-type galaxies will be observed as part of several Herschel Key Programs, including the Herschel Virgo Cluster Survey \citep{Davies2010} and the Herschel Reference Survey \citep{Boselli2010}."
 This research has made use of NASA's Astrophysics Data System and the NASA/TPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory. California Institute of Technology. under contract with the National Aeronautics and Space Administration.," This research has made use of NASA's Astrophysics Data System and the NASA/IPAC Extragalactic Database (NED) which is operated by the Jet Propulsion Laboratory, California Institute of Technology, under contract with the National Aeronautics and Space Administration."
 This research also has made use of the SIMBAD database. operated at CDS. Strasbourg. France.," This research also has made use of the SIMBAD database, operated at CDS, Strasbourg, France."
 JV acknowledges FWO Vlaanderen for the financial support., JV acknowledges FWO Vlaanderen for the financial support.
 ATR was supported by the FONDAP Center for Astrophysics CONICYT 15010003 and by the National Science Foundation grants, AJR was supported by the FONDAP Center for Astrophysics CONICYT 15010003 and by the National Science Foundation grants
0 andο.,0 and.
 To the West oft1e nuclear region the gas ds bluc-shited up lo 220 although at the western-most ti» of the nebula the gas velocity again changes direction and is red-shifted. by," To the West of the nuclear region the gas is blue-shifted up to $-220$, although at the western-most tip of the nebula the gas velocity again changes direction and is red-shifted by."
 The central xut of the SW filament is blue-shifted. whilst the SW. tip is redshifted. ius the filament must either be stretching or Collapsing on itself," The central part of the SW filament is blue-shifted, whilst the SW tip is redshifted, thus the filament must either be stretching or collapsing on itself."
 Similar kinematics can be seen in 11275 in he Perseus cluster (2).. and in the models of ?..," Similar kinematics can be seen in 1275 in the Perseus cluster \citep{Hatch06}, and in the models of \citet{Pope2008}."
 The leneh of the SW filament is 42 kpc. implving a dynamical age of ~107 vvr. so the linc-emitting gas must »e lone-lived.," The length of the SW filament is 42 kpc, implying a dynamical age of $\sim10^8$ yr, so the line-emitting gas must be long-lived."
 The dispersion map of the gas is displaved in Fig., The dispersion map of the gas is displayed in Fig.
 LL bb. All line-wieths are e-widths. and not PNCEITME (fall width at valf maximum).," \ref{fig:ifu_kinematics}b b. All line-widths are $\sigma$ -widths, and not FWHM (full width at half maximum)."
 The e-widths of the Ori] and lines used or creating the dispersion map of the gas were corrected for instrumental broadening of aas measured from the nearby sskvline., The $\sigma$ -widths of the ] and $\beta$ lines used for creating the dispersion map of the gas were corrected for instrumental broadening of as measured from the nearby skyline.
 White bins in this figure are regions in which the spectral resolution is too low to determine a dispersion., White bins in this figure are regions in which the spectral resolution is too low to determine a dispersion.
 A peak in the dispersion is seen slightly West of the nuclear region., A peak in the dispersion is seen slightly West of the nuclear region.
 The dispersion here is whilst the median velocity dispersion. is 7200, The dispersion here is whilst the median velocity dispersion is $\sim$.
 The SW filament has the lowest velocity dispersion., The SW filament has the lowest velocity dispersion.
 These strong internal motions may result. from the superposition of distinct components along the line of sight., These strong internal motions may result from the superposition of distinct components along the line of sight.
 The FUY lisht from a galaxy is dominated by emission fron voung O and B tvpe stars so we can use the rest-frame EUM emission to estimate the star formation rate., The FUV light from a galaxy is dominated by emission from young O and B type stars so we can use the rest-frame FUV emission to estimate the star formation rate.
 This estimate of the SER. does not assume that the line emission has been ionized by the voung stellar population., This estimate of the SFR does not assume that the line emission has been ionized by the young stellar population.
 Since the NUV and FUY emission is spatially extended. approximately in the same direction as the filament. it is likely to be emitted from a spatially extended source. such as stars. rather than the central active nucleus.," Since the NUV and FUV emission is spatially extended, approximately in the same direction as the filament, it is likely to be emitted from a spatially extended source, such as stars, rather than the central active nucleus."
 Assuming all of the flux is emitted from a population of hot voung stars. we use the observed NUM. emission. (equivalent to the rest-frame ISTOA)) to obtain a star formation rate.," Assuming all of the flux is emitted from a population of hot young stars, we use the observed NUV emission (equivalent to the rest-frame ) to obtain a star formation rate."
 The NUV tux of L1.10?'erονtem211. is corrected for Galactic extinction. of 0 13— 0.108 and internal extinction of E(D Vj= 0.211.," The NUV flux of $1.1\times10^{-27}\,{\rm erg\, s^{-1}\, cm^{-2}\, Hz^{-1}}$ is corrected for Galactic extinction of $E$ $B-V$ $=0.108$ and internal extinction of $E$ $B-V$ $=0.211$."
 We use the relation between star formation rate anc rest-[rame EUM. lux. (observed NUY flux) derived by ? assuming a Salpeter mass function., We use the relation between star formation rate and rest-frame FUV flux (observed NUV flux) derived by \citet{Salim2007} assuming a Salpeter mass function.
 This equation is calibrated for FUN. emission centered on1528A.. whils the rest-[rame wavelength. of the light observed. through w NUM filter is1," This equation is calibrated for FUV emission centered on, whilst the rest-frame wavelength of the light observed through the NUV filter is."
866: However 7 show that this wil rave a negligible elfect on the conversion between FUY luminosity and. star. formation rate., However \citet{Salim2007} show that this will have a negligible effect on the conversion between FUV luminosity and star formation rate.
 Star formation rates derived [rom the observed NUWV tux are eiven in Table 6.., Star formation rates derived from the observed NUV flux are given in Table \ref{tab:sfr}.
 2. derive a relation that results in ereater star formation rate for a given NUV us., \citet{Kennicutt98} derive a relation that results in greater star formation rate for a given NUV flux.
 The Ha and Ou] uminositv mav also be used to estimate the star formation rate in the BCC under the assumption that the eas is ionized by UV photons from hot voung stars., The $\alpha$ and ] luminosity may also be used to estimate the star formation rate in the BCG under the assumption that the gas is ionized by UV photons from hot young stars.
 reftab:str lists the total flux. extinction. corrected Duxes and luminosities of both Oii] and Hào. together with the derived star formation rates.," \\ref{tab:sfr} lists the total flux, extinction corrected fluxes and luminosities of both ] and $\alpha$, together with the derived star formation rates."
" Star formation rates (SETU are computed from the Oni] and Ho. Iuminosites according to the empirical relation of ?.. and the theoretical equation of ?.. respectively. where Lou and Lg, are expressed in units of ergs"," Star formation rates (SFR) are computed from the ] and $\alpha$ luminosites according to the empirical relation of \citet*{Kewley04}, and the theoretical equation of \citet{Kennicutt98}, respectively, where $L_{\rm \oii}$ and $L_{{\rm H}\alpha}$ are expressed in units of ${\rm erg\, s^{-1}}$."
 The EUV-derived star formation rates (calculated from the observed NUV. emission) are 23 times lower that the rates derived from the Lo and. ΟΠ) emission lines., The FUV-derived star formation rates (calculated from the observed NUV emission) are 2–3 times lower that the rates derived from the $\alpha$ and ] emission lines.
 In section 4.3.2. we show the number of O anc D stars present in the BCG is barely sullicient to power the Lla emission., In section \ref{sec:stellarpop} we show the number of O and B stars present in the BCG is barely sufficient to power the $\alpha$ emission.
 The discrepaney between the star formation rates derived. from. LUV and emission line luminosity reinforces this conclusion and implies that the Oj. Ho. and. other emission. lines are heated by an additional source.," The discrepancy between the star formation rates derived from FUV and emission line luminosity reinforces this conclusion and implies that the ], $\alpha$ and other emission lines are heated by an additional source."
 Star formation rates of DBCCGs derived. [rom line emission must only be considered upper limits of the true star formation rates., Star formation rates of BCGs derived from line emission must only be considered upper limits of the true star formation rates.
 We present a number of observations which imply. the, We present a number of observations which imply the
It is widely accepted that core-collapse supernovae (CC-SNo) represent the final explosive evolutionary phase of stars having initial (1.0. main sequence) nisses lareer than ~ (e.g.Woosley&Weaver1986:Tanz2003a:Teeeretal. 2003).,"It is widely accepted that core-collapse supernovae (CC-SNe) represent the final explosive evolutionary phase of stars having initial (i.e. main sequence) masses larger than $\sim$ \citep[e.g.][]{woosley86,hamuy03,heger03}."
. As such. CC-SNe are fundamental probes of the stellar evolution of massive stars and can be used to nuderstaud the link among explosion mechiauisui. nature of reciinants. progenitors evolution. auc euvironimoent around progenitors (e.g.Filippeuko1997:Cappellaro&Turatto 2000).," As such, CC-SNe are fundamental probes of the stellar evolution of massive stars and can be used to understand the link among explosion mechanism, nature of remnants, progenitors evolution, and environment around progenitors \citep[e.g.][]{filip97,CT00}."
.. Iu addition to their iutriusic interest. CC-SNe are relevant to many astrophysical issues.," In addition to their intrinsic interest, CC-SNe are relevant to many astrophysical issues."
 For example. the CC-SNe influence the plivsical evolution of galaxies. their frequency being related. to the on-goiug star formation rate (e.g.Cappellaroetal.2005:Madau.DellaValle.&Panagia1998) and determining the kinematics aud the structure of the interstellar mediun (e.g.Ratnatuuga&vaudenBerel 1989).," For example, the CC-SNe influence the physical evolution of galaxies, their frequency being related to the on-going star formation rate \citep*[e.g.][]{cappellaro05,madau98} and determining the kinematics and the structure of the interstellar medium \citep[e.g.][]{RB89}."
. CC-SNo are also important because of their role in the production of ueutriuos. cosmic rays and. probably. gravitational waves (e.g.Pagliarolietal. 2009).," CC-SNe are also important because of their role in the production of neutrinos, cosmic rays and, probably, gravitational waves \citep*[e.g.][]{haungs03,DK09,pagliaroli09}."
". Additionally. the CC-SN ejecta. αμσος in heavy elemeuts. male CC-SNe key actors iu the nucleosvuthesis processes of intermediate anc trans-iron clements and in the chemical evolution of galaxies (ee,Arnett1996:Chieff.Linonei.&Strauiero1998)."," Additionally, the CC-SN ejecta, enriched in heavy elements, make CC-SNe key actors in the nucleosynthesis processes of intermediate and trans-iron elements and in the chemical evolution of galaxies \citep*[e.g.][]{arnett96,chieffi98}."
 Moreover. they seem to be particularly promising to measure cosmological distances im addition to type Ia SNe (e.c.Nugentetal.2006:Zampicri2007:Olivaresctal.2010.andreferences therein)..," Moreover, they seem to be particularly promising to measure cosmological distances in addition to type Ia SNe \citep[e.g.][and references therein]{nugent06,zampieri07,olivares10}."
 Iu spite of the importance of these explosive events iu astroplivsics. there are still basic questions to be answered on C'C-SNe. related to the extreme varietv of their displays.," In spite of the importance of these explosive events in astrophysics, there are still basic questions to be answered on CC-SNe, related to the extreme variety of their displays."
 Indeed they appear to show ciffercut cherectics and to eject rather diverse amouuts of naterial. causing rather heterogeneous behaviors of hei light curves aud spectra (c.g.Turatto.Beuctti.&Pastorello 2007).," Indeed they appear to show different energetics and to eject rather diverse amounts of material, causing rather heterogeneous behaviors of their light curves and spectra \citep*[e.g.][]{turatto07}."
. This apparent heterogeneity may 0. linked το stellar evolution effects and to the explosion iuechanisni (c.e.Punooetal.2009.audreferences therein). but the exact link between the physical properties of the explosion (ejected mass. explosion energv. stellar structure and composition at he explosion) aud the observational characteristics is far roni being well-established.," This apparent heterogeneity may be linked to stellar evolution effects and to the explosion mechanism \citep[e.g.][and references therein]{pumo09b}, but the exact link between the physical properties of the explosion (ejected mass, explosion energy, stellar structure and composition at the explosion) and the observational characteristics is far from being well-established."
 Light curve and spectral modelling lave often stuccessfilly been used to provide information about the shivsical properties of sinele CC-SNe (c.g.SNe1987A.vento.&Παιν 2011).," Light curve and spectral modelling have often successfully been used to provide information about the physical properties of single CC-SNe \citep*[e.g. SNe 1987A, 1993J, 1999em, 2004et, 
and 2005cs; see][]{shige88,woosley88,shigenomoto90,blinnikov98,utrobin04,utrobin07,baklanov05,
UC08,pasto09,bersten11}."
.. Comparatively little effort has been devoted to investigating huge samples of these events and trviug to explain the significant range of properties that they show., Comparatively little effort has been devoted to investigating large samples of these events and trying to explain the significant range of properties that they show.
 Differences im huuinosity. expansion velocity and vields are very large (up to ~ 2 order of magnitude). aud difficult to relate to changes im a single parameter. such as the mass of the progenitor star.," Differences in luminosity, expansion velocity and yields are very large (up to $\sim$ 2 order of magnitude), and difficult to relate to changes in a single parameter, such as the mass of the progenitor star."
 However. despite this. variations m the observed properties appear to obey a certain order.," However, despite this, variations in the observed properties appear to obey a certain order."
 Some correlations amone differeut observalles have Όσοι noted (e.g.Tamu20035:Ziiipieri2007) and used to calibrate," Some correlations among different observables have been noted \citep[e.g.][]{hamuy03b,zampieri07} and used to calibrate"
of the temperature and thermalization depth at 0.75 keV for model m2rlcold.,of the temperature and thermalization depth at 0.75 keV for model m2r1cold.
" Before shock breakout the location of the thermalization layer is the transition region between the wind and the progenitor star, where the density increases rapidly."," Before shock breakout the location of the thermalization layer is the transition region between the wind and the progenitor star, where the density increases rapidly."
" Emission from these regions is, however, negligible in the XRT band due to the relatively low temperatures there."," Emission from these regions is, however, negligible in the XRT band due to the relatively low temperatures there."
" After shock breakout the thermalization layer lies in between the forward and reverse shocks, the region comprise of shocked, accelerated wind material."," After shock breakout the thermalization layer lies in between the forward and reverse shocks, the region comprise of shocked, accelerated wind material."
" Light curves of explosion models m2rlcold and m2rlhot (the smaller progenitor) are shown in Figures 12 and 13 for three viewing angles: along the equator (0 degrees), 45 degrees, and the axis of (90 degrees)."," Light curves of explosion models m2r1cold and m2r1hot (the smaller progenitor) are shown in Figures \ref{fig:lc_m2r1cold} and \ref{fig:lc_m2r1hot} for three viewing angles: along the equator (0 degrees), 45 degrees, and along the axis of symmetry (90 degrees)."
 The widths of the alonglight curves are roughly symmetry100 seconds at all angles for both models and the total X-ray burst time is around 500 to 600 seconds., The widths of the light curves are roughly 100 seconds at all angles for both models and the total X-ray burst time is around 500 to 600 seconds.
" These values are very close to what was observed for SN 2008D, shown in Figure 8.."," These values are very close to what was observed for SN 2008D, shown in Figure \ref{fig:sn08Dlc}."
 Examining the shape of the model light curves shows that for 0 degree viewing angles the light curves rise very quickly., Examining the shape of the model light curves shows that for 0 degree viewing angles the light curves rise very quickly.
 This is due to two effects., This is due to two effects.
 The first is simply because the emission from both poles is visible., The first is simply because the emission from both poles is visible.
 The second is due to the nature of an aspherical breakout and the effects of light travel corrections., The second is due to the nature of an aspherical breakout and the effects of light travel corrections.
" As the bipolar shocks continue to erupt from the surface, the brightest emission is from where the progenitorshocks are just reaching the surface coming(essentially two rings moving across the stellar surface from the poles to the equator)."," As the bipolar shocks continue to erupt from the progenitor surface, the brightest emission is coming from where the shocks are just reaching the surface (essentially two rings moving across the stellar surface from the poles to the equator)."
" Thus, the brightest emitting regions are moving toward the observer a of emission once rapidlylight travel time corrections causingare made."," Thus, the brightest emitting regions are moving rapidly the observer causing a pile up of emission once light travel time corrections are made."
"pile upThis effect is reduced, or eliminated, at higher viewing angles because the brightly emitting rings are no longer moving so much toward the observer."," This effect is reduced, or eliminated, at higher viewing angles because the brightly emitting rings are no longer moving so much toward the observer."
" Table 2 lists the resulting light curve characteristics for our models m2rlcold and m2rlhot, as well as m2rlsph for comparison."," Table \ref{table:lc} lists the resulting light curve characteristics for our models m2r1cold and m2r1hot, as well as m2r1sph for comparison."
" As expected, the spherical explosion has a higher peak luminosity, but a much shorter FWHM and overall burst time, Ar."," As expected, the spherical explosion has a higher peak luminosity, but a much shorter FWHM and overall burst time, $\Delta t$."
 The width of the light curve for m2rlsph is set by the light crossing time of the progenitor., The width of the light curve for m2r1sph is set by the light crossing time of the progenitor.
" The simulated XRT count rates for models m2rlhot and, m2r1cold are very similar to those we find for XRO especially,080109/SN 2008D. The light curve rise time and FWHM for model m2r1cold at a viewing angle of 45 degrees (red dashed curve in Figure 12)) is a good match to SN 2008D. The X-ray spectra for models m2rlcold and m2rlhot are shown in Figure 17."," The simulated XRT count rates for models m2r1hot and, especially, m2r1cold are very similar to those we find for XRO 080109/SN 2008D. The light curve rise time and FWHM for model m2r1cold at a viewing angle of 45 degrees (red dashed curve in Figure \ref{fig:lc_m2r1cold}) ) is a good match to SN 2008D. The X-ray spectra for models m2r1cold and m2r1hot are shown in Figure \ref{fig:spec_m2r1}."
. These spectra are corrected for the detector response function and for X-ray absorpotion (assuming Ny=1.7x107°)., These spectra are corrected for the detector response function and for X-ray absorpotion (assuming $N_{\rm H} = 1.7\times10^{20}$ ).
" They are also averaged over the burst time, Af, as was done for the observations of XRO 080109/SN 2008D (seeFigure9;;??)."," They are also averaged over the burst time, $\Delta t$, as was done for the observations of XRO 080109/SN 2008D \citep[see Figure \ref{fig:sn08Dspec};."
 The shapes of the spectra at the different viewing angles are very similar., The shapes of the spectra at the different viewing angles are very similar.
 Indeed there is not much difference between the two different models., Indeed there is not much difference between the two different models.
 The time-averaging washes away the major differences that are apparent in the light curves., The time-averaging washes away the major differences that are apparent in the light curves.
 The shape of the spherical explosion spectra are also very similar., The shape of the spherical explosion spectra are also very similar.
" They are generally brighter, but this is due to a shorter averaging time."," They are generally brighter, but this is due to a shorter averaging time."
 The spectra in each case are softer than the spectrum of XRO 080109 (see Figure 9))., The spectra in each case are softer than the spectrum of XRO 080109 (see Figure \ref{fig:sn08Dspec}) ).
 The light curves for the simulations in the larger (Figures 14 and 15)) are characterized by muchprogenitor longer time scales and overall less bright emission., The light curves for the simulations in the larger progenitor (Figures \ref{fig:lc_m7r6cold} and \ref{fig:lc_m7r6hot}) ) are characterized by much longer time scales and overall less bright emission.
 The FWHM for models m7r6cold and m7r6hot range from around 100 to 200 seconds at various angles while the total burst times are from 300 to 1000 seconds (see Table 2))., The FWHM for models m7r6cold and m7r6hot range from around 100 to 200 seconds at various angles while the total burst times are from 300 to 1000 seconds (see Table \ref{table:lc}) ).
 The peak luminosities and total radiated energies are less than in the analogous explosions in the smaller , The peak luminosities and total radiated energies are less than in the analogous explosions in the smaller progenitor.
"While the ratio of explosion energies to ejectaprogenitor. masses are roughly equivalent across all simulations, the reduced luminosities in the larger progenitor can be explained by a slightly lower shock the burst and a lower wind density at the radius of velocityshock duringbreakout (i.e., the radius of the progenitor star)."," While the ratio of explosion energies to ejecta masses are roughly equivalent across all simulations, the reduced luminosities in the larger progenitor can be explained by a slightly lower shock velocity during the burst and a lower wind density at the radius of shock breakout (i.e., the radius of the progenitor star)."
" As we discuss in more detail in Section 4.2,, the density of the wind plays an important role in the strength of the X-ray emission because the thermalization depths during the bursts lie in the shocked wind."," As we discuss in more detail in Section \ref{sec:wind}, the density of the wind plays an important role in the strength of the X-ray emission because the thermalization depths during the bursts lie in the shocked wind."
 The light curve for m7r6cold is similar in shape to m2rlcold and m2rlhot but with longer timescales., The light curve for m7r6cold is similar in shape to m2r1cold and m2r1hot but with longer timescales.
" Model m7r6hot, however, exhibits a dramatically double-peaked light curve."," Model m7r6hot, however, exhibits a dramatically double-peaked light curve."
" As discussed in Section 2,, we ran simulations of model m2rlcold at three different resolutions."," As discussed in Section \ref{sec:simulations}, we ran simulations of model m2r1cold at three different resolutions."
 Figure 19 shows the curves for these three simulations plotted together., Figure \ref{fig:m2r1cold_res} shows the X-ray light curves for these three simulations plotted together.
 X-rayThe light curves of the low- and fiducial resolution simulations are extremely similar., The light curves of the low- and fiducial resolution simulations are extremely similar.
 The high-resolution case varies only slightly from the lower-resolution light curves., The high-resolution case varies only slightly from the lower-resolution light curves.
" The early peak is brighter, due to the shock speeds just after breakout (see 4.2)). The sl"," The early peak is brighter, due to the slightly greater shock speeds just after breakout (see \ref{sec:wind}) )."
"ightlysecond greaterpeak is slightly dimmer, due to the lesser degree of extension of the southern shock structure in the high-resolution simulation, leading to a smaller emitting area as seen from a viewing angle of 0 degrees."," The second peak is slightly dimmer, due to the lesser degree of extension of the southern shock structure in the high-resolution simulation, leading to a smaller emitting area as seen from a viewing angle of 0 degrees."
 The absorptive opacities depend strongly on the metallicity of the absorbing gas., The absorptive opacities depend strongly on the metallicity of the absorbing gas.
 To illustrate this we have calculated, To illustrate this we have calculated
planetary systems due to tidal interactions between planets or protoplanets and a gaseous disc. in which the whole system is still embedded.,"planetary systems due to tidal interactions between planets or protoplanets and a gaseous disc, in which the whole system is still embedded."
 The different mass objects. which we expect to [ind in forming planctary systems. will migrate with cdillerent. rates.," The different mass objects, which we expect to find in forming planetary systems, will migrate with different rates."
 The final configurations will depend on the intricate interplay among many physical processes including planet-planet. disc-planet. and. planet-star interactions.," The final configurations will depend on the intricate interplay among many physical processes including planet-planet, disc-planet and planet-star interactions."
 One scenario is that the convergent migration brings the giant planets closer to cach other and they can become locked in low order commoensurability (Drvdenetal.2000:Wles2000) as it is observed at least in three multiplanet. svstems (GJ S76. 11D. 82943 and 55 ςπο)," One scenario is that the convergent migration brings the giant planets closer to each other and they can become locked in low order commensurability \citep{bryden, kley} as it is observed at least in three multiplanet systems (GJ 876, HD 82943 and 55 Cnc)."
" Also. the low mass planets can undergo the convergent migration and form a resonant structure '""apaloizou& Szuszkiewicz2005)."," Also, the low mass planets can undergo the convergent migration and form a resonant structure \citep{papszusz}."
. Phe pulsar. planets around PSR D1257|12 might be an outcome of such scenario., The pulsar planets around PSR B1257+12 might be an outcome of such scenario.
 The dillerences in the migration rates of low and high mass planets may also lead to the convergent migration (llahn&Ward1906)., The differences in the migration rates of low and high mass planets may also lead to the convergent migration \citep{hawa}.
 We will consider this ἵνρο of convergent migration here as a mechanism. for capture of a Super-Exrth by a Jupiter into mean motion commeoensurabilities anc we will follow the pair of planets further in their evolution to establish the stability of their configuration., We will consider this type of convergent migration here as a mechanism for capture of a Super-Earth by a Jupiter into mean motion commensurabilities and we will follow the pair of planets further in their evolution to establish the stability of their configuration.
 Ht has been found already by Phommes(2005) that for typical protoplanetary disc parameters the low mass planets can be captured by massive planet into exterior resonances.," It has been found already by \citet{thommes}
 that for typical protoplanetary disc parameters the low mass planets can be captured by massive planet into exterior resonances."
 The resonant captures play also an important role in moclelline terrestrial planet. formation ane their survival in the presence of migrating gas giant., The resonant captures play also an important role in modelling terrestrial planet formation and their survival in the presence of migrating gas giant.
 Phe early divergent conclusions on the occurrence of terrestrial planets in hot Jupiter systems (Levisonctal.2001:Armitage2003:ltavmondetal.2005:Alancdell&Sigurdsson2003). have been clarified by most recent studies (Fogg&Nelson2005.etal.2007) which predict that terrestrial planets can grow and be retained in the hot-Jupiter svstems both interior and exterior of the gas giant.," The early divergent conclusions on the occurrence of terrestrial planets in hot Jupiter systems \citep{levison, armitage, 
raymond05, mansig} have been clarified by most recent studies \citep{fonel05, 
fonel07a, fonel07b, zhou, raymond, mandell} which predict that terrestrial planets can grow and be retained in the hot-Jupiter systems both interior and exterior of the gas giant."
 The most relevant feature of these investigations for the present study is the possibility of the existence of terrestrial planets on the internal orbits relative to Jupiter., The most relevant feature of these investigations for the present study is the possibility of the existence of terrestrial planets on the internal orbits relative to Jupiter.
 Such planets could become captured in the mean motion commoensurability which is then maintained during their evolution so an outcome in this case will be the Jupiter and the low-mass planet migrating together towards the star., Such planets could become captured in the mean motion commensurability which is then maintained during their evolution so an outcome in this case will be the Jupiter and the low-mass planet migrating together towards the star.
 Our aim here is to consider in fall cetail the evolution of the close pair of planets: a gas giant on the external orbit and a Super-Earth on the internal one in the voung planetary system. when both planets are embedded in a gaseous disc.," Our aim here is to consider in full detail the evolution of the close pair of planets: a gas giant on the external orbit and a Super-Earth on the internal one in the young planetary system, when both planets are embedded in a gaseous disc."
 We have performed 2D hydrodynamical simulations in order to determine the occurrence of the first order mean motion resonances in such svstems and at the same time to examine the possible planet configurations as the outcome of the convergent orbital migration., We have performed 2D hydrodynamical simulations in order to determine the occurrence of the first order mean motion resonances in such systems and at the same time to examine the possible planet configurations as the outcome of the convergent orbital migration.
 Our paper is organized as follows., Our paper is organized as follows.
 In Section 2 we summarize shortly few facts about two types of migration. which we will use in the present study.," In Section 2 we summarize shortly few facts about two types of migration, which we will use in the present study."
 In Section 3 we discuss the Jupiter anc Super-Earth svstem in the context of the restricted three body problem when the disc is absent., In Section 3 we discuss the Jupiter and Super-Earth system in the context of the restricted three body problem when the disc is absent.
 Aloreover we have performed simple N-body calculations in order to determine the Super-Earth dynamics in a presence of Jupiter and compare it with the classical. results of celestial mechanies., Moreover we have performed simple N-body calculations in order to determine the Super-Earth dynamics in a presence of Jupiter and compare it with the classical results of celestial mechanics.
 This provides us with a well studied and understood. framework in which we can present the results of our hyvdrodynamical. simulations of a voung planetary system. where planets are still embedded in a gaseous clise.," This provides us with a well studied and understood framework in which we can present the results of our hydrodynamical simulations of a young planetary system, where planets are still embedded in a gaseous disc."
 In Section 4 we describe our hvdrodynamical simulations and show the planetary orbit evolution together with the changes in the protoplanetary disc structure., In Section 4 we describe our hydrodynamical simulations and show the planetary orbit evolution together with the changes in the protoplanetary disc structure.
 Phe discussion and our conclusions are given in Section 5., The discussion and our conclusions are given in Section 5.
 Migration due to planct-cise interaction might plav an important role in shaping up planet configurations in the Xdanetarv svstems as we have already mentioned. in the »evious section., Migration due to planet-disc interaction might play an important role in shaping up planet configurations in the planetary systems as we have already mentioned in the previous section.
 From the dynamical. point of. view one of the most important. consequences of this process is an occurrence ofplanets locked in mean-motion resonances and his will be our main concern here., From the dynamical point of view one of the most important consequences of this process is an occurrence of planets locked in mean-motion resonances and this will be our main concern here.
" The migration rates for dilferent planet masses has en estimated by number of authors. see the review by ""apaloizouetal.(2006)."," The migration rates for different planet masses has been estimated by number of authors, see the review by \cite{fivebig}."
. Pheir results have been illustrated in Fig., Their results have been illustrated in Fig.
 1 where we plot the migration time of a planet as a unction of its mass., \ref{fig1} where we plot the migration time of a planet as a function of its mass.
" There are two mass regimes which are of interest here. namely (0.1 - 3034, ) and (150 - 150034, ) or which in a tvpical protoplanetary disc we can talk about wo different types of migration. called. simply. tvpe Ll and ype LL respectively."," There are two mass regimes which are of interest here, namely (0.1 - $M_{\oplus}$ ) and (150 - $M_{\oplus}$ ) for which in a typical protoplanetary disc we can talk about two different types of migration, called simply type I and type II respectively."
" The migration time for low mass planets cnibeclelect in a gascous disc (tvpeE migration) has been derived by in the form lere my is mass of the planet. ry is the distance from the central star AM. M, is the disc surface density. £f/r ancl €, are the disc aspect ratio and angular velocity respectively."," The migration time for low mass planets embedded in a gaseous disc (typeI migration) has been derived by \cite{tanaka02} in the form Here $m_p$ is mass of the planet, $r_p$ is the distance from the central star $M_*$, $\Sigma_p$ is the disc surface density, $H/r$ and $\Omega_p$ are the disc aspect ratio and angular velocity respectively."
 The coellicient + depends on the cise surface density. profile. which is expressed as V(r)xr.," The coefficient $\gamma$ depends on the disc surface density profile, which is expressed as $\Sigma(r) \propto r^{-\gamma}$."
" Assuming ~=0 for the lat surface density distribution. V,=2000 keg/m. H/»=0.05 andr,=5.24€ we have calculated the migration tine as a function of the planet mass and draw it in Fig. 1.."," Assuming $\gamma=0$ for the flat surface density distribution, $\Sigma_p= 2000$ $^2$, $ H/r =0.05$ and $r_p=5.2 AU$ we have calculated the migration time as a function of the planet mass and draw it in Fig. \ref{fig1}."
 Type HE migrators open a gap in the cise and. their evolution is determined. by the racial velocity drift in the cdisc where f£ ds a kinematic viscosity parameter., Type II migrators open a gap in the disc and their evolution is determined by the radial velocity drift in the disc where $\nu$ is a kinematic viscosity parameter.
 The migration time can be estimated as (Lin&Papaloizou1993) lt has been shown in Fig., The migration time can be estimated as \citep{linpap93} It has been shown in Fig.
" 1. for r,=5.240 and clilferent values of vw (lO7.2:I07.3IO.7.910.7 and 10 "") expressed in the dimensionless units. cliscussed in Section 4.1.."," \ref{fig1} for $r_p=5.2 AU$ and different values of $\nu$ $10^{-5},
2\cdot 10^{-5}, 3\cdot 10^{-5}, ..., 9\cdot 10^{-5}$ and $10^{-6}$ ) expressed in the dimensionless units discussed in Section \ref{section41}. ."
 Phe migration time in this formulation does not depend on the planet mass., The migration time in this formulation does not depend on the planet mass.
 In drawing lines for a given viscosity parameter we have taken into account the condition [or a gap opening in the dise which reads, In drawing lines for a given viscosity parameter we have taken into account the condition for a gap opening in the disc which reads
Current data from the ΑςΠΟ collaboration (Alcocketal.1997) indicate that in the context of a spherical isothermal model with a Alaxwellian velocity cistribution. some significant fraction of the Galactic halo is composed of NLACIIOs with masses roughly in the range 0.1 to 1.0 Al..,"Current data from the MACHO collaboration \cite{MACHOmass} indicate that in the context of a spherical isothermal model with a Maxwellian velocity distribution, some significant fraction of the Galactic halo is composed of MACHOs with masses roughly in the range 0.1 to 1.0 $\Msol$."
. Such masses are consistent with several astrophysical candidates for ALACΟν white chwarfs. neutron stars. and black holes — each of which presents serious challenges for stellar formation and evolution theories.," Such masses are consistent with several astrophysical candidates for MACHOs – white dwarfs, neutron stars, and black holes – each of which presents serious challenges for stellar formation and evolution theories."
 However. the MACHO component of the halo. if it is not the major component. as in Cold. Dark Matter scenarios. may have a very dilferent. distribution from the typically assumed spherical isothermal model.," However, the MACHO component of the halo, if it is not the major component, as in Cold Dark Matter scenarios, may have a very different distribution from the typically assumed spherical isothermal model."
 Phe ΑςΠΟ distribution may be in a significantly. Lattened halo. and/or. due to dissipation. more centrally. condensed.," The MACHO distribution may be in a significantly flattened halo and/or, due to dissipation, more centrally condensed."
 In. addition. such a distribution might have a significant rotational component.," In addition, such a distribution might have a significant rotational component."
 These possibilities have strong implications not only for the MACHO fraction of the halo. but also for the mass estimates derived from the event clurations.," These possibilities have strong implications not only for the MACHO fraction of the halo, but also for the mass estimates derived from the event durations."
 The velocity. dispersion and. mass density. distribution which describe the halo model are crucial input. parameters in extracting an estimate of the lens mass from the data., The velocity dispersion and mass density distribution which describe the halo model are crucial input parameters in extracting an estimate of the lens mass from the data.
 The event duration is given by the radius of the Einstein ring divided by the NLACIHIO velocity transverse to the line of sight. where the Einstein ring radius is a function of the position of the NLACTIIO along the line of sight. ancl the lens mass.," The event duration is given by the radius of the Einstein ring divided by the MACHO velocity transverse to the line of sight, where the Einstein ring radius is a function of the position of the MACHO along the line of sight and the lens mass."
 Phe masses of the lenses can only be determined statistically. in the context of an a priori assumption about the distribution and. velocity dispersion of the lenses.," The masses of the lenses can only be determined statistically, in the context of an a priori assumption about the distribution and velocity dispersion of the lenses."
 Such an analysis. using a spherical isothermal niocdel vields a central mass estimate of zzO4AL..," Such an analysis, using a spherical isothermal model yields a central mass estimate of $\approx 0.4\Msol$."
 llowever. as discussed above. the NLACTIO distribution ancl velocity dispersion. may be very. dillerent. from. that assumed in the standard halo model.," However, as discussed above, the MACHO distribution and velocity dispersion may be very different from that assumed in the standard halo model."
 Earlier work (αν&Gates1998). in exploring models with a highlv llattened halo and a bulk rotational component has suggested that à very highlv condensed model for the NLACTIO clistribution. such as a thick disk. might reduce the NLACTIO mass estimate from the current data to a level consistent with brown dwarl candidates.," Earlier work \cite{us_rotate} in exploring models with a highly flattened halo and a bulk rotational component has suggested that a very highly condensed model for the MACHO distribution, such as a thick disk, might reduce the MACHO mass estimate from the current data to a level consistent with brown dwarf candidates."
 Previous explorations (Gatesetal.1998). have indicated that such models may be able to reproduce the observed optical cepths toward the Large Magellanic Cloud (LM?) and the Galactic bulge., Previous explorations \cite{us_nomacho} have indicated that such models may be able to reproduce the observed optical depths toward the Large Magellanic Cloud (LMC) and the Galactic bulge.
 The overall shape of the Galactic halo is unknown., The overall shape of the Galactic halo is unknown.
 Attempts to determine the shape of ealactic halo potentials from Uarine of the outer Galactic gas laver (Ollinge1996:Sackettetal.1994). point to a Uattened halo: Hattening of the potential is also supported by simulations of the cold dark matter halo formation (DubinskiJ.&Carlberg. 1991).," Attempts to determine the shape of galactic halo potentials from flaring of the outer Galactic gas layer \cite{Olling,Sackett} point to a flattened halo; flattening of the potential is also supported by simulations of the cold dark matter halo formation \cite{nbody}."
. While it is highly unlikely that the entire halo is in a clisk-like configuration. it is not unreasonable to assume that the ΑΕΛΠΟ component of the halo is significantly. more llattened than the dark matter halo.," While it is highly unlikely that the entire halo is in a disk-like configuration, it is not unreasonable to assume that the MACHO component of the halo is significantly more flattened than the dark matter halo."
 The condensation of a gascous halo component to form a very thick disk is likely to result in significant star formation and thus the production of a population of lens candidates., The condensation of a gaseous halo component to form a very thick disk is likely to result in significant star formation and thus the production of a population of lens candidates.
 In this paper we explore in detail the consequences, In this paper we explore in detail the consequences
"frame rate and the need to correctly sample the fringe at the highest frequency fj,,,=6v/..",frame rate and the need to correctly sample the fringe at the highest frequency $f_{max}=6v/\lambda$.
 In the observations presentec here. the velocity was set to v=7Oym/s. which corresponds to fringes at 80. 160. and 240 Hz.," In the observations presented here, the velocity was set to $v=70\,\mu$ m/s, which corresponds to fringes at $80$ , $160$, and $240\,$ Hz."
 This is just fast enough to freeze the turbulent optical path delay (OPD) under mediar atmospheric conditions., This is just fast enough to freeze the turbulent optical path delay (OPD) under median atmospheric conditions.
 The outputs of the chip are dispersec over seven spectral channels across the H-band. providing low (R= 35) spectral resolution.," The outputs of the chip are dispersed over seven spectral channels across the H-band, providing low $R=35$ ) spectral resolution."
 During the observation. each seu IS processed by a quick-look algorithm that also implements slow group tracking.," During the observation, each scan is processed by a quick-look algorithm that also implements slow group tracking."
 All the data have been reduced with the pndrs package described 11 ?.., All the data have been reduced with the $\texttt{pndrs}$ package described in \citet{LeBouquin11}.
 For each file and each spectral channel. the pndrs package provides six visibility and four closure phase measurements.," For each file and each spectral channel, the $\texttt{pndrs}$ package provides six visibility and four closure phase measurements."
 The statistical uncertainties are estimated by the dispersion over the 100 scans contained in each file. and typically range from 0.25 to 3 degrees for the closure phase. depending on the target brightness and atmospheric conditions.," The statistical uncertainties are estimated by the dispersion over the 100 scans contained in each file, and typically range from 0.25 to 3 degrees for the closure phase, depending on the target brightness and atmospheric conditions."
 The statistical uncertainty includes. contributions from. the classical detection noises (detector and photon noise). as well as from the atmospheric OPD fluctuations faster than the scan rates.," The statistical uncertainty includes contributions from the classical detection noises (detector and photon noise), as well as from the atmospheric OPD fluctuations faster than the scan rates."
The instrumental and atmospheric contributions to the visibilities and closure phases (the so-called transfer function) is Monitored by interleaving the observation of science stars with calibration stars., The instrumental and atmospheric contributions to the visibilities and closure phases (the so-called transfer function) is monitored by interleaving the observation of science stars with calibration stars.
 Some nights show an extremely stable transfer function. down to 0.1 deg. while other nights require fitting the transfer function by a drift. indicating that we are facing some non-stationary biases.," Some nights show an extremely stable transfer function, down to $0.1\,$ deg, while other nights require fitting the transfer function by a drift, indicating that we are facing some non-stationary biases."
 This aspect requires additional commissioning before being fully understood., This aspect requires additional commissioning before being fully understood.
 In this section. we discuss how the PIONIER calibrated data set is used to search for faint companions. compute confidence levels for tentative detections. and derive upper limits to the presence of off-axis companions in case of non-detections.," In this section, we discuss how the PIONIER calibrated data set is used to search for faint companions, compute confidence levels for tentative detections, and derive upper limits to the presence of off-axis companions in case of non-detections."
 Our analysis is based solely on the calibrated closure phases. which have two significant advantages over (differential) visibilities and (differential) phases.," Our analysis is based solely on the calibrated closure phases, which have two significant advantages over (differential) visibilities and (differential) phases."
 On the one hand. at first order. closure phases are not affected by phase disturbances ahead of the beam combiner (seee.g..?)..," On the one hand, at first order, closure phases are not affected by phase disturbances ahead of the beam combiner \citep[see e.g.,][]{Monnier03}."
 Therefore. closure phases are not affected by atmospheric turbulence and other instrumental effects such as mechanical. vibrations. which strongly reduce the accuracy of other observable quantities.," Therefore, closure phases are not affected by atmospheric turbulence and other instrumental effects such as mechanical vibrations, which strongly reduce the accuracy of other observable quantities."
 On the other hand. closure phase are especially sensitive to faint off-axis companions. as they differ from zero only for non point-symmetric targets.," On the other hand, closure phase are especially sensitive to faint off-axis companions, as they differ from zero only for non point-symmetric targets."
 One of the associated benefits is that the detection of a faint companion does not depend on the knowledge of the parent star’s photospheric parameters. provided that it is point-symmetric: however. poor knowledge of the primary stellar diameter may lead to larger ucertainty on the photometry of a potential companion.," One of the associated benefits is that the detection of a faint companion does not depend on the knowledge of the parent star's photospheric parameters, provided that it is point-symmetric; however, poor knowledge of the primary stellar diameter may lead to larger uncertainty on the photometry of a potential companion."
 Our companion search strategy is based on the conputation of the y distance betweer our data set and several models where a faint companion ts added to a user-specified photospheric model for the central star., Our companion search strategy is based on the computation of the $\chi^2$ distance between our data set and several models where a faint companion is added to a user-specified photospheric model for the central star.
 The y goodness-of-fit is computed for a series of angular separations along east and north between the primary star and its companion., The $\chi^2$ goodness-of-fit is computed for a series of angular separations along east and north between the primary star and its companion.
 The first step in this analysis is therefore to define the region around the star in which companions will be searched for., The first step in this analysis is therefore to define the region around the star in which companions will be searched for.
 There are three limitations to the maximum angular separation at which a companion can be searchec for in the PIONIER data., There are three limitations to the maximum angular separation at which a companion can be searched for in the PIONIER data.
 The first one is related to the use of single-mode fibres in the instrument., The first one is related to the use of single-mode fibres in the instrument.
 The coupling efficiency of point-like sources into a single-mode fibre depends on the angular separation between the source and the optical axis., The coupling efficiency of point-like sources into a single-mode fibre depends on the angular separation between the source and the optical axis.
 The transmission profile. which results from the overlap integral between the turbulent image of a point-like source and the fundamental mode of the fibre can be approximated by a Gaussian profile. whose full width at half maximum (FWHM) depends on the fibre parameters. the telescope diameter. the observing wavelength. and the strength of the atmospheric turbulence.," The transmission profile, which results from the overlap integral between the turbulent image of a point-like source and the fundamental mode of the fibre can be approximated by a Gaussian profile, whose full width at half maximum (FWHM) depends on the fibre parameters, the telescope diameter, the observing wavelength, and the strength of the atmospheric turbulence."
 In the case of PIONIER. working with the 1.8mm ATs under typical atmospheric conditions (seeing of 05 in the visible). the FWHM of the Gaussian profile is about mmas.," In the case of PIONIER, working with the m ATs under typical atmospheric conditions (seeing of $0\farcs8$ in the visible), the FWHM of the Gaussian profile is about mas."
 The transmitted flux of off-axis companions located at angular separations larger than mmas from the central will therefore be attenuated by more than5065., The transmitted flux of off-axis companions located at angular separations larger than mas from the central will therefore be attenuated by more than.
 A second limitation to the angular distance of detectable companions comes from the separation of the fringe packets associated to the two stars in terms of optical path delay (OPD)., A second limitation to the angular distance of detectable companions comes from the separation of the fringe packets associated to the two stars in terms of optical path delay (OPD).
 To properly detect a potential companion. we ask the fringe packet separation to be less than half the size of the smallest scan. which amounts to 40rm for the baselines with the slowest scanning speed (?)..," To properly detect a potential companion, we ask the fringe packet separation to be less than half the size of the smallest scan, which amounts to $40\,\mu$ m for the baselines with the slowest scanning speed \citep{LeBouquin11}."
 The OPD separation of the fringe packets is directly related to the angular separation AG of the two objects in the sky: AOPD=BA0cos&. with Bcos8 the projected baseline length.," The OPD separation of the fringe packets is directly related to the angular separation $\Delta\theta$ of the two objects in the sky: $\Delta{\rm OPD} = B \, \Delta\theta \, \cos\theta$, with $B \, \cos \theta$ the projected baseline length."
" Using a mean projected baseline of about mm for the compact configuration (E0-GO-HO-I. used for most targets) and taking AOPD,ja,= 204m. we end up with AG,=100 mmas."," Using a mean projected baseline of about m for the compact configuration (E0-G0-H0-I1, used for most targets) and taking $\Delta{\rm OPD}_{\rm max}=20\,\mu$ m, we end up with $\Delta\theta_{\rm max} \simeq 100$ mas."
" A fringe packet separation smaller than 20, also ensures that the two fringe packets are (at least partially) superposed in all cases. as the size of the fringe packet is UL/AAz 63m when the H-band signal is dispersed onto seven spectral channels."," A fringe packet separation smaller than $20\,\mu$ m also ensures that the two fringe packets are (at least partially) superposed in all cases, as the size of the fringe packet is $\lambda^2/\Delta\lambda \simeq 63\,\mu$ m when the H-band signal is dispersed onto seven spectral channels."
 The last limitation to the size of the search region comes from the spectral sampling of the closure phase signal., The last limitation to the size of the search region comes from the spectral sampling of the closure phase signal.
 Because the angular resolution of the array is proportional to the observing wavelength. the closure phase signature of an off-axis companion oscillates around zero as a function. of wavelength.," Because the angular resolution of the array is proportional to the observing wavelength, the closure phase signature of an off-axis companion oscillates around zero as a function of wavelength."
 These oscillations must be properly sampled by the spectral resolution of the instrument in order to produce a unique (non-aliased). solutior in the fitting process., These oscillations must be properly sampled by the spectral resolution of the instrument in order to produce a unique (non-aliased) solution in the fitting process.
" As discussed by 9?.. the period in the closure phase signal. which is roughly given by P,=A-/(OB4AG.—2). must be larger than four times the spectral charnel size."," As discussed by \citet{Absil10a}, the period in the closure phase signal, which is roughly given by $P_{\lambda} = \lambda^2 / (B\,\Delta\theta-\lambda)$, must be larger than four times the spectral channel size."
 In the present case. considering a mean sky-projected baseline B= 430mm for the compact configuration (the nost frequentlyused in our observations) and taking the charnel width of 0.045 jim in the PIONIER dispersed data into account. this constraint translates," In the present case, considering a mean sky-projected baseline $B=40$ m for the compact configuration (the most frequentlyused in our observations) and taking the channel width of $0.045\,\mu$ m in the PIONIER dispersed data into account, this constraint translates"
mass estimates are only based on the dust.,mass estimates are only based on the dust.
 On the other hand. there is certainly enough circum-sub-stellar material to form less massive planets.," On the other hand, there is certainly enough circum-sub-stellar material to form less massive planets."
 Our mm observations for brown chwarls extend (he mass range of objects with constraints on disk masses down to ~0.02Af. (see Fig. 3)).," Our mm observations for brown dwarfs extend the mass range of objects with constraints on disk masses down to $\sim 0.02\,M_{\odot}$ (see Fig. \ref{f3}) )."
 In our age-spread corrected sample. we do nol see any significant change of relative disk mass with object mass. as claimed previously in (he literature.," In our age-spread corrected sample, we do not see any significant change of relative disk mass with object mass, as claimed previously in the literature."
 The dominant feature in this plot is a large scatter of at least (wo orders of magnitude. independent of the mass of the central object.," The dominant feature in this plot is a large scatter of at least two orders of magnitude, independent of the mass of the central object."
 Apparently. the disk masses only scale down with object mass.," Apparently, the disk masses only scale down with object mass."
 As outlined in Sect. 3..," As outlined in Sect. \ref{diskmass},"
 binarity may plaw à more importaat role than central object mass for the evolution of the disk., binarity may play a more important role than central object mass for the evolution of the disk.
 It is a pleasure to acknowledge the support from the IRAAI pool observing team. in parüceular Stephane Leon.," It is a pleasure to acknowledge the support from the IRAM pool observing team, in particular Stephane Leon."
 Christina Walker and Mark O'Sullivan kindly helped with the modeling section of this paper., Christina Walker and Mark O'Sullivan kindly helped with the modeling section of this paper.
 We thank Mirza Ahmic lor assisting us in the selection of Spitzer images for our targets., We thank Mirza Ahmic for assisting us in the selection of Spitzer images for our targets.
 We are grateful to an anonymous referee ancl Matthew Date for helphu suggestions that improved the paper., We are grateful to an anonymous referee and Matthew Bate for helpful suggestions that improved the paper.
 This work is based in part on observations made with the Spitzer Space Telescope aud on data products from the Two Micron All Sky Survey., This work is based in part on observations made with the Spitzer Space Telescope and on data products from the Two Micron All Sky Survey.
 We make use of observations made with the NASA/ESA IIubble Space Telescope. obtained from the ESO/ST-ECF Science Archive Facility.," We make use of observations made with the NASA/ESA Hubble Space Telescope, obtained from the ESO/ST-ECF Science Archive Facility."
 This research was supported by an NSERC grant and University of Toronto startup funds (o R.J. Facilities:, This research was supported by an NSERC grant and University of Toronto startup funds to R.J. Facilities:
The times of the six Nav. bursts observed during the RNTE observatious are listed in Table 1. aud the light curves for the three brighest bursts are shown iu Figure |.,"The times of the six X-ray bursts observed during the RXTE observations are listed in Table 1, and the light curves for the three brightest bursts are shown in Figure 4."
 The bursts shown in Figures la aud th were observed during observation |. aud the burst shown i Figure le was observed during observation 2.," The bursts shown in Figures 4a and 4b were observed during observation 4, and the burst shown in Figure 4c was observed during observation 2."
 The properties of the observatio nol bursts are considerably different from the four dimuner biysts., The properties of the observation 4 bursts are considerably different from the four dimmer bursts.
 Their rise times are [ s compared ο about 15s for the cinuuer bursts. aud the observation [| burwis are double-peaked. which may indicate that hey are radius exsion bursts (Lewinetal.1995)).," Their rise times are 4 s compared to about 15 s for the dimmer bursts, and the observation 4 bursts are double-peaked, which may indicate that they are radius expansion bursts \cite{lewin95}) )."
 Siice the peak burst hIuuinositv is less than «0 equid to the Eddineto unuinositv (Lea). a distance upper luit can be derived.," Since the peak burst luminosity is less than or equal to the Eddington luminosity $L_{edd}$ ), a distance upper limit can be derived."
 Using equation. L.10b 0| Lewi1 et al. (, Using equation 4.10b of Lewin et al. (
1992) aud ASSIwing a neutron star mass of 1.14 anda neutji star atiiosphiere with a cosmic annudauce of hydroexορreives Logg=1.89.10 eres,1992) and assuming a neutron star mass of 1.4 $_{\sun}$ and a neutron star atmosphere with a cosmic abundance of hydrogen gives $L_{edd} = 1.89\times 10^{38}$ erg $^{-1}$.
" tI should be noteκα that Log, would be larger for a more massive neutron star or a neturon star atmosphere with a lower hydrogen abundance.", It should be noted that $L_{edd}$ would be larger for a more massive neutron star or a neturon star atmosphere with a lower hydrogen abundance.
 We fine the peak bolometric flux by xoduciug 15 energv spectra near the peak oft1ο burst and fitting the spectra with a model consisting of the persistent emission plus a blacsbodx., We find the peak bolometric flux by producing 1 s energy spectra near the peak of the burst and fitting the spectra with a model consisting of the persistent emission plus a blackbody.
 For the bursts shown in Figures la aud tb. the peak bolometric fhixes are GNs109 re >1 and τιντοὉ erg cni E7. respectively. eiviug a distance upper liuit of 15 kpc.," For the bursts shown in Figures 4a and 4b, the peak bolometric fluxes are $6.8\times 10^{-9}$ erg $^{-2}$ $^{-1}$ and $7.1\times 10^{-9}$ erg $^{-2}$ $^{-1}$, respectively, giving a distance upper limit of 15 kpc."
 The to the high Calactic latitude of NTE J2123058. its distance is of considerable interest.," Due to the high Galactic latitude of XTE J2123–058, its distance is of considerable interest."
" Tere. we derive the cistauce from estimates of the spectral type of he optical companion and the quiesceut naenitude of the source,"," Here, we derive the distance from estimates of the spectral type of the optical companion and the quiescent magnitude of the source."
 Observations mace on August 26/27 eave R=21.5040.06 L998)). and we measured R=21.60+0.06 on September 20.," Observations made on August 26/27 gave $21.50\pm 0.06$ \cite{zc98}) ), and we measured $21.60\pm 0.06$ on September 20."
 There was no evidence for X-ray heating of he optica companion during either observation., There was no evidence for X-ray heating of the optical companion during either observation.
 Since X-ray jieating was not observed and the source Hux renired constant between August 26/27 iud September 20. tlhe source was probably in quiescence.," Since X-ray heating was not observed and the source flux remained constant between August 26/27 and September 20, the source was probably in quiescence."
 Asstuing that the companion is a he star (Pattersou 1981) aid that the cuescent emissjou is donünated WwW endsson from the optical companion gives V = 22.52 aud :ui absolute V“qnaenitude of 8.5. from which we derive a distance of 5.5 kpc.," Assuming that the companion is a K8 star (Patterson 1984) and that the quiescent emission is dominated by emission from the optical companion gives V = 22.52 and an absolute V-magnitude of 8.5, from which we derive a distance of 5.5 kpc."
 In deriviug the quiescent W-hzuid maeguitule from the quiescent R-baucdl naenitude we use Ay = 0.30 and Ay = 0.21 to account for exinction (Schοσο]etal. 1998))., In deriving the quiescent V-band magnitude from the quiescent R-band magnitude we use $_{\rm V}$ = 0.30 and $_{\rm R}$ = 0.24 to account for extinction \cite{schlegel98}) ).
 We note that he derived value for V is consistent with our September 20 V-1uid imieasureiment., We note that the derived value for V is consistent with our September 20 V-band measurement.
 A range of spectral types roni ALO to IND gives a distauce range of 1.5 to 9.1 kpc., A range of spectral types from M0 to K5 gives a distance range of 4.5 to 9.1 kpc.
 We cosider L5 kpe as a distance lower limit since lis method underestimates he distance if light comes from the accretion cisk d1 quiescence aud conclude hat the NTE J2123O58 disance is between 15 Ipc and 15 Ipc., We consider 4.5 kpc as a distance lower limit since this method underestimates the distance if light comes from the accretion disk in quiescence and conclude that the XTE J2123–058 distance is between 4.5 kpc and 15 kpc.
 From the distaice lower lint aud the Galactic latitude. we find that the source must be at least 2.6 spe from the Galacic plane.," From the distance lower limit and the Galactic latitude, we find that the source must be at least 2.6 kpc from the Galactic plane."
 This is unusual since the scale heigif for neuron star LAINB is only Ἰ kpc (wanParadijs&Whie 19953)., This is unusual since the scale height for neutron star LMXB is only 1 kpc \cite{vw95}) ).
 Four optical bursts occurred during our observations., Four optical bursts occurred during our observations.
 O1i June 30. a burst of at least 0.15 V magnitudes occurred caving a 60s exposure.," On June 30, a burst of at least 0.15 V magnitudes occurred during a 60 s exposure."
 We observed two bursts during the iue from Juvy laud July 16. which are shown in Figure 3.," We observed two bursts during the time from July 1 and July 16, which are shown in Figure 3."
 One of these bursts produced increases of 0.38 aid 0.10 πιαστὰudes for two consecutive 30 s exposures and the other produced a 0.21 maenitude iICLCASC ¢luring a GOs expostre., One of these bursts produced increases of 0.38 and 0.10 magnitudes for two consecutive 30 s exposures and the other produced a 0.21 magnitude increase during a 60 s exposure.
 The largest optical burst was observed durius the first exposure taken on Auest 6., The largest optical burst was observed during the first exposure taken on August 6.
 As show iiu Fieure 9. the 120 x burst exposure was 0.17 magnitudes brighter than the following exposures.," As shown in Figure 9, the 120 s burst exposure was 0.47 magnitudes brighter than the following exposures."
 Although we speculate that these bursts are reprocessed X-ray bursts. there were no simultaue«mis ταν observatiois during auv of the optical bursts.," Although we speculate that these bursts are reprocessed X-ray bursts, there were no simultaneous X-ray observations during any of the optical bursts."
 To see if the data are consistent with this hivpotjesijs. we conipared t16 enerev released durimg N-rav bursts to that curing the optical bursts.," To see if the data are consistent with this hypothesis, we compared the energy released during X-ray bursts to that during the optical bursts."
" For a burst where the V inagnitude increases from 16.9 to 16.5 aud lasts 30 s. the burst fluence in the V-MC ij. m5""410UI2 ore 27."," For a burst where the V magnitude increases from 16.9 to 16.5 and lasts 30 s, the burst fluence in the V-band is $8.5\times 10^{-12}$ erg $^{-2}$."
 Td1ο X-ray flueuce for an average NTE J2123058 burst is 2.3«10SN creaeg 2 givingquay an optical to N-rav mrst fluence ratio of ]«10.1., The X-ray fluence for an average XTE J2123–058 burst is $2.3\times 10^{-8}$ erg $^{-2}$ giving an optical to X-ray burst fluence ratio of $4\times 10^{-4}$.
 For the persistent emission. we find Lf Lniios between j3«10 ‘and 1.5«10.," For the persistent emission, we find $_{opt}$ $_{x}$ ratios between $3\times 10^{-4}$ and $1.5\times 10^{-3}$."
" The shuilarity of the persistent aud the burst L;,4/L, ratios iuxicates that the optical bursts are likely to be", The similarity of the persistent and the burst $_{opt}$ $_{x}$ ratios indicates that the optical bursts are likely to be
The fact that the merger predictions do not match the observed major merger history is potentially an explanation for other known problems in the seni-analvtic moclels associated with the Millennium simulation.,The fact that the merger predictions do not match the observed major merger history is potentially an explanation for other known problems in the semi-analytic models associated with the Millennium simulation.
 These include. [or example. the overproduction of red. low mass galaxies in colour-magnituce diagrams.," These include, for example, the overproduction of red, low mass galaxies in colour-magnitude diagrams."
 This problem might be partially alleviated if the galaxy merger history is more substantial than predicted., This problem might be partially alleviated if the galaxy merger history is more substantial than predicted.
 We would like to thank €. De Lucia. 8S. Foucaud. P. Jonsson. EF. Pearce. J. Primack and 8. White for useful cliscussions and the anonymous Referee for many constructive comments that helped to improve the MEE," We would like to thank G. De Lucia, S. Foucaud, P. Jonsson, F. Pearce, J. Primack and S. White for useful discussions and the anonymous Referee for many constructive comments that helped to improve the manuscript."
 SB acknowledges support from. NSE Crant AS'T-0507117., SB acknowledges support from NSF Grant AST-0507117.
""" The Millennium simulation was carried out. by Vireo Consortium. at the Max. Planck Society in Garching.", The Millennium simulation was carried out by the Virgo Consortium at the Max Planck Society in Garching.
 Data on the galaxy population produced by the models used in this work. as well as on the parent halo population. are publicly available at httpz//www.mpa-earching.mpsg.cde/millennium.," Data on the galaxy population produced by the models used in this work, as well as on the parent halo population, are publicly available at http://www.mpa-garching.mpg.de/millennium/."
The finite age of the universe limits the time available for stars to evolve.,The finite age of the universe limits the time available for stars to evolve.
 The lowest mass stars have had insullicient time to evolve to the point where the remnant. is seen as a white chwarf star., The lowest mass stars have had insufficient time to evolve to the point where the remnant is seen as a white dwarf star.
 There is à monotonically increasing relationship between the initial mass of the star and. the resulting white dwarf star and so this implies à minimum mass for white dwarf stars., There is a monotonically increasing relationship between the initial mass of the star and the resulting white dwarf star and so this implies a minimum mass for white dwarf stars.
 The exact value is uncertain. but is around (Dragaglia. Itenzini Bergeron. 1905).," The exact value is uncertain, but is around (Bragaglia, Renzini Bergeron, 1995)."
 Nevertheless. white dwarf stars are found. with masses well below this limit.," Nevertheless, white dwarf stars are found with masses well below this limit."
 These white cwarls are the consequence of binary star evolution in which the evolution of a star through the red giant phase is interrupted by a commonenvelope phase in which a companion star is engulleck by the expanding envelope ancl then rapidlv spirals in towards the core. ejecting the envelope.," These white dwarfs are the consequence of binary star evolution in which the evolution of a star through the red giant phase is interrupted by a common–envelope phase in which a companion star is engulfed by the expanding envelope and then rapidly spirals in towards the core, ejecting the envelope."
 This arrests the formation of the degenerate red. elant core resulting in an anomalously low mass white dwarf star., This arrests the formation of the degenerate red giant core resulting in an anomalously low mass white dwarf star.
 Dramatic evidence for this scenario was provided. by Alarsh. Dhillon Duck (1995).," Dramatic evidence for this scenario was provided by Marsh, Dhillon Duck \shortcite{Marsh95}."
. They observed 7 DA white chvarls selected. for the low mass derived. for them by Bergeron. Saller Liebert (1992) [rom their spectra.," They observed 7 DA white dwarfs selected for the low mass derived for them by Bergeron, Saffer Liebert \shortcite{Berg92} from their spectra."
 Marsh et wwere able to measure racial velocities with accuracics of a few bby using the narrow core of the aabsorption line., Marsh et were able to measure radial velocities with accuracies of a few by using the narrow core of the absorption line.
 Periodic radial velocity variations showed alt least 5 of the 7 stars to be binary stars., Periodic radial velocity variations showed at least 5 of the 7 stars to be binary stars.
 No evidence for binarity was found for | 136 or | 409., No evidence for binarity was found for $+$ 136 or $+$ 409.
 Lt should be emphasized that the observations of Marsh et ceanmot rule out. the possibility that. these two white εναν stars are binaries., It should be emphasized that the observations of Marsh et cannot rule out the possibility that these two white dwarf stars are binaries.
 Although a mainsequence companion more massive than ccan be ruled. out in both cases. the presence of another cool white cdiwarl. very low mass M ονα or a brown ναι. perhaps in a low inclination orbit. cannot be ruled out.," Although a main–sequence companion more massive than can be ruled out in both cases, the presence of another cool white dwarf, very low mass M dwarf or a brown dwarf, perhaps in a low inclination orbit, cannot be ruled out."
 With such strong evidence for the scenario outlined above it would seem to be inevitable that these white cwarls were once members of binary svstems., With such strong evidence for the scenario outlined above it would seem to be inevitable that these white dwarfs were once members of binary systems.
 The nature of any companion. or its fate if it is no longer present. remain open questions.," The nature of any companion, or its fate if it is no longer present, remain open questions."
 Vhe failure of Marsh. et tto. detect binarity in | 136 and | 409 Leck Iben. Tutukoy Yuneleson (LOOT) to suggest these stars are now single stars that are the result of a merger between a white cwarl and the companion responsible for the commonenvelope phase.," The failure of Marsh et to detect binarity in $+$ 136 and $+$ 409 led Iben, Tutukov Yungleson \shortcite{Iben97} to suggest these stars are now single stars that are the result of a merger between a white dwarf and the companion responsible for the common–envelope phase."
 In this paper we show that the detection of à narrow core to the munakes this suggestion verv unlikely unless angular momenttun loss from the merger product is extremely elficient., In this paper we show that the detection of a narrow core to the makes this suggestion very unlikely unless angular momentum loss from the merger product is extremely efficient.
Then. if Wis a solution of (1)) and Wis an operator from this aleebra. so that [ITS|20. then SW is also solution of (1)).,"Then, if $\Psi$ is a solution of \ref{Schrodinger}) ) and $W$ is an operator from this algebra, so that $[W, S] = 0$, then $S \Psi$ is also solution of \ref{Schrodinger}) )."
" For a given classical dispersion £=E(p). Ly=E(0) we define E-polvnonials LE) : ⋅∫∫⊔≺⇁≀⋅∙∩↴⋝⋅↖↽↑↕∐∖∶↴∙⊾↸∖∐↸∖↥⋅⋜↧⊓∐∶↴∙⊾↕∏∐↸⊳↑↕∪∐ It is equivalent to πο that WPGut is a solution of with the initial value Z1,(r0)=e""."," For a given classical dispersion $E = E(p)$, $E_0 \equiv E(0)$ we define $E$ -polynomials $H^{(E)}_n (x,t)$ by the generating function It is equivalent to so that $H^{(E)}_n(x,t)$ is a solution of with the initial value $H_n^{(E)}(x,0) = x^n$."
 From coummtativity [οςN|=0. time evolution of the operator A satisfies and has the form A(t)2e5?μαμιaeZUeU ," From commutativity $[S,K] = 0$, time evolution of the operator $K$ satisfies and has the form $K(t) = e^{-\frac{i}{\hbar}{\cal H} t}K(0)\, e^{\frac{i}{\hbar}{\cal H} t} = e^{-\frac{i}{\hbar}{\cal H} t} x \,e^{\frac{i}{\hbar}{\cal H} t}$."
"Then as follows. operator A generates the infinite hierarchy. of polvnomials according to The non-relativistic dispersion £(p)=p/2ii miplies the IEuuiltouiau operator and the Callilean boost operator From the OoeeneratineC» function we have the Schróddinger polvuonmials If HIPoutEcepitLe"" ds the Kampe de Feniet polvnomial. then HVotoHISTG.Zt or iu terius of the Wermit polvnonial"," Then as follows, operator $K$ generates the infinite hierarchy of polynomials according to The non-relativistic dispersion $E(p) = p^2/2m$ implies the Hamiltonian operator and the Gallilean boost operator From the generating function we have the Schröddinger polynomials If $H^{(KF)}_n(x,t) = exp (t \frac{d^2}{dt^2}) x^n$ is the Kampe de Feriet polynomial, then $H^{(S)}_n (x,t) = H^{KF} (x, \frac{i\hbar}{2m}t)$ or in terms of the Hermit polynomial"
"some of the cross terms from the lensing of galaxies in a nearer tomographic bin by galaxies in a farther tomographic bin, due to the overlap in redshift distributions (solid blue lines).","some of the cross terms from the lensing of galaxies in a nearer tomographic bin by galaxies in a farther tomographic bin, due to the overlap in redshift distributions (solid blue lines)."
 At highest redshifts the magnification terms start to dominate., At highest redshifts the magnification terms start to dominate.
 The cross-correlation between magnification and intrinsic alignment is mostly one of the smallest terms., The cross-correlation between magnification and intrinsic alignment is mostly one of the smallest terms.
 The LA model is currently our best description of galaxy intrinsic alignment over a range of cosmologies., The LA model is currently our best description of galaxy intrinsic alignment over a range of cosmologies.
 It has been shown that ignoring IAs in an analysis can significantly bias the measurement of fundamental cosmological parameters (?).., It has been shown that ignoring IAs in an analysis can significantly bias the measurement of fundamental cosmological parameters \citep{bridleandking}.
 In this section we quantify this bias and examine the effect of moving from the old standard LA implementation to the new one., In this section we quantify this bias and examine the effect of moving from the old standard LA implementation to the new one.
 We also explore the bias produced by using the old implementation rather than the new., We also explore the bias produced by using the old implementation rather than the new.
 A naive approach to cosmic shear measurements would ignore IAs and measure values for cosmological parameters which are systematically biased., A naive approach to cosmic shear measurements would ignore IAs and measure values for cosmological parameters which are systematically biased.
 By employing a flexible IA model and marginalising over a set of nuisance parameters we lose precision but hope to produce unbiased cosmological measurements., By employing a flexible IA model and marginalising over a set of nuisance parameters we lose precision but hope to produce unbiased cosmological measurements.
 In this section we quantify the cosmological bias which results from ignoring IAs and also the cosmological bias which results from employing the HSO4NL LA model rather than the current LA model., In this section we quantify the cosmological bias which results from ignoring IAs and also the cosmological bias which results from employing the HS04NL LA model rather than the current LA model.
 'To this end we employ the cosmological bias formalism of? (seealso ? and appendix A of ?))., To this end we employ the cosmological bias formalism of \citet{huterertbj06} (see also \citet{amarar07} and appendix A of \citet{joachimiea_megazlrg}) ).
" where dpq is the cosmological bias on each parameter considered, Ci; are the projected angular power spectra, acea are the derivatives of these power spectra with respect C)to the cosmological parameters, Cov[Ον(1),Cy(1)] is the covariance matrix of the power spectra, calculated according to eqn."," where $\delta p_{\alpha}$ is the cosmological bias on each parameter considered, $C_{ij}$ are the projected angular power spectra, $\frac{\partial C_{mn}(l)}{\partial p_{\beta}}$ are the derivatives of these power spectra with respect to the cosmological parameters, ${\rm Cov} \left[ C_{ij}(l),C_{mn}(l) \right]$ is the covariance matrix of the power spectra, calculated according to eqn."
" 39 in ?,, and is the inverse of the Fisher Matrix (FM) for our set of FGcosmological parameters, calculated as Each of the FM terms is calculated using the assumed IA model."," 39 in \citet{joachimi_bridle_2009}, and $F_{\alpha\beta}^{-1}$ is the inverse of the Fisher Matrix (FM) for our set of cosmological parameters, calculated as Each of the FM terms is calculated using the assumed IA model."
" AC;; is the difference between the data vector calculated using the assumed model and the true model, Fig."," $\Delta C_{ij}$ is the difference between the data vector calculated using the assumed model and the true model, Fig."
" |2} shows the cosmological bias on the dark energy equation of state parameters if [As are assumed not to exist but they are truly present according to one of: the HSO4NL implementation; the new model; or an intermediate between the two, called HS10NL."," \ref{fig:bias_ignore} shows the cosmological bias on the dark energy equation of state parameters if IAs are assumed not to exist but they are truly present according to one of: the HS04NL implementation; the new model; or an intermediate between the two, called HS10NL."
" The bias is the distance between the centre of the “true” contour [-1,0] and the centre of the contours which assume no IAs."," The bias is the distance between the centre of the “true” contour [-1,0] and the centre of the contours which assume no IAs."
 Note that the bias parameterisation used is based on the FM formalism which assumes purely Gaussian errors in the parameters and observables which respond linearly with respect to the parameters in question., Note that the bias parameterisation used is based on the FM formalism which assumes purely Gaussian errors in the parameters and observables which respond linearly with respect to the parameters in question.
 The parameterisation is very likely to break down for parameter values more than ~2—3c away fromthe fiducial cosmology., The parameterisation is very likely to break down for parameter values more than $\sim 2-3 \sigma$ away fromthe fiducial cosmology.
 As such the more extreme bias values seen in Figs., As such the more extreme bias values seen in Figs.
 and B] should not be read as exact., \ref{fig:bias_ignore} and \ref{fig:bias_oldnew} should not be read as exact.
 What they can tell us is rough relative bias between different scenarios and it is clear that any scenario producing an absolute bias on wo of order unity or above can be considered to be “catastrophically biased”., What they can tell us is rough relative bias between different scenarios and it is clear that any scenario producing an absolute bias on $w_0$ of order unity or above can be considered to be “catastrophically biased”.
 For clarity we use the notation orm when quoting the size of biases on parameters with respect to the errors on the unbiased constraints to remind the reader that they are calculated assuming the Fisher Matrix formalism is valid., For clarity we use the notation $\sigma_{\textrm{FM}}$ when quoting the size of biases on parameters with respect to the errors on the unbiased constraints to remind the reader that they are calculated assuming the Fisher Matrix formalism is valid.
 The intermediate model (HS10NL) is introduced to disentangle the two effects which change between HS04NL and the new LA implementation., The intermediate model (HS10NL) is introduced to disentangle the two effects which change between HS04NL and the new LA implementation.
" It includes the correct factors of a, introduced by ? and shown in equations [38], but it always applies the non-linear matter power spectrum, PIT(k,2), to the IA terms, following the ? interpretation of HSO4NL."," It includes the correct factors of $a$, introduced by \citet{hiratas10_posterratum} and shown in equations \ref{eqn:HS04_P_II_GI_erratum}, but it always applies the non-linear matter power spectrum, $P^{lin}_{\delta\delta}(k,z)$, to the IA terms, following the \citet{bridleandking} interpretation of HS04NL."
" So this HSIONL model applies the correct redshift evolution to IAs but assumes that the effects of nonlinear clustering are always present, even in the II term, by adopting the NLA ansatz."," So this HS10NL model applies the correct redshift evolution to IAs but assumes that the effects of nonlinear clustering are always present, even in the II term, by adopting the NLA ansatz."
" Ignoring IAs causes strong biasing on the dark energy parameters, no matter which IA model is assumed to be true."," Ignoring IAs causes strong biasing on the dark energy parameters, no matter which IA model is assumed to be true."
 The new implementation is the least biased at ~80rM away from the true model., The new implementation is the least biased at $\sim 8\sigma_{\textrm{FM}}$ away from the true model.
" However, ignoring IAs appears to still bias wo by of order ~1.5."," However, ignoring IAs appears to still bias $w_0$ by of order $\sim1.5$."
 The HS10NL model is more biased at ~20σεμ., The HS10NL model is more biased at $\sim20\sigma_{\textrm{FM}}$.
" The HSO4NL implementation produces the strongest cosmological bias, with the contour far outside the plotted area."," The HS04NL implementation produces the strongest cosmological bias, with the contour far outside the plotted area."
 At a point this far away from the fiducial parameter values the Gaussian assumptions of the Fisher matrix and bias formalisms certainly break down and it would be wrong to put much faith in the exact direction/distance of the predicted bias., At a point this far away from the fiducial parameter values the Gaussian assumptions of the Fisher matrix and bias formalisms certainly break down and it would be wrong to put much faith in the exact direction/distance of the predicted bias.
 What is clear however is that the effect is very strong., What is clear however is that the effect is very strong.
" This finding is what we would expect given that the move from HSO4NL to the new implementation has reduced the impact of IAs, not only through changes to their redshift evolution, but also through the removal of IA power on small scales."," This finding is what we would expect given that the move from HS04NL to the new implementation has reduced the impact of IAs, not only through changes to their redshift evolution, but also through the removal of IA power on small scales."
" 'The general biasing trend is the same when we calculate for a stage-III type survey, like the Dark Energy Survey (DES), with biases of ~4orm for the latest model, ~ὅσεν for the HSIONL implementation and ~30orw for the HSO4NL model."," The general biasing trend is the same when we calculate for a stage-III type survey, like the Dark Energy Survey (DES), with biases of $\sim 4 \sigma_{\textrm{FM}}$ for the latest model, $\sim 8 \sigma_{\textrm{FM}}$ for the HS10NL implementation and $\sim 30 \sigma_{\textrm{FM}}$ for the HS04NL model."
" So far we have not taken into account any uncertainty in the IÀ model, but assumed they are zero in the parameter fitting, whichever model we employ to describe them in the true model."," So far we have not taken into account any uncertainty in the IA model, but assumed they are zero in the parameter fitting, whichever model we employ to describe them in the true model."
" In reality, we are aware that our knowledge of IAs from simulations and observations is still developing and relatively uncertain."," In reality, we are aware that our knowledge of IAs from simulations and observations is still developing and relatively uncertain."
" The caseis similar for galaxy bias, introduced in section [2.2}, and employed below."," The caseis similar for galaxy bias, introduced in section \ref{sec:LA_nn}, and employed below."
" To parameterise our ignorance of botheffects and their correlations we use a grid of nuisance parameters where Ax is a constant amplitude parameter, free to vary about a fiducial value of 1, and Qx(k,z) is a grid of NxxNz nodes logarithmically spaced in k/z space, each of which is allowed to vary independently around a fiducial value of 1."," To parameterise our ignorance of botheffects and their cross-correlations we use a grid of nuisance parameters where $A_X$ is a constant amplitude parameter, free to vary about a fiducial value of 1, and $Q_{X}(k,z)$ is a grid of $N_{k} \times N_{z}$ nodes logarithmically spaced in k/z space, each of which is allowed to vary independently around a fiducial value of 1."
 A final smooth grid is created by spline interpolation over the values of the grid nodes., A final smooth grid is created by spline interpolation over the values of the grid nodes.
" For more details of this nuisance parameter grid see ?,, ?,, ?.."," For more details of this nuisance parameter grid see \citet{joachimi_bridle_2009}, , \citet{MGpaper1}, , \citet{MGpaper2}. ."
 This was inspired by the marginalisation of ? which led to the grid implementation in ?.., This was inspired by the marginalisation of \citet{Bernstein_2008} which led to the grid implementation in \citet{joachimi_bridle_2009}. .
"are produced by ions of the carbon isoelectronic sequence, our formalism applies to both.","are produced by ions of the carbon isoelectronic sequence, our formalism applies to both."
" The final expression is independent of temperature, density, and metallicity and depends only on rotational velocity and central mass."," The final expression is independent of temperature, density, and metallicity and depends only on rotational velocity and central mass."
" We assume the gas is photoionized and collisionally excited and begin by writing a simple two-level balance equation where n, and n» are the populations of the lower and upper levels, A», is the spontaneous radiative transition probability (s~!) from level 2 to level 1, and q1? and q», are the collisional excitation and de-excitation coefficients (cm? s!) and are related by q12=9:qiexp( "," We assume the gas is photoionized and collisionally excited and begin by writing a simple two-level balance equation where $n_1$ and $n_2$ are the populations of the lower and upper levels, $A_{21}$ is the spontaneous radiative transition probability $\s^{-1}\,$ ) from level 2 to level 1, and $q_{12}$ and $q_{21}$ are the collisional excitation and de-excitation coefficients $\cm^{3}\s^{-1}\,$ ) and are related by $q_{12} = \frac{\omega_2}{\omega_1} q_{21} \exp(-h\nu/kT)$ ."
The critical density nai;=A21/q21 is the density at which the —hv/KT).collisional de-excitation rate (7542) equals the radiative decay rate., The critical density $n_{\mathrm{crit}} \equiv A_{21}/q_{21}$ is the density at which the collisional de-excitation rate $n_e q_{21}$ ) equals the radiative decay rate.
" For ng«πρµι, the radiative decay dominates and the population of level 2 is For ne>Merit, theusual Boltzmann equation of local thermodynamic equilibrium (LTE) applies."," For $n_e \ll n_{\mathrm{crit}}$, the radiative decay dominates and the population of level 2 is For $n_e \gg n_{\mathrm{crit}}$, theusual Boltzmann equation of local thermodynamic equilibrium (LTE) applies."
" The local emissivity is €—nohv54A»,, and it follows that The critical density marks the transition from the regime where radiative cooling is proportional to ngNoy to the LTE regime where cooling is proportional to rog."," The local emissivity is $\epsilon = n_2 h\nu_{21} A_{21}$, and it follows that The critical density marks the transition from the regime where radiative cooling is proportional to $n_e~n_{\mathrm{ion}}$ to the LTE regime where cooling is proportional to $n_{\mathrm{ion}}$."
 The emitted energy per unit volume per unit time increases monotonically with increasing density., The emitted energy per unit volume per unit time increases monotonically with increasing density.
 The two regimes are illustrated in Figure 1 for the 45007 line., The two regimes are illustrated in Figure \ref{fig:emis} for the $\lambda5007$ line.
 The critical density is represented by the knee near 106 cm~3., The critical density is represented by the knee near $10^{6}$ $^{-3}$.
" Both I10 and Z08 considered the low-density case, and we have no need to reconsider that here."," Both I10 and Z08 considered the low-density case, and we have no need to reconsider that here."
 The remainder of this work will assume Πρ>ng., The remainder of this work will assume $n_e > n_{\mathrm{crit}}$.
" Authors who have discussed the high-density behavior of emission include Nussbaumer Storey (1981), Keenan Aggarwal (1987), Kastner Bhatia (1989), and Osterbrock Ferland (2006)."," Authors who have discussed the high-density behavior of emission include Nussbaumer Storey (1981), Keenan Aggarwal (1987), Kastner Bhatia (1989), and Osterbrock Ferland (2006)."
" From an observational perspective, we note, for example, Andreà,, Dreschel, Starrfield (1994) found n, às high as 108cm~? from lines in classical novae."," From an observational perspective, we note, for example, Andreä,, Dreschel, Starrfield (1994) found $n_e$ as high as $10^{8}$ $^{-3}$ from lines in classical novae."
 We generalize the high-density emissivity to multi-level configurations by using the Ot? calculations of Nussbaumer Storey (1981)., We generalize the high-density emissivity to multi-level configurations by using the $^{+2}$ calculations of Nussbaumer Storey (1981).
" Their calculations are valid for all practical densities and include temperatures as high 40,000 K, safely above the temperature of photoionized O? gas considered in section 3 below."," Their calculations are valid for all practical densities and include temperatures as high $40,000$ K, safely above the temperature of photoionized $^{+2}$ gas considered in section \ref{discmass} below."
 The upper level of the 45007 and 16548 transitions is !02., The upper level of the $\lambda5007$ and $\lambda6548$ transitions is $^1D_2$.
 The fractional population f(!D2) in Table 5 of Nussbaumer Storey does not exceed ο”0.2., The fractional population $f(^1D_2)$ in Table 5 of Nussbaumer Storey does not exceed $\approx0.2$.
" A comparable value applies to N*, and our emissivities take the maximum value where ΠΑ is the density of the ionization stage."," A comparable value applies to $^+$, and our emissivities take the maximum value where $n_A$ is the density of the ionization stage."
 Next we posit a column of gas having local emissivity given by Equation 4.., Next we posit a column of gas having local emissivity given by Equation \ref{final_emis_max}.
 The intensity emitted through the column is where z represents position along the column (or above an annulus perpendicular to the plane of the accretion disc)., The intensity emitted through the column is where $z$ represents position along the column (or above an annulus perpendicular to the plane of the accretion disc).
" To impose an upper limit on the integration variable, we introduce a photon escape probability, Pesc=τς, where 7 is the (Napier) line-center optical depth."," To impose an upper limit on the integration variable, we introduce a photon escape probability, $P_{\mathrm{esc}} = \frac{1}{1+\tau}$, where $\tau$ is the (Napier) line-center optical depth."
 Equation 5 is then written where we have used ΝΑ=ngfdz., Equation \ref{eqn:intensity} is then written where we have used $N_{A} = n_A \int dz$.
" Equation 6 saturates at t7:1, consistent with the thumb that we can see into a cloud only up to optical depth unity (e.g., Rybicki Lightman 1979)."," Equation \ref{eqn:maxintensity} saturates at $\tau \approx 1$, consistent with the rule-of-thumb that we can see into a cloud only up to optical depth unity (e.g., Rybicki Lightman 1979)."
" The upper limit corresponds to N(O+?) =10? cm7?, assuming local line widths are thermal at ez104 K. Optical depth in the lines causes their relative strengths to decrease from the canonical 3:1 ratio to approximate parity in the optically-thick limit."," The upper limit corresponds to $N($ $^{+2}$ ) $\approx 10^{22}$ $^{-2}$, assuming local line widths are thermal at $\approx10^{4}$ K. Optical depth in the lines causes their relative strengths to decrease from the canonical 3:1 ratio to approximate parity in the optically-thick limit."
 This behavior is shown in Figure 2.., This behavior is shown in Figure \ref{fig:doublet}.
 The plotted ratio is visibly greater than unity in both Z08 and 110., The plotted ratio is visibly greater than unity in both Z08 and I10.
" Our limit, therefore, has the consequence of preventing model line-ratios inconsistent with the observations."," Our limit, therefore, has the consequence of preventing model line-ratios inconsistent with the observations."
 Note that Equation 6 does not depend upon density., Note that Equation \ref{eqn:maxintensity} does not depend upon density.
 The green dashed curve in Figure 1 illustrates this and is an important point in our analysis., The green dashed curve in Figure \ref{fig:emis} illustrates this and is an important point in our analysis.
" Above the critical density, the intensity emitted through a fixed column density is independent of the volume density of the gas."," Above the critical density, the intensity emitted through a fixed column density is independent of the volume density of the gas."
" This means that, with no penalty on the total intensity, we can increase the density and, with all else fixed, squeeze the gas into a smaller volume, potentially allowing a stellar-mass black hole explanation."," This means that, with no penalty on the total intensity, we can increase the density and, with all else fixed, squeeze the gas into a smaller volume, potentially allowing a stellar-mass black hole explanation."
" Finally, we obtain a maximum luminosity via an effective surface area."," Finally, we obtain a maximum luminosity via an effective surface area."
" We define the inner radius, ro, of our line-emitting region by assuming the line-emitting gas is in a Keplerian orbit about central mass M with velocity v (which we relate to observed line-widths below) so that We assume the line-emitting region comprises an annulus with width comparable to ro."," We define the inner radius, $r_{0}$ , of our line-emitting region by assuming the line-emitting gas is in a Keplerian orbit about central mass $M$ with velocity $v$ (which we relate to observed line-widths below) so that We assume the line-emitting region comprises an annulus with width comparable to $r_{0}$ ."
 Line luminosities are written, Line luminosities are written
Iarker. D. E.. Woodward. C. E.. Wooden. D. II. The dust grains from 9P/Tempel 1 before aud after the encounter with Deep Science 310. 278-280.,"Harker, D. E., Woodward, C. E., Wooden, D. H. The dust grains from 9P/Tempel 1 before and after the encounter with Deep Science 310, 278-280."
 Jehin. E.. Manfroid. J. IIutsemékkers. D. Cochran. A. L.. Arpigny. C.. Jackson. W. M.. Dauer. IL. Schulz. R.. Zucconi. J.-M. Astrophysical Journal. in press.," Jehin, E., Manfroid, J, Hutsemékkers, D, Cochran, A. L., Arpigny, C., Jackson, W. M., Rauer, H., Schulz, R., Zucconi, J.-M. Astrophysical Journal, in press."
" Jewitt. D.. Meech. Ix. J. Conmeltary grain scattering versus wavelength. or. ""What color is comet The Astrophvsical Journal 310. 937-952."," Jewitt, D., Meech, K. J. Cometary grain scattering versus wavelength, or, “What color is comet The Astrophysical Journal 310, 937-952."
 Ixissler-Patig. M.. Copin. Y.. Ferruit. P.. Péecontal-Rousset. À.. Roth. M. M. “The Euro3D data format: A common FITS data format for integral field Astronomical Notes 325. 159-162.," Kissler-Patig, M., Copin, Y., Ferruit, P., Péccontal-Rousset, A., Roth, M. M. “The Euro3D data format: A common FITS data format for integral field Astronomical Notes 325, 159-162."
 Lantz. D.. Aldering. G.. Antilogus. P.. Donnaud. C.. Capoani. L.. Castera. Αν Copin. Y.. Dubet. D.. Gangler. E.. Hlénnault. F.. Lemonnier. J.-P.. Pain. R.. Péecontal. À..," Lantz, B., Aldering, G., Antilogus, P., Bonnaud, C., Capoani, L., Castera, A., Copin, Y., Dubet, D., Gangler, E., Hénnault, F., Lemonnier, J.-P., Pain, R., Péccontal, A.,"
(2007a).. where the optically thick monochromatic flux in the ze aud o-directions is computed by flux-limited radiative diffusion and where the radiative (ransport of energy in the z-direc(on is solved using a one-rav discrete ordinate method in both optically thin and thick regions.,", where the optically thick monochromatic flux in the $\varpi$ and $\phi$ -directions is computed by flux-limited radiative diffusion and where the radiative transport of energy in the $z$ -direction is solved using a one-ray discrete ordinate method in both optically thin and thick regions."
 Although the central star remains fixed at the grid center. we account [or acceleration of (he reference frame by the planet and by the disk via indirect potentials. as in Michael&Durisen(2010).," Although the central star remains fixed at the grid center, we account for acceleration of the reference frame by the planet and by the disk via indirect potentials, as in \citet{michael2010}."
. The planet integration is done with a Verlet integrator 1995).. aud the indirect potential terms are treated as in Nelsonetal.(2000)..," The planet integration is done with a Verlet integrator \citep[e.g.,][]{hut1995}, and the indirect potential terms are treated as in \citet{nelson2000a}."
 The Rosseland mean aud Planck mean opacities and molecular weights in our simulations are the same as those in Boleyetal.(2006.2007a).. except that we correct the gas mean molecular weight to j( = 2.33.," The Rosseland mean and Planck mean opacities and molecular weights in our simulations are the same as those in \citet{boley2006,boley2007b}, except that we correct the \citet{dalessio2001} gas mean molecular weight to $\mu$ = 2.33."
" The model disk. based on an equlibrium model from Pickettetal.(2003)... orbits a 1 M. star and has à mass M,=0.14... inner and outer radii at 5 and 40 AU. and an initial surface density distribution X—c'?."," The model disk, based on an equlibrium model from \citet{pickett2003}, orbits a 1 $M_{\odot}$ star and has a mass $M_d = 0.14 M_{\odot}$, inner and outer radii at 5 and 40 AU, and an initial surface density distribution $\Sigma \sim \varpi^{-1/2}$."
 The time unit of one ORP (= outer rotation period) is defined as the rotation period of the initial disk at zex 32 AU. or about 150 vr.," The time unit of one ORP (= outer rotation period) is defined as the rotation period of the initial disk at $\varpi \approx$ 32 AU, or about 180 yr."
 The disk is initially located between radial grid zones 30and 240 and is close to isentropic. which results in à Toonmre-Q distribution with a mareinally unstable 2007) minimum value of 1.38 at radial grid zone 161 (26.7 AU).," The disk is initially located between radial grid zones 30and 240 and is close to isentropic, which results in a $Q$ distribution with a marginally unstable \citep[see ][]{durisen2007} minimum value of 1.38 at radial grid zone 161 (26.7 AU)."
 The computational erid extends racially to 512 zones to accommodate expansion of the outer disk once GlIs become nonlinear., The computational grid extends radially to 512 zones to accommodate expansion of the outer disk once GIs become nonlinear.
 An outflow boundary condition is enforced at the upper vertical eric boundary. the outer radial erid boundary. and an inner radial boundary at 2 AU.," An outflow boundary condition is enforced at the upper vertical grid boundary, the outer radial grid boundary, and an inner radial boundary at 2 AU."
 To seed nonaxisvmnietiry. the density distribution is given an initial 0.01 random cell-to-cell perturbation.," To seed nonaxisymmetry, the density distribution is given an initial 0.01 random cell-to-cell perturbation."
 This Letter presents three simulations., This Letter presents three simulations.
 The first. which we call the run. is," The first, which we call the , is"
" lcm lcm For the LLAGNs, the radius of the BLR should be lower than that for broad-line AGNs, if the correlation between Regier and Ly) (Eq. 15))"," 1cm 1cm For the LLAGNs, the radius of the BLR should be lower than that for broad-line AGNs, if the correlation between $R_{\rm BLR}$ and $L_{\rm bol}$ (Eq. \ref{r_blr}) )"
 still holds for low-luminosity sources (butalsoseeWang&Zhang2003).., still holds for low-luminosity sources \citep*[but also see][]{2003MNRAS.340..793W}.
" Our estimate shows that the radiative cooling of the outflow in the source accreting at a rate significantly lower than mi is inefficient, which means that the outflow being expanding adiabatically is a good approximation."," Our estimate shows that the radiative cooling of the outflow in the source accreting at a rate significantly lower than $\dot{m}_{\rm crit}$ is inefficient, which means that the outflow being expanding adiabatically is a good approximation."
" Considering a small volume V in the outflow with gas temperature Τοις and particle number density n, we have for an adiabatic expanding outflow, where ps4,nkT;4,."," Considering a small volume $V$ in the outflow with gas temperature $T_{\rm gas}$ and particle number density $n$, we have for an adiabatic expanding outflow, where $p_{\rm gas}=nkT_{\rm
gas}$."
 The conservation of particles requires Substituting Eq. (27)), The conservation of particles requires Substituting Eq. \ref{conserv}) )
" into (26)), we arrive at le. Tyasοςn?P."," into \ref{energy}) ), we arrive at i.e., $T_{\rm gas}\propto n^{2/3}$."
" As the number density ncr3/? in the outflow (see Eq. 4)),"," As the number density $n\propto
r^{-3/2}$ in the outflow (see Eq. \ref{dens_w2}) ),"
 we find that the gas temperature in an adiabatically expanding outflow., we find that the gas temperature $T_{\rm gas}\propto r^{-1}$ in an adiabatically expanding outflow.
" The broad-line AGNs are relatively luminous, which contain cold accretion disks."," The broad-line AGNs are relatively luminous, which contain cold accretion disks."
 The accretion flows transit to hot ADAFs when the sources are accreting at very low rates., The accretion flows transit to hot ADAFs when the sources are accreting at very low rates.
" Strong outflows may probably be present in LLAGNs, as the ADAFs have a positive Bernoulli parameter (Narayan&Yi 1995a).."," Strong outflows may probably be present in LLAGNs, as the ADAFs have a positive Bernoulli parameter \citep{1995ApJ...444..231N}."
 This implies that the disappearance of BLR in LLAGNs cannot be simply attributed to the lack of outflows from the accretion disk., This implies that the disappearance of BLR in LLAGNs cannot be simply attributed to the lack of outflows from the accretion disk.
" We estimate the cooling of the hot outflows from the ADAF, and find that the radiative cooling of the outflows is always inefficient within the radius of the BLR with any values of"," We estimate the cooling of the hot outflows from the ADAF, and find that the radiative cooling of the outflows is always inefficient within the radius of the BLR with any values of"
iner source eric. and a corresponding increase in the density of rays required. (provided. we can relax the constraint that he source plane must be completely covered by sources with he RBAI).,"finer source grid, and a corresponding increase in the density of rays required (provided we can relax the constraint that the source plane must be completely covered by sources with the RBM)."
 The BM should only be applied with caution to strong ensing scenarios. where the RSAL is far superior.," The RBM should only be applied with caution to strong lensing scenarios, where the RSM is far superior."
 As only one image is followed to the source. there will be an error in the otal magnification when the image is near a critical curve.," As only one image is followed to the source, there will be an error in the total magnification when the image is near a critical curve."
 The RDM is. however. particularly well suited to ooblenms were we want to investigate in detail individual ines-o[-sight for various lens geometries and models.," The RBM is, however, particularly well suited to problems were we want to investigate in detail individual lines-of-sight for various lens geometries and models."
 An important advantage of the RBAL is that we can associate a xwticular image position and shape with the corresponding source position and shape., An important advantage of the RBM is that we can associate a particular image position and shape with the corresponding source position and shape.
 This provides us with the opportunity of following the development of the shape of a rav bundle through a sequence of lens planes. as used in models of cosmological lensing (for example. Wambsganss. Con Ostriker (1998))).," This provides us with the opportunity of following the development of the shape of a ray bundle through a sequence of lens planes, as used in models of cosmological lensing (for example, Wambsganss, Cen Ostriker \shortcite{WAMBSGANSS98}) )."
 An important application of weak lensing is to determine the ellect of small changes in the magnification of standard candle sources (such as Type la Supernovac) on the derived: values of cosmological parameters (Wambseanssetal. 1997)., An important application of weak lensing is to determine the effect of small changes in the magnification of standard candle sources (such as Type Ia Supernovae) on the derived values of cosmological parameters \cite{WAMBSGANSS97}.
. The high accuracy of the RBAI in the weak lensing limit makes it a value tool for such stuclics., The high accuracy of the RBM in the weak lensing limit makes it a value tool for such studies.
 The authors would like to thank Peter Thomas and Andrew Barber (University of Sussex). and Lugh Couchman (University of Western. Ontario) for helpful. discussions.," The authors would like to thank Peter Thomas and Andrew Barber (University of Sussex), and Hugh Couchman (University of Western Ontario) for helpful discussions."
 The authors are erateful to the referee. Joachim Wambseanss. for his insightful comments.," The authors are grateful to the referee, Joachim Wambsganss, for his insightful comments."
 CJL and. DJM ave funcdec by Australian Postgraduate: Awarels., CJF and DJM are funded by Australian Postgraduate Awards.
 CJE is eratelul for financial assistance from the University of Alelbourne’s Melbourne Abroad scholarships scheme. and the Astronomical Society of Australia’s travel grant. scheme.," CJF is grateful for financial assistance from the University of Melbourne's Melbourne Abroad scholarships scheme, and the Astronomical Society of Australia's travel grant scheme."
 The simplest lens model is the point-mass or Schwarzschild lens (for example Schneider et al. (1992).," The simplest lens model is the point-mass or Schwarzschild lens (for example Schneider et al. \shortcite{SCHNEIDER92},"
" Naravan Bartelman (1996))) . for which the deflection angle due to à mass. M. is The climensionless lens equation for a point source with this naoclel is so that there are two images (one located on either side of the lens. and co-linear with the lens) at with corresponding magnifications The total magnification is the sum of the absolute values of the individual image magnifications: 4,=[ye||&|."," Narayan Bartelman \shortcite{NARAYAN96}) ), for which the deflection angle due to a mass, $M$, is The dimensionless lens equation for a point source with this model is so that there are two images (one located on either side of the lens, and co-linear with the lens) at with corresponding magnifications The total magnification is the sum of the absolute values of the individual image magnifications: $\mu_{\rm p} = \vert \mu_+ \vert
+ \vert \mu_- \vert$."
 When the source is far [rom the lens axis (jy> 1). one of the images (savor ) will be significantly demagnilied (ji—<< 1).," When the source is far from the lens axis $y \gg 1$ ), one of the images (say, $x_-$ ) will be significantly demagnified $\mu_- \ll 1$ )."
" The total magnification is then 6,zfi ", The total magnification is then $\mu_{\rm p} \approx \mu_{\rm +}$.
"]tis possible to. derive. an analyic form for the magnification probability. p(y.2). for the Schwarzschild lens [or large values of fy, (Schneideretal.1992).. which is reasonably generic for most lens models (Peacock 1982): On integrating to form the MPLL we ive (1993)... (1992) "," Itis possible to derive an analytic form for the magnification probability, $p(\mu, z)$, for the Schwarzschild lens for large values of $\mu_{\rm p}$ \cite{SCHNEIDER92}, which is reasonably generic for most lens models \cite{PEACOCK82}: : On integrating to form the MPH we have \shortcite{PEI93}, \shortcite{SCHNEIDER92} "
Variability estimates are listed in Table 1.. they can be compared to Table 4 in ?.. where we made the same estimates for the CoRoT exofield database.,"Variability estimates are listed in Table \ref{varfrac}, they can be compared to Table 4 in \cite {COROT-JD}, where we made the same estimates for the CoRoT exofield database."
 A detailed description of the variability selection criteria can be found there., A detailed description of the variability selection criteria can be found there.
" In short. we take a light curve to be variable if at least one of the 3 highest peaks in the amplitude spectrum ας significant (significance parameter Pp«P, ). and has a frequeney value above a certain threshold (f£> finn)."," In short, we take a light curve to be variable if at least one of the 3 highest peaks in the amplitude spectrum is significant (significance parameter $P_{f_i}<P_{max}$ ), and has a frequency value above a certain threshold $f_i>f_{min}$ )."
" We list the resulting percentages for a few combinations of f,;, and P,,,.", We list the resulting percentages for a few combinations of $f_{min}$ and $P_{max}$.
 If we compare these with a short CoRoT observing run. having approximately the same time span as the data. we find a significantly smaller fraction of variables.," If we compare these with a short CoRoT observing run having approximately the same time span as the data, we find a significantly smaller fraction of variables."
" Kepler""s noise levels per measurement are significantly lower. as shown in 2.. but the time sampling 1s less dense: 29.4 min versus 8.5 min or even 32 sec for a significant fraction of the CoRoT data."," 's noise levels per measurement are significantly lower, as shown in \cite{Blomme-Kepler}, but the time sampling is less dense: 29.4 min versus 8.5 min or even 32 sec for a significant fraction of the CoRoT data."
 Probably. the estimates for CoRoT. though conservative. were still influenced by instrumental effects. amongst other things caused by the passage through the South Atlantic Anomaly (2).," Probably, the estimates for CoRoT, though conservative, were still influenced by instrumental effects, amongst other things caused by the passage through the South Atlantic Anomaly \citep{Auvergne-COROT}."
 This passage causes impacts of charged particles on the CCDs. influencing the pixel responses in several ways.," This passage causes impacts of charged particles on the CCDs, influencing the pixel responses in several ways."
 Measured flux levels can temporary increase or decrease. and this translates to discontinuities in the light curves.," Measured flux levels can temporary increase or decrease, and this translates to discontinuities in the light curves."
 We refer to e.g. ? for a more detailed description of these instrumental effects., We refer to e.g. \cite{COROT-jumps} for a more detailed description of these instrumental effects.
 Often. several discontinuities are present in a single CoRoT light curves. causing peaks in various regions of the amplitude spectrum. but always with significant power at frequencies below 0.15 d7!.," Often, several discontinuities are present in a single CoRoT light curves, causing peaks in various regions of the amplitude spectrum, but always with significant power at frequencies below 0.15 $d^{-1}$."
 Figure 3. plots the fraction of objects having significant variability (P-value of the dominant frequency f| below 0.1 ). and with a corresponding amplitude below a certain threshold. as a function of this. threshold value.," Figure \ref{vars-amp} plots the fraction of objects having significant variability (P-value of the dominant frequency $f_1$ below 0.1 ), and with a corresponding amplitude below a certain threshold, as a function of this threshold value."
 It is clear that the majority of variables have very low amplitudes. only reliably detectable using space-based instruments.," It is clear that the majority of variables have very low amplitudes, only reliably detectable using space-based instruments."
 This figure can be compared with Fig., This figure can be compared with Fig.
 6 in ?.. where similar result were obtained (they are included in Fig. 4 2).," 6 in \cite{COROT-JD}, , where similar result were obtained (they are included in Fig. \ref{vars-amp}) )."
 Table 2 summarizes the class statistics. including the remaining numbers of objects using different thresholds for the contamination level (taken from the KIC catalogue) of the light curves.," Table \ref{classes} summarizes the class statistics, including the remaining numbers of objects using different thresholds for the contamination level (taken from the KIC catalogue) of the light curves."
 We determined the numberof good candidates for each, We determined the numberof good candidates for each
"where ó5, indicates cosine phase convention. dillering from our sine convention by s.","where $\phi^*_{31}$ indicates cosine phase convention, differing from our sine convention by $\pi$."
 The form of this formula resembles Eq. (7)), The form of this formula resembles Eq. \ref{e:2}) )
 except for a linear transformation of units., except for a linear transformation of units.
 Our sample contains 54 RRe stars from v Cen observed ον Ixaluzny et al. (, Our sample contains 54 $c$ stars from $\omega$ Cen observed by Kaluzny et al. (
1997. 2003) within the OGLE and CASE oxojects.,"1997, 2003) within the OGLE and CASE projects."
 For this sample we selected only stars free of any complications. such as period changes. multiple periods. or ow-amplitude noisy light curves.," For this sample we selected only stars free of any complications, such as period changes, multiple periods, or low-amplitude noisy light curves."
" For the Rite stars we ollowecl the same procedure as for the ία,", For the RRc stars we followed the same procedure as for the RRab.
 First. we determined their Fourier parameters and then proceeded to it⋅ formulae⋅ akin. to Eq.4 (5)), First we determined their Fourier parameters and then proceeded to fit formulae akin to Eq. \ref{e:1}) )
 by minimizing.ο. 2x7., by minimizing $\chi^2$.
 On the one and. the quality of our photometry for Itc is as good as ever.," On the one hand, the quality of our photometry for RRc is as good as ever."
 On the other hand. the small amplitude and nearly sinusoical light curve of Rite stars produced increased errors or high harmonies. so that above the 6-th harmonic they exceeded: of the value.," On the other hand, the small amplitude and nearly sinusoidal light curve of RRc stars produced increased errors for high harmonics, so that above the 6-th harmonic they exceeded of the value."
 Phe periods. magnitudes. and Fourier parameters of the RRe stars from our sample are isted in Table 3..," The periods, magnitudes and Fourier parameters of the $c$ stars from our sample are listed in Table \ref{t:3}."
" Ehe amplitudes 21). amplitudes ratios 2), and. phase combinations 6), in the function of the period ? or all RRe variables from our sample are shown in Fig."," The amplitudes $A_j$, amplitudes ratios $R_{j1}$ and phase combinations $\phi_{j1}$ in the function of the period $P$ for all $c$ variables from our sample are shown in Fig."
 5., 5.
 These data were fitted with formulae of the type of Eq. (72).," These data were fitted with formulae of the type of Eq. \ref{e:2}) ),"
 using P? and up to 3 Fourier coellicients derived for up o the 4-th harmonic., using $P$ and up to 3 Fourier coefficients derived for up to the 4-th harmonic.
 The results were a bit. disappointing in that all types of formulae vielded rather laree standard deviations and of similar order Dzz0.1 mag., The results were a bit disappointing in that all types of formulae yielded rather large standard deviations and of similar order ${\cal D}\approx 0.1$ mag.
 In Table 4. we ist selection of the best formulae for each length category., In Table \ref{t:4} we list selection of the best formulae for each length category.
 lt must be kept in mind. that according to the £ test. none of hese formulae performed significantly better than the rest. at 0.95 level.," It must be kept in mind, that according to the $F$ test, none of these formulae performed significantly better than the rest, at 0.95 level."
 Note that the phase termi present in formulae 199 and E3 does not vary by more than 0.1 mag., Note that the phase term present in formulae F2 and F3 does not vary by more than 0.1 mag.
 For all these reasons we recommend the use of the simplest E1 formula: The correspondingὃν AM<Vrelations shown in Fig., For all these reasons we recommend the use of the simplest F1 formula: The corresponding $M_V - <V>$ relations shown in Fig.
e 6 reveal that the predicted: Ady vary across fewer than half of the observed range., 6 reveal that the predicted $M_V$ vary across fewer than half of the observed range.
 Lt is rather disturbing to note that all stars except for the 3 most outlving ones form a broad horizontal clump consistent with no correlation of My: with V, It is rather disturbing to note that all stars except for the 3 most outlying ones form a broad horizontal clump consistent with no correlation of $M_V$ with $V$.
olt is hard to imagine a more vivid demonstration that the observed. Luminosity and shape of the RRe light curve do not follow as tight a relation as the theoretical one (see point iv in Sect. ?7)), It is hard to imagine a more vivid demonstration that the observed luminosity and shape of the RRc light curve do not follow as tight a relation as the theoretical one (see point iv in Sect. \ref{s311}) ).
 The question whether the shape of the RRe light curve is or is not correlated to the luminosity in principle does not exelude their use as standard candles., The question whether the shape of the RRc light curve is or is not correlated to the luminosity in principle does not exclude their use as standard candles.
 More problems in this respect stem from the large intrinsic scatter of their magnitudes. of order 0.1. mag.," More problems in this respect stem from the large intrinsic scatter of their magnitudes, of order 0.1 mag."
 For LAIC Zecbru etal. (, For LMC Żeebruń et al. (
2001) list OGLE observations of 450 Rite stars from the fields SC2. SC3. SCA and SCS.,"2001) list OGLE observations of 450 $c$ stars from the fields SC2, SC3, SC4 and SC5."
 As these observations were only a byproduct of a project designed for other purposes. the number of V. filter observations. their exposure leneth and distribution are not optimal for measuring such fine ellects as the third harmonic in the small amplitude. near sinusoidal light curves of RRe.," As these observations were only a byproduct of a project designed for other purposes, the number of $V$ filter observations, their exposure length and distribution are not optimal for measuring such fine effects as the third harmonic in the small amplitude, near sinusoidal light curves of RRc."
 For this practical reason we could not test Simon Clement (1993) formula (Iq. 11)), For this practical reason we could not test Simon Clement (1993) formula (Eq. \ref{e:11}) )
" as none ofthe OGLE light curves vielded Os, with the required accuracy of 0.2 radians.", as none of the OGLE light curves yielded $\phi_{31}$ with the required accuracy of 0.2 radians.
 Application. of our formula (leq. 12)), Application of our formula (Eq. \ref{e:12}) )
 for these data was straightforward. as finding the amplitude of the OGLE V light curves posed. no dillicultv at all.," for these data was straightforward, as finding the amplitude of the OGLE $V$ light curves posed no difficulty at all."
 From the total sample of 450 RRe stars we selected only 57 variables with sullicient amplitude 24)>0.1 mag and errors not exceeding 0.010 mag., From the total sample of 450 $c$ stars we selected only 57 variables with sufficient amplitude $A_1>0.1$ mag and errors not exceeding 0.010 mag.
 In Fig., In Fig.
 7 we plot Ady computed (rom Iq. (12)), 7 we plot $M_V$ computed from Eq. \ref{e:12}) )
 against the observed average magnitude., against the observed average magnitude.
In this respect we were peassured. in a perverse wav. that the LMC RRe stars behave quite similarly to our stars from zc Cen: they both,"In this respect we were reassured, in a perverse way, that the LMC RRc stars behave quite similarly to our stars from $\omega$ Cen: they both"
Spectropolarimetry continues to be a valuable technique for a broad range of astroplivsical ;Doe foo ↴∖↾∏↸⊔↸↴∖≼↸↰⋎↸⋞⊔⊔↴∖∪∐↙∣↑∪∣⋅.: 22005: Bastian 2010). iucludiug applications for circumstellar euvelopes.,"Spectropolarimetry continues to be a valuable technique for a broad range of astrophysical studies (e.g., Adamson 2005; Bastian 2010), including applications for circumstellar envelopes."
 Advances in technology and access to larger telescopes means an ever erowing body of; high quality spectropolariumetric data., Advances in technology and access to larger telescopes means an ever growing body of high quality spectropolarimetric data.
 It is therefore nuportaut that the arsenal of diagnostic methods aud theoretical models iu different astrophysical scenarios keep pace., It is therefore important that the arsenal of diagnostic methods and theoretical models in different astrophysical scenarios keep pace.
 This paper represcuts the fifth in a series devoted toward developing the Hanle effect as tool for casting magnetic fields iu cireuistellar media from resonance line scattering polarization., This paper represents the fifth in a series devoted toward developing the Hanle effect as tool for measuring magnetic fields in circumstellar media from resonance line scattering polarization.
 The observational requirements for the effects examined iu this series are ambitious: high signal-to-noise (S/N) data and high spectral resolving power., The observational requirements for the effects examined in this series are ambitious: high signal-to-noise (S/N) data and high spectral resolving power.
 However. thesedemands are being moet.as illustrated by Harrington Ikulin (20092) in a spectropolarimetyic survey of circumstellar disks at Ho.," However, these demands are being met, as illustrated by Harrington Kuhn (2009a) in a spectropolarimetric survey of circumstellar disks at $\alpha$."
 There are numerous effects that can Παποιος the polarization across resolved lines., There are numerous effects that can influence the polarization across resolved lines.
 A uuuber of researchers have investigated the effects of line opacity for polarization from Thomson scattering (Wood. Brown. Fox 1993: Tarrics 2000: Vins. Harvies. Drew 2005: Wane Wheeler 2008: Tole. Kasen. Nordsicck 2010).," A number of researchers have investigated the effects of line opacity for polarization from Thomson scattering (Wood, Brown, Fox 1993; Harries 2000; Vink, Harries, Drew 2005; Wang Wheeler 2008; Hole, Kasen, Nordsieck 2010)."
 Geatterine bv resonance lines. cau generate polarization.η similar∙∙ to dipole scattering (e... Ieuace 2000a).," Scattering by resonance lines can generate polarization similar to dipole scattering (e.g., Ignace 2000a)."
 With lieh S/N aud high spectral resolution. Warrington Kuhn (20095) have ideitified a new polarizing effect for lines that coincides with Ime absorption.," With high S/N and high spectral resolution, Harrington Kuhn (2009b) have identified a new polarizing effect for lines that coincides with line absorption."
 An explanation for this previously mobserved effect i stars is discussed by Kuhn (2007) and is attributed to t1ο samme dichroic processes detailed bv Trujillo Bueno Laudi DeelTunoceuti (1997) for iuterpretiug polarizations in some solar spectral lines., An explanation for this previously unobserved effect in stars is discussed by Kuhn (2007) and is attributed to the same dichroic processes detailed by Trujillo Bueno Landi Degl'Innocenti (1997) for interpreting polarizations in some solar spectral lines.
 Generally. associated with circular polarization. the Zecman effect. has received acute attention of late as a result of techuiques that co-add many lines (Donati 11997).," Generally associated with circular polarization, the Zeeman effect has received acute attention of late as a result of techniques that co-add many lines (Donati 1997)."
 The method has been used successfully in many studies (soe the review of Donati Landstrect 2009)., The method has been used successfully in many studies (see the review of Donati Landstreet 2009).
 In relation to massive stars. the technique las led to the detection of maguetisin in several stars (e.¢.. Donati 22002. 2006a. 2006b: Neiner 22003: Gruuhut," In relation to massive stars, the technique has led to the detection of magnetism in several stars (e.g., Donati 2002, 2006a, 2006b; Neiner 2003; Grunhut"
Our work sugeestsMOD cli'ections for futre research.,Our work suggests directions for future research.
 For example. a detailed study of the inner regions o ‘the AGB envelope is necessa‘y to examine the eflects of gas heating/cooling. dust anisotropi€ distribution. axd possibly stellar pulsation on the properties of the spiral arm patteru.," For example, a detailed study of the inner regions of the AGB envelope is necessary to examine the effects of gas heating/cooling, dust anisotropic distribution, and possibly stellar pulsation on the properties of the spiral arm pattern."
 In addition. investigations oL binary systeus over a wider range of mass ratios shoulcl be examinect to explore he moclilicatiois in the spiral/arc pattert sin the regime where the effect of the relex motion of the AGB star is also important.," In addition, investigations of binary systems over a wider range of mass ratios should be examined to explore the modifications in the spiral/arc patterns in the regime where the effect of the reflex motion of the AGB star is also important."
 Finally. radiative trausler modeliug of the molecular ine CLUISSIOL ain dust continuuu are encouraged in order to compare theoretical models with observed structures iu detail.," Finally, radiative transfer modeling of the molecular line emission and dust continuum are encouraged in order to compare theoretical models with observed structures in detail."
 Η.]ν is [n]oOateful to Fraucisca ]xeiper aud ]xanak 5E1a for fruitful comanents throug Lreacling tle WALUSCrIp. Paul RRicker for advice on computaticdLal issues. aud oat Soker for cISCUSSIOLL o1 vertical exeusion.," H.K is grateful to Francisca Kemper and Kanak Saha for fruitful comments through reading the manuscript, Paul Ricker for advice on computational issues, and Noam Soker for discussion on vertical extension."
 This research is supported by the Tjeoretical Institute for Advanced Researcl in Astrophysics (TIARA) in the Academia Sinica Institute of Astronny and Astroplysics (ASIAA)., This research is supported by the Theoretical Institute for Advanced Research in Astrophysics (TIARA) in the Academia Sinica Institute of Astronomy and Astrophysics (ASIAA).
 The colhdutatious presented here have been perforlied through tle ASIAA/TIARA computitο resource. Isine FLASH3.0 code developed by the DOE-supported ASC/Alliauce Cener for Astrophysical Tjeriionuclear Flashes at the University of Cicago.," The computations presented here have been performed through the ASIAA/TIARA computing resource, using FLASH3.0 code developed by the DOE-supported ASC/Alliance Center for Astrophysical Thermonuclear Flashes at the University of Chicago."
{ο decrease wilh Gime.,to decrease with time.
 Iowever. the merging or the collapse of low column density structures is likely to trigger star formation ancl outfows of hot eas into the intergalactic medium: this hot gas at temperatures T~LO? 10* IX would be the likely confining medium of low cohunn density Lya clouds and it would contribute mostly to the density of barvonic matter in the Local Universe (Fukugitaetal.1998:Cen&Ostriker1999:Dave2001:Tripp2001).," However, the merging or the collapse of low column density structures is likely to trigger star formation and outflows of hot gas into the intergalactic medium; this hot gas at temperatures $\sim 10^5$ $10^7$ K would be the likely confining medium of low column density $\alpha$ clouds and it would contribute mostly to the density of baryonic matter in the Local Universe \citep{fuk98,cen99,dav01,tri01}."
". In (his paper we have shown that ionizing radiation affects the neutral gas distribution in (he local Universe ancl that any attempt to fit ΝΕ) will a simple power law across (he (transition region (10!*<Nyy10?"" 7) can lead to misleading results.", In this paper we have shown that ionizing radiation affects the neutral gas distribution in the local Universe and that any attempt to fit $f(N_{HI})$ with a simple power law across the transition region $10^{17}<N_{HI}<10^{20}$ $^{-2}$ ) can lead to misleading results.
 A power law distribution for the total gas column density instead fits the data reasonably well. from the Lya forest to the high gas column density in galaxies.," A power law distribution for the total gas column density instead fits the data reasonably well, from the $\alpha$ forest to the high gas column density in galaxies."
 Additional 21-em data. as well as new Lvman limit searches at low redshifts. are needed to implement the knowledge of /(Nj) in (his delicate region and draw a more definite conclusion on the slope of the total gas column density distribution.," Additional 21-cm data, as well as new Lyman limit searches at low redshifts, are needed to implement the knowledge of $f(N_{HI})$ in this delicate region and draw a more definite conclusion on the slope of the total gas column density distribution."
 Furthermore. if 21-cm data are analyzed in detail to constrain the gas and dark matter distribution in outer disks. one can assign a more definite value to the gas compression and picture a unique evolutionary scenario lor the gas distribution aud the UV ionizing background in the Universe.," Furthermore, if 21-cm data are analyzed in detail to constrain the gas and dark matter distribution in outer disks, one can assign a more definite value to the gas compression and picture a unique evolutionary scenario for the gas distribution and the UV ionizing background in the Universe."
 We are grateful to the anonymous referee for a careful reading of the original manuscript., We are grateful to the anonymous referee for a careful reading of the original manuscript.
"where D, is the huninosity distance for an object at redshift :.",where $D_L$ is the luminosity distance for an object at redshift $z$ .
 The AzDj?feLain for this equation las tho units of Jv cung.," The $4 \pi{D_L}^2 f_{24, {\rm obs}}$ for this equation has the units of Jy $^2$."
 The relationship between the observed mouochromatic flux from cach template and the teiiplates monochromatic Iuninositv at a eivenu τοςκατ is approxiniatelv Buear. which allows for a linear fit at cach redshift (.c.. for a redshift exid) to determine a set of coefficients that convert the observed fux at a given redshift to the buninosity or any quantity associated witli the template (c.e.. aad SPR).," The relationship between the observed monochromatic flux from each template and the template's monochromatic luminosity at a given redshift is approximately linear, which allows for a linear fit at each redshift (i.e., for a redshift grid) to determine a set of coefficients that convert the observed flux at a given redshift to the luminosity or any quantity associated with the template (e.g., and SFR)."
" For example. a relatiouship between aud the observed 21 sau flux. foyon.. has the form where A(2) and Bez) are the intercept aud the slope from the linear fit. respectively, Cis a zoro-poiut to reduce covariance in the ft parameters. and £(TIR).9 refers to the associatedM with each» of the ROO unSED templates;"," For example, a relationship between and the observed 24 $\micron$ flux, $f_{24,{\rm obs}}$ , has the form where $A(z)$ and $B(z)$ are the intercept and the slope from the linear fit, respectively, $C$ is a zero-point to reduce covariance in the fit parameters, and $_{z = 0}$ refers to the associated with each of the R09 SED templates."
" Rieke!etal.(2009) use this method to tabulate the mscocfficicuts"" to couvert the monochromatic fluxes i various bauds to the SER of galaxies.", \citet{Rieke09} use this method to tabulate the coefficients to convert the monochromatic fluxes in various bands to the SFR of galaxies.
 This foriialisii successfully estimates aand SER for local galaxies., This formalism successfully estimates and SFR for local galaxies.
 However. a modification is needed to apply the formalisin at hieh + (Rickeetal. 2009).," However, a modification is needed to apply the formalism at high $z$ \citep{Rieke09}."
. If we assmune that IR star-forming galaxies bevond the local Universe are in the main sequence (sve discuss the validity aud applicability of this assuniptiou in Section ??)). it is possible to use the corresponding aas a euide to estimate aand SER at high + without the need for nuueasurenients of individual galaxies.," If we assume that IR star-forming galaxies beyond the local Universe are in the main sequence (we discuss the validity and applicability of this assumption in Section \ref{sec:discuss_sfmodes}) ), it is possible to use the corresponding as a guide to estimate and SFR at high $z$ without the need for measurements of individual galaxies."
 Followine BRujopakunetal.(2011). an appropriate choice of SED to for calelations of the k-correction andl bolometric Tsecorrection hevoud the local Universe is the one corresponding to the same Hocally.," Following \citet{Rujopakarn11}, an appropriate choice of SED to use for calculations of the k-correction and bolometric correction beyond the local Universe is the one corresponding to the same locally."
 This is equivalent to reassiguing the aassociated with cach of the R09 SED templates to new values determined by the ratio of huninosities with equal oon the local U/LIRC: trend to those on the main sequence., This is equivalent to reassigning the associated with each of the R09 SED templates to new values determined by the ratio of luminosities with equal on the local U/LIRG trend to those on the main sequence.
" We will refer to this ratio as the . : ⊓∣↙, ↴⋝↸∖≺∐∖↑↸∖↥⋅↕⊔↕↕∐∖≼↧⋝↖"," We will refer to this ratio as the, $S_i$ , for each SED template."
↽↥⋅↸∖≓∐↑↑↕∐∩⊾↸∖≺∏⋪↧↑↕∪∐↓∙∙↖↖⊽↸∖∐⋪∏⇁↸∖∐∐∐↑↸∖≼↧ ⋅⋅ ΤΗΝ).yall ythe = 1041.20rangeOP of ↴≻ ↕↿∐∐∐∪↴∖↴↕↑↖⇁∪↕↕∪≼↴∙⊾⊓⊺∐↴⋝⊔⊑⋅∖∖↽⊔↴⋝↸∖↸⊳⋜⋯↴∖↴↸∖↑↕∐∖↴∖↴⊓⋅↸∖↑↸⊳↕∐∖≼↧ ⋜⋯∐≼⊓∖⋯⊳∐↸∖↴∖," From Figure \ref{lir_lirsd}, $S_i$ is nearly negligible at $_{z=0} = 10^{11}$ and reaches three orders of magnitude at $_{z=0}$ of $10^{14}$."
"↑↕∐↸∖↸∖∪⊓∐∖↥↴∖∪↕⊔⋜↧∶↴∙↕∐↑⋯∐∖⋜↧↑⊓↽∐↕↴⋝∶⋃ ⋅⋅⋅⋅ . quautitatively,↕↿∐∐∐∪↴∖↴↕⊓↸∖↴∖↴⋜↧↕⋜∐⋅≼↴∙⊾↸∖∆∡↰∣⇁⋜⋯∖↕⋜∐⋅≼↴∙⊾↥⋅↸∖⋜↧↑↸∖↥⋅↑∐⋜↧∐↑∐↸∖ 10we⋅⋅ D. umL... ""nTo . determine"" ↻⋜∐⋜∐⊔↸∖↑↸∖∐∑↸∖↑∐↸∖∐∖↕⋜↧⊓∪∐↴∖∐↻↴⋝∙↖∏↑⊓∐∶↴∙⋜⊓≻⋜∐⋜∏⋯↕⋜↧↑∪⋅ ⋅ the local U/LIRC: relatiou (excludiug the lower linüts of ffrom CO observations)‘ )gand ‘a liuear‘ fit to the main‘ sequence.]qucuce. and then take- the ratio.1 oof oon the main sequence ft to the ROO template a“aa tte anecune‘ ({sec Figure 15s the $; values are tabulated in Table 1))"," To determine $S_i$ quantitatively, we parameterize the relationship by fitting a parabola to the local U/LIRG relation (excluding the lower limits of from CO observations) and a linear fit to the main sequence, and then take the ratio of on the main sequence fit to the R09 template at the same (see Figure \ref{lir_lirsd}; the $S_i$ values are tabulated in Table \ref{table_si}) )."
 Since we effectively increase the Ipuuinositv of cach teiiplate by a factor of 5;. the observed flux will also be increased by the same factor.," Since we effectively increase the luminosity of each template by a factor of $S_i$, the observed flux will also be increased by the same factor."
" Equation ο can thus be rewritten as The new set of. cocfficieuts⋅⋅ 4C). BY=). and C"" can ⋅ a."," Equation \ref{eq3} can thus be rewritten as The new set of coefficients, $A'(z)$, $B'(z)$, and $C'$, can be determined by re-fitting equation \ref{eq4}."
 to only:) encompass. the “stretched nearly negligible ancatTT qni ‘. ⋡∙ - 5; Q4nll ∙⋅ . ∙↴∙ ⋅ ↕↿⋯∐⋯↴∖↴↕↑↖⇁↥⋅⋜⋯∶∙⊾↸∖∪↸⊳⋂∏⋯∖≼↧↴⋝↖⇁↥⋅↸∖⋜↧↕∶↴∙⊾⋜↕↕⋜⋯↕↸∖↴∖↴⇈↕∐∖⊥∩∠∙ ccut off is chosen because we expect this to be approximately at luminosity. lit of star-forming aealaxies)., We have limited the fitting range to only encompass the “stretched” luminosity of $_{{\rm new}}$ $< 14$ because the stretched luminosities at large $S_i$ are far greater than the luminosity range occupied by real galaxies (the $10^{14}$ cut off is chosen because we expect this to be approximately at luminosity limit of star-forming galaxies).
 For a eiven template. theshape as a function ofredshift: remains⋅ the same.," For a given template, theshape as a function ofredshift remains the same."
 The effect of the5; is to change the spacing: between templates., The effect of the$S_i$ is to change the spacing between templates.
 The relationship: between)etween huninositvluminosity and observed fuxflux reaisremains wewell approximated. as lear.. with. residuals. <0.05 dex.," The relationship between luminosity and observed flux remains well approximated as linear, with residuals $<0.05$ dex."
Just when the issue of neutrino mass is clarifving. the cosmological cousequenuces of such mass Las become more puzzling.,"Just when the issue of neutrino mass is clarifying, the cosmological consequences of such mass has become more puzzling."
 About 100/01? of relic neutrinos are everywhere. and the evidence for neutrino mass ix now ecnerally accepted.," About $^3$ of relic neutrinos are everywhere, and the evidence for neutrino mass is now generally accepted."
 It is even likely that there are three neutrino mass differences. the basis for which will be reviewed briefly. but experiments have not vet proved that the mass ds sufficient to have important cosmological effects.," It is even likely that there are three neutrino mass differences, the basis for which will be reviewed briefly, but experiments have not yet proved that the mass is sufficient to have important cosmological effects."
 On the basis not of experiments but of observations. the situation is remarkably coufusing.," On the basis not of experiments but of observations, the situation is remarkably confusing."
" On the one haud. there are may observations showing the matter density of the iudverse (Q,,) to be less than critical (Q,,< 1)."," On the one hand, there are many observations showing the matter density of the universe $\Omega_m$ ) to be less than critical $\Omega_m<1$ )."
" Ou the other haud. the ouly model which fits observations of uulverse structure over more than three orders of magnitude in distance scale is one having cold plus lot (presumably neutrinos) dark matter aud @,,=1."," On the other hand, the only model which fits observations of universe structure over more than three orders of magnitude in distance scale is one having cold plus hot (presumably neutrinos) dark matter and $\Omega_m=1$."
" These observations exclude models of an open universe (low @,,) or oue adding to cold dark matter a cosmological constant CX) so as to male critical density (0,,|O4= 1).", These observations exclude models of an open universe (low $\Omega_m$ ) or one adding to cold dark matter a cosmological constant $\Lambda$ ) so as to make critical density $\Omega_m+\Omega_\Lambda=1$ ).
" Adding hot dark matter helps ouly a little if O,,<1.", Adding hot dark matter helps only a little if $\Omega_m\ll1$.
" Either some of these observations are wrong, or are being wrongly interpreted. or soniething entirely new is needed."," Either some of these observations are wrong, or are being wrongly interpreted, or something entirely new is needed."
 While much more complete aud precise observations will be available soon. the definitive answer regarding the hot dark matter component provided by ucutrinos ust come from terrestrial experiments.," While much more complete and precise observations will be available soon, the definitive answer regarding the hot dark matter component provided by neutrinos must come from terrestrial experiments."
 The ouly experiment providing possible direct evidence for neutrino mass which could be cosinologicallv significant is that of LSND. ancl its results are interpreted here in a wav whichenhances that possibility.," The only experiment providing possible direct evidence for neutrino mass which could be cosmologically significant is that of LSND, and its results are interpreted here in a way whichenhances that possibility."
 The experiucuts on the solar ~., The experiments on the solar $\nu_e$ deficiency and the anomalous $\nu_\mu/\nu_e$ ratio from atmospheric neutrinos are important in this context in establishing the pattern of neutrino masses to learn which neutrinos could contribute to the hot dark matter.
 deficien, The pattern which emerges is made more likely by indirect evidence from the need to rescue heavy-element production by supernovae and possibly by removing a small lack of concordance in the initial abundance of light elements.
cy aud the anomalo," There is now evidence from the first Doppler peak observed in the cosmic microwave background radiation that the total energy density of the universe is the critical value; i.e., $\Omega=1$, and the universe will expand forever at an ever decreasing rate."
us VyfMe ratio fro," Such a flat universe has the only time-stable value of density and is expected in all but very contrived models of an early era of exponential expansion, or “inflation”."
m atin," It has usually been assumed that $\Omega=\Omega_m$; that is, the energy density is the matter density."
osplieric neu," Recent evidence points to $0.3\le\Omega_m\le0.6$, however, based on a variety of observations: high-redshift supernovae type Ia, evolution of galactic clusters, high baryon content of clusters, lensing arcs in clusters, and dynamical estimates from infrared galaxy surveys."
trinos a," On this basis it has become popular to assume that $\Omega_m\approx0.3$, but $\Omega=1$ through the addition of a vacuum energy density, usually designated as a cosmological constant, $\Lambda$."
re inportant in this cou," In stark contrast to this information, either a low-density universe or a critical density universe with a cosmological constant certainly does not fit universe structure as measured over three orders of magnitude in distance scale by the cosmic microwave background and galaxysurveys."
text in establishing t," The only model (CHDM) which fits these extensive data is one having $\Omega_m=1$, of which $\Omega_\nu=0.2$ is in neutrinos, and $\Omega_b=0.1$ in baryons, with the main component being cold dark matter."
he patteru , This is work of Gawiser and \cite{ref:1} who used all published data from the cosmic microwave background and galaxy surveys.
of neutrino," They compared the data with ten models of universe structure, but of concern here are only three of these, CHDM, an open universe model (OCDM) having $\Omega_m=0.5$ , and one $\Lambda$ CDM) having $\Omega_m=0.5$ and $\Omega_\Lambda=0.5$ ."
, In the latter two cases the
NGC 4274. UGC 2082. UGC 4278 and UGC 7774. dy is obtained with method 3 and is always larger than «ως (the last galaxy is the most extreme case. with a difference of 17.6 Mpe between the two estimates).,"NGC 4274, UGC 2082, UGC 4278 and UGC 7774, $d_\mathrm{best}$ is obtained with method 3 and is always larger than $d_\mathrm{T88}$ (the last galaxy is the most extreme case, with a difference of 17.6 Mpc between the two estimates)."
 In general. djs 1s always one of the lowest (when not the lowest) distance estimates available.," In general, $d_\mathrm{T88}$ is always one of the lowest (when not the lowest) distance estimates available."
 The method that we use to calculate star formation rates also bears some discussion., The method that we use to calculate star formation rates also bears some discussion.
 The galaxies that make up the HALOGAS sample are not otherwise included in à uniform observational sample at other wavelengths. in. particular the infrared.," The galaxies that make up the HALOGAS sample are not otherwise included in a uniform observational sample at other wavelengths, in particular the infrared."
 We therefore use different methods for individual targets to calculate SFR. depending on what data are available.," We therefore use different methods for individual targets to calculate SFR, depending on what data are available."
 We prefer to use the combination of Ha and TIR luminosity given by ?.. where Lg is either calculated from IRAS or Spitzer fluxes via the preseriptions given by ?..," We prefer to use the combination of $\alpha$ and TIR luminosity given by \citet{kennicutt_etal_2009}, where $L_{\mathrm{TIR}}$ is either calculated from IRAS or Spitzer fluxes via the prescriptions given by \citet{dale_helou_2002}."
 The IRAS form is used where possible. and the Spitzer form where IRAS fluxes are not available or provide only upper limits for the 25um flux.," The IRAS form is used where possible, and the Spitzer form where IRAS fluxes are not available or provide only upper limits for the $\mu$ m flux."
 IRAS fluxes are obtained from the references listed in Table [.. and the Spitzer fluxes are from the Local Volume Legacy (LVL:?)..," IRAS fluxes are obtained from the references listed in Table \ref{table:sample}, and the Spitzer fluxes are from the Local Volume Legacy \citep[LVL;][]{dale_etal_2009}."
 Πα fluxes are obtained first from ?.. and for targets not included 1n that work. the literature was searched for fluxes.," $\alpha$ fluxes are obtained first from \citet{kennicutt_etal_2008}, and for targets not included in that work, the literature was searched for fluxes."
 In all cases lummosities are calculated using the adopted distances listed in Table 1.., In all cases luminosities are calculated using the adopted distances listed in Table \ref{table:sample}.
" In two cases. UGC 2082 and NGC 4448. no He fluxes are available: but since IRAS 60um and 1005 fluxes are available. we calculate Ly), using the equation given by ?.. and use SER(L(44) as described by ?.."," In two cases, UGC 2082 and NGC 4448, no $\alpha$ fluxes are available; but since IRAS $\mu$ m and $\mu$ m fluxes are available, we calculate $L_{\mathrm{FIR}}$ using the equation given by \citet{rice_etal_1988}, , and use $L_{\mathrm{FIR}}$ ) as described by \citet{kennicutt_1998}."
 The references given in Table | were used to obtain the Ha and IR fluxes., The references given in Table \ref{table:sample} were used to obtain the $\alpha$ and IR fluxes.
 The derivations of the SFR for the HALOGAS sample will become more uniform with the release of the Wide-field Infrared Survey Explorer (WISE:?) data. which will also be well matched in angular resolution to our WSRT observations.," The derivations of the SFR for the HALOGAS sample will become more uniform with the release of the Wide-field Infrared Survey Explorer \citep[WISE;][]{wright_2008} data, which will also be well matched in angular resolution to our WSRT observations."
 In Table 2.. we show environmental information about the target galaxies.," In Table \ref{table:environment}, , we show environmental information about the target galaxies."
 The data are from the group catalog developed from the Nearby Optical Galaxies (NOG) sample (?).., The data are from the group catalog developed from the Nearby Optical Galaxies (NOG) sample \citep{giuricin_etal_2000}.
 The columns in the table are: (1) galaxy ID: (2) number of group members (using the percolation-1 grouping algorithm. see below): (3) NOG group (NOGG) number: (4) total group B luminosity: (5) distance from group center of mass to outermost member: (6.7) minimum and maximum group member velocity relative to group center of mass: (8) contribution of the target to the group B luminosity; (9) fraction of maximum group member B luminosity; (10) normalized distance to group center of mass; (11) target velocity relative to group center of mass.," The columns in the table are: (1) galaxy ID; (2) number of group members (using the percolation-1 grouping algorithm, see below); (3) NOG group (NOGG) number; (4) total group $B$ luminosity; (5) distance from group center of mass to outermost member; (6,7) minimum and maximum group member velocity relative to group center of mass; (8) contribution of the target to the group $B$ luminosity; (9) fraction of maximum group member $B$ luminosity; (10) normalized distance to group center of mass; (11) target velocity relative to group center of mass."
" We have chosen to use group definitions resulting from the 9 ""PI"" algorithm (which ts a percolation friends-of-friendsalgorithm that keeps the velocity grouping condition fixed irrespective of the galaxies” redshift).", We have chosen to use group definitions resulting from the \citet{giuricin_etal_2000} “P1” algorithm (which is a percolation friends-of-friendsalgorithm that keeps the velocity grouping condition fixed irrespective of the galaxies' redshift).
 Because thisdefinition, Because thisdefinition
 , 
Double degenerate (DD) binaries consisting of two white dwarf (WD) stars experience a shrinkage of their orbits caused by gravitational wave radiation.,Double degenerate (DD) binaries consisting of two white dwarf (WD) stars experience a shrinkage of their orbits caused by gravitational wave radiation.
 Sufficiently close binaries with orbital periods of less than half a day will eventually merge within less than a Hubble time., Sufficiently close binaries with orbital periods of less than half a day will eventually merge within less than a Hubble time.
 The outcome of such a merger can be a single compact object like a white dwarf or a neutron star., The outcome of such a merger can be a single compact object like a white dwarf or a neutron star.
 But the merger event may also lead to the ignition of nuclear burning and trigger the formation of rare objects such as extreme helium stars (Saio Jeffery 2002)). R CrB stars (Webbink 1984)) or He sdO stars (Heber 2008:: Napiwotzki 2008).," But the merger event may also lead to the ignition of nuclear burning and trigger the formation of rare objects such as extreme helium stars (Saio Jeffery \cite{saio}) ), R CrB stars (Webbink \cite{webbink}) ) or He sdO stars (Heber \cite{heber3}; Napiwotzki \cite{napiwotzki10}) )."
 The merger of two sufficiently massive C/O white dwarfs may lead to a Supernova of type Ia (SN la. Webbink. 1984)).," The merger of two sufficiently massive C/O white dwarfs may lead to a Supernova of type Ia (SN Ia, Webbink, \cite{webbink}) )."
 SN la play a key role in the study of cosmic evolution., SN Ia play a key role in the study of cosmic evolution.
 They are utilised às standard candles for determining the cosmologicalparameters (e.g. Riess et al. 1998:: , They are utilised as standard candles for determining the cosmologicalparameters (e.g. Riess et al. \cite{riess}; ;
Leibundgut 2001: Perlmutter et al. 1999))., Leibundgut \cite{leibundgut}; Perlmutter et al. \cite{perlmutter}) ).
 There is general consensus that the thermonuclear explosion of a white dwarf. of Chandrasekhar mass causes a SN Ila. The DD merger (Iben Tutukov 1984)) is one of two main scenarios to feed a white dwarf to the Chandrasekhar mass., There is general consensus that the thermonuclear explosion of a white dwarf of Chandrasekhar mass causes a SN Ia. The DD merger (Iben Tutukov \cite{iben}) ) is one of two main scenarios to feed a white dwarf to the Chandrasekhar mass.
 Progenitor candidates for the DD scenario have to merge in less than a Hubble time and the total binary mass must exceed the Chandrasekhar limit., Progenitor candidates for the DD scenario have to merge in less than a Hubble time and the total binary mass must exceed the Chandrasekhar limit.
 In general. DD candidates are assumed to consist. of two white dwarfs.," In general, DD candidates are assumed to consist of two white dwarfs."
 The discovery of the sdB-WD binary 11930-2752. however. has highlighted the importance of such systems as SN [a progenitor candidates (Billérres et al. 2000:;," The discovery of the sdB+WD binary $+$ 2752, however, has highlighted the importance of such systems as SN Ia progenitor candidates (Billèrres et al. \cite{billeres};"
 Maxted et al. 2000:;, Maxted et al. \cite{maxted1};
 Geier et al. 2007))., Geier et al. \cite{geier1}) ).
 Hot subdwarf stars (sdBs) are core helium-burning stars situated at the hot end of the horizontal branch with masses around 0.5M. (Heber 1986)).," Hot subdwarf stars (sdBs) are core helium-burning stars situated at the hot end of the horizontal branch with masses around $0.5\,M_{\rm \odot}$ (Heber \cite{heber1}) )."
 Although the formation of these objects 1s still under debate. it is general consensus that they must loose most of their envelope at the tip of the RGB and evolve directly from the EHB to the WD cooling tracks without ascending the AGB (see Heber 2009 for a review).," Although the formation of these objects is still under debate, it is general consensus that they must loose most of their envelope at the tip of the RGB and evolve directly from the EHB to the WD cooling tracks without ascending the AGB (see Heber \cite{heber} for a review)."
 Since a high fraction of sdBs resides in close binaries with unseen companions (Maxted et., Since a high fraction of sdBs resides in close binaries with unseen companions (Maxted et.
 al 2001:: Naprwotzki et al. 2004)).," al \cite{maxted2}; Napiwotzki et al. \cite{napiwotzki6}) ),"
 such systems can qualify as SN Ia progenitors. if the companion is a white dwarf and the lifetime of the sdB is shorter than the merging time of the binary.," such systems can qualify as SN Ia progenitors, if the companion is a white dwarf and the lifetime of the sdB is shorter than the merging time of the binary."
 This is the case for the sdB+WD binary 11930-2752. which is the best known DD progenitor candidate for SN Ia (Maxted et al. 2000::," This is the case for the sdB+WD binary 1930+2752, which is the best known DD progenitor candidate for SN Ia (Maxted et al. \cite{maxted1};"
 Geier et al. 2007)., Geier et al. \cite{geier1}) ).
 Systematic radial velocity (RV) searches for DDs have been undertaken (e.g. Naprwotzki 2003 and references therein)., Systematic radial velocity (RV) searches for DDs have been undertaken (e.g. Napiwotzki \cite{napiwotzki2} and references therein).
 The largest of these projects was the ESO SN Ia Progenitor Survey (SPY)., The largest of these projects was the ESO SN Ia Progenitor Survey (SPY).
 More than 1000 WDs were checked for RV-variations (Napiwotzki et al. 2001. 2003).," More than $1000$ WDs were checked for RV-variations (Napiwotzki et al. \cite{napiwotzki9,napiwotzki2}) )."
 SPY detected ~100 new DDs (only 18 were known before)., SPY detected $\sim 100$ new DDs (only $18$ were known before).
 One of them may fulfil the criteria for SN Ia progenitor candidates (Naprwotzki et al. 2007)., One of them may fulfil the criteria for SN Ia progenitor candidates (Napiwotzki et al. \cite{napiwotzki3}) ).
 This sample includes about two dozen radial velocity variable sdB stars with invisible companions., This sample includes about two dozen radial velocity variable sdB stars with invisible companions.
 As the sdB stars are intrinsically brighter than the white dwarfs. the nature of the companion 1s not so easily revealed.," As the sdB stars are intrinsically brighter than the white dwarfs, the nature of the companion is not so easily revealed."
 Most sdBs in close binary systems are single-lined and no features of the companions are visible in their spectra., Most sdBs in close binary systems are single-lined and no features of the companions are visible in their spectra.
 In this case the nature of the unseen companion is. hard to constrai since only lower limits can be derived for the companion masses., In this case the nature of the unseen companion is hard to constrain since only lower limits can be derived for the companion masses.
 Given the fact that sdBs are subluminous stars. main sequence companions earlier than K-type can easily be excluded.," Given the fact that sdBs are subluminous stars, main sequence companions earlier than K-type can easily be excluded."
 But if the derived lower limit for the companion mass ts lower than about half a solar mass it is m general not possible to distinguish between a main sequence companion and a compact object like a white dwarf., But if the derived lower limit for the companion mass is lower than about half a solar mass it is in general not possible to distinguish between a main sequence companion and a compact object like a white dwarf.
 Despite the fact that the catalogue of Ritter Kolb (2009)) lists more than 80 sdB binaries. the nature of their companions can only be constrained in special cases (e.g. For et al. 2010:: ," Despite the fact that the catalogue of Ritter Kolb \cite{ritter}) ) lists more than 80 sdB binaries, the nature of their companions can only be constrained in special cases (e.g. For et al. \cite{for}; ;"
Geter etal. 2008)).," Geier etal. \cite{geier2}) ),"
 in particular. from photometric variations due to eclipses. reflection effects or ellipsoidal deformations.," in particular from photometric variations due to eclipses, reflection effects or ellipsoidal deformations."
doublv-ionized helium (Ile HI) at redshifts z<3 (Shull 2004).,doubly-ionized helium (He III) at redshifts $z \leq 3$ (Shull 2004).
 We assume a helium mass fraction bv mass Y=0.244 and define y=(Y/4)/(1—Y)80.0807 (Le fraction by number)., We assume a helium mass fraction by mass $Y = 0.244$ and define $y = (Y/4)/(1-Y) \approx 0.0807$ (He fraction by number).
" The critical density is p,=(1.879x10.7? & 1? where /=(14/100 km ! +).", The critical density is $\rho_{\rm cr} = (1.879 \times 10^{-29}$ g $^{-3}) h^2$ where $h = (H_0$ /100 km $^{-1}$ $^{-1}$ ).
 The above integral can be done analytically: where (1—Y(12y)=39/4)., The above integral can be done analytically: where $(1-Y)(1+y) = (1 - 3Y/4)$.
" For large redshifts. Q,,(1+2)*2Q4. and the integral simplifies to Ilere. we substituted the WALAP-3 parameters (Spergel 2006): 0.0009. Q,,h?=0.12PN and h=0.73£0.03. and adopted a primordial helium abundance Y,=0.244z0.002 (Izotov Thuan 1993: Olive Skillman 2001)."," For large redshifts, $\Omega_m (1+z)^3 \gg \Omega_{\Lambda}$, and the integral simplifies to Here, we substituted the WMAP-3 parameters (Spergel 2006): $\Omega_b h^2 = 0.0223 \pm 0.0009$ , $\Omega_m h^2 = 0.127^{+0.007}_{-0.013}$, and $h = 0.73 \pm 0.03$, and adopted a primordial helium abundance $Y_p = 0.244 \pm 0.002$ (Izotov Thuan 1998; Olive Skillman 2001)."
" A recent paper (Peimbert 2007) suggests a revised Y,=0.2474x0.0028 based on new LHe I atomic data.", A recent paper (Peimbert 2007) suggests a revised $Y_p = 0.2474 \pm 0.0028$ based on new He I atomic data.
" From the approximateκ expression. 33).ey we see that tTx(p,1/27Hy LyA "," From the approximate expression 3), we see that $\tau_e \propto (\rho_{\rm cr} \Omega_b \Omega_m^{-1/2} H_0^{-1}$ )."
":Thus. 7, is nearly independent of the Hubble constant. since p;x5? while the combined parameters. ΟΕ and O,,h7. ave inferred from D/II. CAIB. ancl galaxy clvuamics."," Thus, $\tau_e$ is nearly independent of the Hubble constant, since $\rho_{\rm cr} \propto h^2$ while the combined parameters, $\Omega_b h^2$ and $\Omega_m h^2$, are inferred from D/H, CMB, and galaxy dynamics."
 The scaling with / cancels to lowest order: a slisht dependence remains From the small O4 term in equation (2)., The scaling with $h$ cancels to lowest order; a slight dependence remains from the small $\Omega_{\Lambda}$ term in equation (2).
" If we invert the approximate equation (3). we can estimate (he primary reionization redshift. (1+2,)&(6.6)|.(z,)/0.04|?7. where we scaled to the value. 7,=0.04. expected for [ull ionization back to z,£26."," If we invert the approximate equation (3), we can estimate the primary reionization redshift, $(1+z_r) \approx (6.6) [\tau_e(z_r)/0.04]^{2/3}$, where we scaled to the value, $\tau_e = 0.04$, expected for full ionization back to $z_r \approx 6$."
" As discussed in 2.2. this is approximately the WMAP-3 value ol optical depth (7,2 0.09) reduced by Ar,zz 0.05."," As discussed in 2.2, this is approximately the WMAP-3 value of optical depth $\tau_e \approx 0.09$ ) reduced by $\Delta \tau_e \approx 0.05$ ."
 This extra scattering may arise [rom high--z star formation. X-ray preionization. and residual electrons left after incompleterecombination.," This extra scattering may arise from $z$ star formation, X-ray preionization, and residual electrons left after incompleterecombination."
" The latter electrons are computed to have fractional ionization .r,2(0.5—3.0)xI0.* between 2 = 10700 (Seager 2000).", The latter electrons are computed to have fractional ionization $x_e \approx (0.5-3.0) \times 10^{-3}$ between $z$ = 10–700 (Seager 2000).
 Inaccuracies in computing their contribution therefore add systematic uncertainty to the CMD-derived value of 7., Inaccuracies in computing their contribution therefore add systematic uncertainty to the CMB-derived value of $\tau_e$.
 Partial ionization may also arise from (he first stars (Venkatesan. Tumlinson. Shull 2003. herealter VTS03) aud [rom penetrating X-rays produced by early black holes (VGS01: RicotG Ostriker 2004. 2005).," Partial ionization may also arise from the first stars (Venkatesan, Tumlinson, Shull 2003, hereafter VTS03) and from penetrating X-rays produced by early black holes (VGS01; Ricotti Ostriker 2004, 2005)."
 These additional jionizalion sources contribute electron scattering that must be subtracted from the WALAP-3 values., These additional ionization sources contribute electron scattering that must be subtracted from the WMAP-3 values.
" Thev will lower the reionization redshift. ο,=10.7)23. derived (Spergel 2006) in the absence of partial ionization al z>z,. and they may bring the WAIAP-3 and results into agreement for the epoch of complete reionization."," They will lower the reionization redshift, $z_r = 10.7^{+2.7}_{-2.3}$, derived (Spergel 2006) in the absence of partial ionization at $z > z_r$, and they may bring the WMAP-3 and Gunn-Peterson results into agreement for the epoch of complete reionization."
 In our calculations. described in 3. we make several kev assumptions.," In our calculations, described in 3, we make several key assumptions."
" First. we assume a fully ionized ICM oul (ο z,26. accounting for both and ionized helium. ("," First, we assume a fully ionized IGM out to $z_r \approx 6$, accounting for both $^+$ and ionized helium. ("
Ilelium,Helium
"The density—weighted wy is calculated according to the definition: where ως is the longitude of pericentre computed inside each grid cell, and Γη=3 for the simulations presented here.","The density–weighted $\omega_d$ is calculated according to the definition: where $\omega_c$ is the longitude of pericentre computed inside each grid cell, and $r_{max}=3$ for the simulations presented here."
 Test calculations have shown that wy does not depend strongly on the value of Γη., Test calculations have shown that $\omega_d$ does not depend strongly on the value of $r_{max}$.
" The figure shows that both the binary and disk precess in a prograde sense, and that at late times the apsidal lines of the disk and binary come into almost perfect alignment."," The figure shows that both the binary and disk precess in a prograde sense, and that at late times the apsidal lines of the disk and binary come into almost perfect alignment."
" This alignment coincides with the time when the binary eccentricity reaches a steady value, showing that the secular interaction between disk and binary has a significant effect on the eccentricity evolution of the binary."," This alignment coincides with the time when the binary eccentricity reaches a steady value, showing that the secular interaction between disk and binary has a significant effect on the eccentricity evolution of the binary."
" The disk eccentricity, eg, shown in the lower panel of fig. 3,,"," The disk eccentricity, $e_d$, shown in the lower panel of fig. \ref{w_disk+bin},"
" is defined in the simulations by: where e, is the eccentricity computed at the center of each grid cell.", is defined in the simulations by: where $e_c$ is the eccentricity computed at the center of each grid cell.
 The disk eccentricity evolves similarly to the binary eccentricity and saturates at the relatively small value of eg 0.01., The disk eccentricity evolves similarly to the binary eccentricity and saturates at the relatively small value of $e_d\sim 0.01$ .
" In line with the expectations outlined in section 4.1.1,, the disk and binary become eccentric."," In line with the expectations outlined in section \ref{theory}, the disk and binary become eccentric."
 They achieve a steady state configuration when the apsidal lines of disk and binary are aligned and both disk and binary precess at the same rate in a prograde sense., They achieve a steady state configuration when the apsidal lines of disk and binary are aligned and both disk and binary precess at the same rate in a prograde sense.
 Figures 4and 5 show the results of higher resolution simulations with 256x380 grid cells., Figures \ref{orbit_bin_hr} and \ref{w_disk+bin_hr} show the results of higher resolution simulations with $256 \times 380$ grid cells.
" Figure 4 shows the evolution of the binary semimajor axis and eccentricity, and fig."," Figure \ref{orbit_bin_hr} shows the evolution of the binary semimajor axis and eccentricity, and fig."
" 5 shows the longitudes of pericentre of disk and binary (upper panel), and the disk eccentricity (lower panel)."," \ref{w_disk+bin_hr} shows the longitudes of pericentre of disk and binary (upper panel), and the disk eccentricity (lower panel)."
" Because of their expensive computing time, high-resolution simulations cannot be run for the same length of time as low-resolution ones, and we have been unable to run this simulation until the binary and disk eccentricities have completely saturated."," Because of their expensive computing time, high-resolution simulations cannot be run for the same length of time as low-resolution ones, and we have been unable to run this simulation until the binary and disk eccentricities have completely saturated."
" Nevertheless, we can see by comparing the results from the low and high resolutuion runs, that the two simulations are in excellent agreement."," Nevertheless, we can see by comparing the results from the low and high resolutuion runs, that the two simulations are in excellent agreement."
" At a time of~1.5x10° binary orbits, we see that binary eccentricity is ep~0.07 in both cases, and the disk eccentricity is 0.01 in each case."," At a time of$\simeq 1.5 \times 10^5$ binary orbits, we see that binary eccentricity is $e_b \simeq 0.07$ in both cases, and the disk eccentricity is $\simeq 0.01$ in each case."
 We see from fig., We see from fig.
" 5 that the apsidal lines of the disk and the binary are almost aligned at the end of the simulations, such that it appears that the system is close to an equilibrium state."," \ref{w_disk+bin_hr} that the apsidal lines of the disk and the binary are almost aligned at the end of the simulations, such that it appears that the system is close to an equilibrium state."
 It is therefore reasonable to use the final result of this high-resolution run as the initial condition for simulations dealing with the evolution of protoplanets embedded in circumbinary disks., It is therefore reasonable to use the final result of this high-resolution run as the initial condition for simulations dealing with the evolution of protoplanets embedded in circumbinary disks.
 We present the results of such simulations in the following section., We present the results of such simulations in the following section.
" We have performed three simulations with embedded protoplanets of mass m,= 5, 10 and 20 Ms, respectively."," We have performed three simulations with embedded protoplanets of mass $m_p=5$ , 10 and 20 $M_\oplus$ , respectively."
"To find the operatorT', note that the fiducial model is only used in converting redshifts into distances for the galaxies in our data sample.","To find the operator$T$, note that the fiducial model is only used in converting redshifts into distances for the galaxies in our data sample."
" In the analysis of galaxy clustering, we only need the separation of a galaxy pair, and not the absolute distances to the galaxies."," In the analysis of galaxy clustering, we only need the separation of a galaxy pair, and not the absolute distances to the galaxies."
" For a thin redshift shell, we can convert the separation of one pair of galaxies from the fiducial model to another model by performing the scaling (see, e.g., Seo&Eisenstein (2003))) where 0 is the angle between the radial direction and the direction of the line connecting the pair of galaxies."," For a thin redshift shell, we can convert the separation of one pair of galaxies from the fiducial model to another model by performing the scaling (see, e.g., \cite{SE03}) ) where $\theta$ is the angle between the radial direction and the direction of the line connecting the pair of galaxies."
" Eisensteinetal.(2005) argued that we can use one rescaling parameter, Dv(2). to convert the observed correlation function from the fiducial model to another model as long as the new model is not very different from the fiducial one, and the redshift range of the sample is not large."," \cite{Eisenstein:2005su} argued that we can use one rescaling parameter, $D_V(z)$, to convert the observed correlation function from the fiducial model to another model as long as the new model is not very different from the fiducial one, and the redshift range of the sample is not large."
" Then the separation of one pair of galaxies is converted from the fiducial model to another by In this section, we discuss methods with one and two rescaling parameters, and show that these two methods are equivalent for spherically-averaged data when certain conditions hold (see refsec:proof))."," Then the separation of one pair of galaxies is converted from the fiducial model to another by In this section, we discuss methods with one and two rescaling parameters, and show that these two methods are equivalent for spherically-averaged data when certain conditions hold (see \\ref{sec:proof}) )."
" From ((15)), the observed correlation function with the different model can be written as follows: where ze is the effective redshift of the sample and Dy(2) is defined by Eq.(5))."," From \ref{eq:convert_s_1d}) ), the observed correlation function with the different model can be written as follows: where $z_{eff}$ is the effective redshift of the sample and $D_V(z)$ is defined by \ref{eq:dv}) )."
 The effective redshift we use in this study is z;;;=0.33., The effective redshift we use in this study is $z_{eff}=0.33$.
" Since the results are insensitive to ze; (see refsec:test)), we rescale our result to z;;;=0.35 for comparing with previous works. ((16))"," Since the results are insensitive to $z_{eff}$ (see ), we rescale our result to $z_{eff}=0.35$ for comparing with previous works. \ref{eq:xi_obs}) )"
 can be rewritten as We can apply the same inverse rescalingoperation to the theoretical correlation function: X? can be calculated by substituting ((18)) into ((9))., can be rewritten as We can apply the same inverse rescalingoperation to the theoretical correlation function: $\chi^2$ can be calculated by substituting \ref{eq:inverse_theory}) ) into \ref{eq:chi2_2}) ).
" From ((14), we can convert the spherically-averaged correlation function from some model to the fiducial model by wherethe weighting function w(r,0) is given by where npp(s,0) is the number density of the data pairs."," From \ref{eq:convert_s_2d}) ), we can convert the spherically-averaged correlation function from some model to the fiducial model by wherethe weighting function $w(r,\theta)$ is given by where $n_{DD}(s,\theta)$ is the number density of the data pairs."
" We define inverse operation, T'!, directly since T' is not necessary in our calculation."," We define inverse operation, $T^{-1}$, directly since $T$ is not necessary in our calculation."
 We now apply the inverse operation to the theoretical correlation function: X? can be calculated by substituting ((21)) into ((9))., We now apply the inverse operation to the theoretical correlation function: $\chi^2$ can be calculated by substituting \ref{eq:inverse_theory_2d}) ) into \ref{eq:chi2_2}) ).
 We now show that using one and two rescaling parameters while calculating the spherically-averaged correlation function are equivalent to first order in approximation., We now show that using one and two rescaling parameters while calculating the spherically-averaged correlation function are equivalent to first order in approximation.
" If the size of the survey is much larger than the scales of interest, npp(s,0) would be proportional to ssin0."," If the size of the survey is much larger than the scales of interest, $n_{DD}(s,\theta)$ would be proportional to $s\sin\theta$."
" Hence Next, if the model is close to the fiducial model, we can just consider the first order terms of Dy/D//, H/'4/H, and Da/pii which can be written as following: where |dv|,|5r|,[δα«& 1."," Hence Next, if the model is close to the fiducial model, we can just consider the first order terms of $D_V/D_V^{fid}$, $H^{fid}/H$ , and $D_A/D_A^{fid}$ which can be written as following: where $|\delta_V|, |\delta_r|, |\delta_a| \ll 1$ ."
" From the definition of Dy (see Eq.[5]]). one can obtain a simple relation, 3dy6,+ 26a."," From the definition of $D_V$ (see \ref{eq:dv}] ]), one can obtain a simple relation, $3\delta_V\simeq\delta_r+2\delta_a$ ."
 Let's consider a power law correlation function:, Let's consider a power law correlation function:
factor of —2-3.,factor of $\sim$ 2-3.
 Larec-magnitude. lous-terii changes im the strength of Πα enuüssion and iutriusic polarization are also observed during disk-oss/disk-reunewal eveuts (Wisniewskietal.2010:Draper2011).," Large-magnitude, long-term changes in the strength of $\alpha$ emission and intrinsic polarization are also observed during disk-loss/disk-renewal events \citep{wis10,dra11}."
. Some ITÀeDe stars are also known to exhibit laree (~1 iaagnuitude) IR photometric variability. which sometimes correlates with optical photometric variations and other times is nucorrelated. ou time-scales as short as 1-2 davs (Eiroaetal. 2002).," Some HAeBe stars are also known to exhibit large $\sim$ 1 magnitude) IR photometric variability, which sometimes correlates with optical photometric variations and other times is uncorrelated, on time-scales as short as 1-2 days \citep{eir02}."
. The origin of the photometric variation in these UN Ori objects has been suggested to be variable obscuration by dust chimps (απetal.1991:Waters&Waoelkeus 1998).. although alternate miechauisuis such as variable accretion have also been suggested. (ITerbst&Shevchenko 1999).," The origin of the photometric variation in these UX Ori objects has been suggested to be variable obscuration by dust clumps \citep{gri91, wat98}, although alternate mechanisms such as variable accretion have also been suggested \citep{her99}."
. UX Oxi events also produce siguificaut enhancements in the observed linear polarization. albeit over short time-scales (Cainetal.1991).," UX Ori events also produce significant enhancements in the observed linear polarization, albeit over short time-scales \citep{gri91}."
. Given our new high precision ucar-IR photometry of 27 of 28 ESIICs aud ELIICs aud uew moderate resolution IIo aud ITJ spectroscopic observations of 21 of 28 of these sources. we now discuss how our new data add additional constraints on the evolutionary status of ESITC aud ELIIC stars.," Given our new high precision near-IR photometry of 27 of 28 ESHCs and ELHCs and new moderate resolution $\alpha$ and $\beta$ spectroscopic observations of 21 of 28 of these sources, we now discuss how our new data add additional constraints on the evolutionary status of ESHC and ELHC stars."
 ELIIC 6 (this work) and ELUC 18 (this work aud deWitetal. 20053) have now each been coufiriued to be comprised of two sources: hence. we sugeest that previous efforts to ideutify aud characterize the optical/IR variability and IR excess of these sources should be treated with some caution.," ELHC 6 (this work) and ELHC 18 (this work and \citealt{dew05}) ) have now each been confirmed to be comprised of two sources; hence, we suggest that previous efforts to identify and characterize the optical/IR variability and IR excess of these sources should be treated with some caution."
 We found ELIIC. £ and ELIIC 11 had secondary sources located 172-175 from each star: hence. previous photometric observations which had poorer spatial resolution (or seciug) than these values iia sutter from contamination.," We found ELHC 4 and ELHC 11 had secondary sources located $\farcs$ $\farcs$ 5 from each star; hence, previous photometric observations which had poorer spatial resolution (or seeing) than these values may suffer from contamination."
 We find no support for possible source confusion for either ELIIC 5 or ELIIC δ. which contrasts with the suggestions of previous works (deWitetal.2005:Wisniewski&Bjorkinan 2006).. and therefore we sugeest they continue to be classified as candidate TAcBe targets.," We find no support for possible source confusion for either ELHC 5 or ELHC 8, which contrasts with the suggestions of previous works \citep{dew05,wis06}, and therefore we suggest they continue to be classified as candidate HAeBe targets."
 We detected evidence of photometric variability at the 3-0 level in one IR filter for ELIIC 3 (I[Xs-band) aud ELIIC 5 (J-band. though II aud Is filters also exhibit variability at the 2-7 level).," We detected evidence of photometric variability at the $\sigma$ level in one IR filter for ELHC 3 (Ks-band) and ELHC 5 (J-band, though H and Ks filters also exhibit variability at the $\sigma$ level)."
 ELIIC 21. ELIIC 7. aud ELUC 12 exhibited variability in all 3 filters.," ELHC 21, ELHC 7, and ELHC 12 exhibited variability in all 3 filters."
 The large (1l magnitude) IR photometric brightening observed iu both ELIIC 7 aud ELIIC 12 is at the extreme eud of IR photometric variability observed iu classical Be stars., The large $\sim$ 1 magnitude) IR photometric brightening observed in both ELHC 7 and ELHC 12 is at the extreme end of IR photometric variability observed in classical Be stars.
 ELUC 7 exhibited a very large Ho cussion line in both 1991 aud 2005: hence. it is plausible that the svstei also exhibited siguificaut Πο cussion at the earlier epoch of IR photometric observations (CJauuarv 1998).," ELHC 7 exhibited a very large $\alpha$ emission line in both 1994 and 2005; hence, it is plausible that the system also exhibited significant $\alpha$ emission at the earlier epoch of IR photometric observations (January 1998)."
 Moreover. ELIIC 7 has previously been classified as à UN Oxi variable (deWitetal.2005).," Moreover, ELHC 7 has previously been classified as a UX Ori variable \citep{dew05}."
. Tlence. the IR variability we observe is completely cousisteut with the classification of this star as a IAcBe star.," Hence, the IR variability we observe is completely consistent with the classification of this star as a HAeBe star."
 ELIIC 12 also exhibited a Πα cussion hue with a very similar peak-to-contiuuuna ratio near the epoch of its 1998 IR photometric observation (~3.1) as near the epoch of its 2005 IR photometric observation (3.8)., ELHC 12 also exhibited a $\alpha$ emission line with a very similar peak-to-continuum ratio near the epoch of its 1998 IR photometric observation $\sim$ 3.1) as near the epoch of its 2005 IR photometric observation (3.8).
 Although: we do not have information about the streneth of the EW and FWIIM of Πα for the archival 1998 data. we suggest that there is no clear evidence to indicate that the Πα ciission line was substantially different between the two epochs of obscrvations.," Although we do not have information about the strength of the EW and FWHM of $\alpha$ for the archival 1998 data, we suggest that there is no clear evidence to indicate that the $\alpha$ emission line was substantially different between the two epochs of observations."
 We are rot aware of any classical Be star imnodel which would xediet the maenitude of near-IR flux increase we see for ELIIC 12. aud correspouding lack of strong enhancement in the Πα profile. owing to a disk buiklius event.," We are not aware of any classical Be star model which would predict the magnitude of near-IR flux increase we see for ELHC 12, and corresponding lack of strong enhancement in the $\alpha$ profile, owing to a disk building event."
 We herefore sugeest that the IR photometric variabilitv we observe for ELIIC 12 might indicate that the svstem is a UX Ovi-tvpe ILÀeDoe star., We therefore suggest that the IR photometric variability we observe for ELHC 12 might indicate that the system is a UX Ori-type HAeBe star.
 The IR brightening we observe for ELIIC 21. ~0.5 magnitudes. is within the range of typical variability observed at near-IR waveleneths for both WAcBe stars and classical Be stars.," The IR brightening we observe for ELHC 21, $\sim$ 0.5 magnitudes, is within the range of typical variability observed at near-IR wavelengths for both HAeBe stars and classical Be stars."
 If ELIIC. 21 is a classical Be star. this level of variability could iudicate a factor of ~2-3 chauge in the density of the disk (Dougherty&Tavlor1991).," If ELHC 21 is a classical Be star, this level of variability could indicate a factor of $\sim$ 2-3 change in the density of the disk \citep{dou94}."
. Comparison of the ratio of the peak-to-contiuuun line streneth of the Ta spectra obtained near the time of the 1998 epoch aud 2005 epoch IR photometry reveals little evidence of substantial changes., Comparison of the ratio of the peak-to-continuum line strength of the $\alpha$ spectra obtained near the time of the 1998 epoch and 2005 epoch IR photometry reveals little evidence of substantial changes.
 The region of a Be stars disk responsible for producing Πα emissiou is known to be predominautly located at larger radii than that which produces near-IR cinission (Wisniewskietal.2007b:Car-clofi2011).," The region of a Be star's disk responsible for producing $\alpha$ emission is known to be predominantly located at larger radii than that which produces near-IR emission \citep{wi07b,car11}."
. Thus the ~0.5 maguitude IR brightening we observe in our 2005 data could indicate that a significant. recent (immer) disk building event occurred. which had not vot manifested itself iu diagnostics more scusitive to the outer disk region (IIo spectroscopy).," Thus the $\sim$ 0.5 magnitude IR brightening we observe in our 2005 data could indicate that a significant, recent (inner) disk building event occurred, which had not yet manifested itself in diagnostics more sensitive to the outer disk region $\alpha$ spectroscopy)."
 Alternatively. we cau not rule out that the mocerate-scale variability we have observed could also arise from circumstellar dust clouds which attenuate light from the ceutral source (see c.g. deWitetal. 2005)). and therefore tudicate the svstenaü is a TWAcBe star.," Alternatively, we can not rule out that the moderate-scale variability we have observed could also arise from circumstellar dust clouds which attenuate light from the central source (see e.g. \citealt{dew05}) ), and therefore indicate the system is a HAeBe star."
 Hieher cadence. contemporaneous optical and IR observations of ELIIC 21 should be pursued to differentiate between these two scenarios.," Higher cadence, contemporaneous optical and IR observations of ELHC 21 should be pursued to differentiate between these two scenarios."
 Iu contrast to the observed optical variability of ELIIC stars. which exhibits either bluer when fainter (negative color eradient) or redder when fainter (positive color eradicut) colors. most of the IR variability we observe is erev.," In contrast to the observed optical variability of ELHC stars, which exhibits either bluer when fainter (negative color gradient) or redder when fainter (positive color gradient) colors, most of the IR variability we observe is grey."
 The color effects of uebular scattering. dust reddening. and variable bff eruissiou should all be snall at near-IR wavelengths. so it is perhaps not unexpected that the variability we observe ix erev.," The color effects of nebular scattering, dust reddening, and variable bf-ff emission should all be small at near-IR wavelengths, so it is perhaps not unexpected that the variability we observe is grey."
 We note that the sources which exhibit the ereatest IR variability (ELIIC 7. 12. aud 21) were found to either have positive optical color eradicuts (redder when fainter) or egrev optical color gradients (deWitetal.2005).," We note that the sources which exhibit the greatest IR variability (ELHC 7, 12, and 21) were found to either have positive optical color gradients (redder when fainter) or grey optical color gradients \citep{dew05}."
. Variable dust cloud obscuration has been suggested as the most likely explanation for this optical behavior (deWitetal.S 305)... which is consistent with our previous interpretation for the origin of the IR variability we observe in these sources.," Variable dust cloud obscuration has been suggested as the most likely explanation for this optical behavior \citep{dew05}, which is consistent with our previous interpretation for the origin of the IR variability we observe in these sources."
 Two ELUC stars which exhibit bluer when fainter optical colors were also identified to be variable in one IR filter iu our data (ELIIC 3 aud ELIIC 5): however. the single baudpass IR variability detection preclides us from drawing strone conclusions on the color variability of these Syvstenas.," Two ELHC stars which exhibit bluer when fainter optical colors were also identified to be variable in one IR filter in our data (ELHC 3 and ELHC 5); however, the single bandpass IR variability detection precludes us from drawing strong conclusions on the color variability of these systems."
(2010).. and is interpreted as evidence of energy conservation as a wave expands into a larger area.,", and is interpreted as evidence of energy conservation as a wave expands into a larger area."
 The specific geometry of the wave expansion can be derived by measuring the decay in pulse intensity and growth of pulse width. although previous efforts to measure these parameters accurately have been hampered by the small number of data points associated with each observed event (Warmuthetal.2004b).," The specific geometry of the wave expansion can be derived by measuring the decay in pulse intensity and growth of pulse width, although previous efforts to measure these parameters accurately have been hampered by the small number of data points associated with each observed event \citep{Warmuth:2004ab}."
. In this paper we determine the kinematics of several CBFs and examine the variation in their width and intensity with increasing time (and hence distance)., In this paper we determine the kinematics of several CBFs and examine the variation in their width and intensity with increasing time (and hence distance).
 The data are presented and discussed in Section 2.. with Section 3. detailing the analysis method.," The data are presented and discussed in Section \ref{sect:observations}, with Section \ref{sect:methods} detailing the analysis method."
 The results are discussed in Section 4 and some conclusions drawn in Section 5.. along with some thoughts on future work.," The results are discussed in Section \ref{sect:results} and some conclusions drawn in Section \ref{sect:conclusion}, along with some thoughts on future work."
 The data discussed here were obtained using EUVI. part of the Sun Earth Connection Coronal and Heliospheric Investigation (SECCHI:Howardetal.2008) suite of instruments onboard the STEREO--A and STEREO--B spacecraft.," The data discussed here were obtained using EUVI, part of the Sun Earth Connection Coronal and Heliospheric Investigation \citep[SECCHI;][]{Howard:2008fq} suite of instruments onboard the -A and -B spacecraft."
 EUVI is a normal-incidence telescope of Ritchey-Chréttien design with a pixel scale of1., EUVI is a normal-incidence telescope of Ritchey-Chréttien design with a pixel scale of.
"6"".. It observes the Sun in four passbands (304À.. 171Á.. 195Α.. and 284 Ad). with the peak temperature sensitivities for each passband being approximately 0.07 MK (304 Ay). 1 MK (171 ÀJ. 1.5 MK (195 ΑΗ. and 2.25 MK (284 A))."," It observes the Sun in four passbands (304, 171, 195, and 284 ), with the peak temperature sensitivities for each passband being approximately 0.07 MK (304 ), 1 MK (171 ), 1.5 MK (195 ), and 2.25 MK (284 )."
 The imaging cadence of the EUVI instrument ranges from 1.5 minutes (in the 171 ppassband) to 20 minutes (in the 284 ppassband)., The imaging cadence of the EUVI instrument ranges from 1.5 minutes (in the 171 passband) to 20 minutes (in the 284 passband).
 Although EUVI takes observations in all four of these passbands. only the 171 aand 195 ppassbands were used here.," Although EUVI takes observations in all four of these passbands, only the 171 and 195 passbands were used here."
 This ts due to the high temporal cadence of both passbands (1.5-2.5 minutes for 171 aand 5-10 minutes for 195 Aj) and also as CBFs are more readily observed oof higher contrast) in these two passbands., This is due to the high temporal cadence of both passbands (1.5–2.5 minutes for 171 and 5–10 minutes for 195 ) and also as CBFs are more readily observed of higher contrast) in these two passbands.
 CBFs have been observed in the 284 ppassband (Zhukov&Auchére2004) and the 304 ppassband (Longetal.2008).. but the nature of the data make it difficult to use these passbands for more rigorous analysis.," CBFs have been observed in the 284 passband \citep{Zhukov:2004kh} and the 304 passband \citep{Long:2008eu}, but the nature of the data make it difficult to use these passbands for more rigorous analysis."
 The CBFs were studied using de-rotated base-difference (BD) and percentage base-difference (PBD) images., The CBFs were studied using de-rotated base-difference (BD) and percentage base-difference (PBD) images.
 This involved de-rotating all images for à given event to the same pre-event time. in order to correct for solar rotation between images. and then subtracting a pre-event image to produce a BD image.," This involved de-rotating all images for a given event to the same pre-event time, in order to correct for solar rotation between images, and then subtracting a pre-event image to produce a BD image."
 The ratio of the BD image and the pre-event image was then calculated. giving a PBD image.," The ratio of the BD image and the pre-event image was then calculated, giving a PBD image."
" This is described by the equation. where /, is the Image at any time ¢ and /y Is the pre-event image."," This is described by the equation, where $I_t$ is the image at any time $t$ and $I_0$ is the pre-event image."
 This technique produces images that highlight the CBF and any associated dimming regions. with the intensity values of any given pixel corresponding to the percentage change in intensity with respect to the pre-event image (formoreseeWills-Daveyetal. 2007)..," This technique produces images that highlight the CBF and any associated dimming regions, with the intensity values of any given pixel corresponding to the percentage change in intensity with respect to the pre-event image \citep[for more details see][]{Wills-Davey:2007oa}."
 Figure 1. shows PBD images for the four events studied here from STEREO--B. with the event dates indicated in the lower right of each panel.," Figure \ref{fig:image} shows PBD images for the four events studied here from -B, with the event dates indicated in the lower right of each panel."
 The erupting AR was close to disk centre for both spacecraft for both the 2007 May 19 and 2007 Dec 07 events due to their small separation., The erupting AR was close to disk centre for both spacecraft for both the 2007 May 19 and 2007 Dec 07 events due to their small separation.
 However the eruption was close to disk centre for STEREO--B but on the limb for STEREO--A for both the 2009 Feb 12 and 2009 Feb 13 events., However the eruption was close to disk centre for -B but on the limb for -A for both the 2009 Feb 12 and 2009 Feb 13 events.
 As a result. only the STEREO--B observations are considered here.," As a result, only the -B observations are considered here."
 The observed pulse in each case extended over a large fraction of the solar disk and was observed in multiple images from both passbands., The observed pulse in each case extended over a large fraction of the solar disk and was observed in multiple images from both passbands.
 The observing cadence of the 195 ppassband was 600 s for each event. while the 171 ppassband operated at a cadence of 150 s for each event with the exception of the 2009 Feb 13 event where it had a cadence of 300 s.34567 It is well known that the use of point-and-click techniques in conjunction with PBD images results in large errors that are unique to each user.," The observing cadence of the 195 passband was 600 s for each event, while the 171 passband operated at a cadence of 150 s for each event with the exception of the 2009 Feb 13 event where it had a cadence of 300 s. It is well known that the use of point-and-click techniques in conjunction with PBD images results in large errors that are unique to each user."
 As a result. we developed an algorithm to automatically identify the presence of CBFs in," As a result, we developed an algorithm to automatically identify the presence of CBFs in"
and dashed approach the equilibrium spin rate at an age of a lines)few times 10° yr.,and dashed lines) approach the equilibrium spin rate at an age of a few times $10^5$ yr.
" After this time, the initial condition for the spin rate of these models has effectively been “erased,” in a sense that the subsequent spin rate is insensitive to the initial rate."," After this time, the initial condition for the spin rate of these models has effectively been “erased,” in a sense that the subsequent spin rate is insensitive to the initial rate."
" For the models with low accretion rate (black solid and dashed lines), the torque is not strong enough to drive the spin rate to the equilibrium value within 3 Myr."," For the models with low accretion rate (black solid and dashed lines), the torque is not strong enough to drive the spin rate to the equilibrium value within 3 Myr."
" The case with a slow initial spin rate solid line) evolves near the equilibrium value, but this (blackis only due to a coincidence between the initial and equilibrium spin rates."," The case with a slow initial spin rate (black solid line) evolves near the equilibrium value, but this is only due to a coincidence between the initial and equilibrium spin rates."
" Figure 3 shows results for the W2 case, which is the same as W1, except that the magnetic field is 4 times stronger (B,=2000 G)."," Figure \ref{fig_10b2000} shows results for the W2 case, which is the same as W1, except that the magnetic field is 4 times stronger $B_*=2000$ G)."
" By comparing Figures 2 and 3,, it is clear that the stronger magnetic field results in slower spin rates."," By comparing Figures \ref{fig_10b500} and \ref{fig_10b2000}, it is clear that the stronger magnetic field results in slower spin rates."
" In the age range of 1-3 Myr, all four models in Figure 3 have spin periods in the range of 4-12 days, corresponding to spin rates between and of breakup."," In the age range of 1–3 Myr, all four models in Figure \ref{fig_10b2000} have spin periods in the range of 4–12 days, corresponding to spin rates between and of breakup."
" The models with high accretion rate in the W2 case reffig1052000, , (Fig. redsolidanddashedlines)approachtheequilibriugspittate has different, yr, sooner than the weaker field case (Wl)."," The models with high accretion rate in the W2 case \\ref{fig_10b2000}, red solid and dashed lines) approach the equilibrium spin rate after $\sim 10^5$ yr, sooner than the weaker field case (W1)."
" Furthermore, the models with low accretion rate (black solid and dashed lines) approach the equilibrium spin rate after ~10 yr."," Furthermore, the models with low accretion rate (black solid and dashed lines) approach the equilibrium spin rate after $\sim 10^6$ yr."
" In all models and at all times, the truncation radius is very close to the corotation radius that the thick and thin lines overlap in panel (d)), and (suchthe Alfvénn radius in the winds remains in the range of 15-50 R,."," In all models and at all times, the truncation radius is very close to the corotation radius (such that the thick and thin lines overlap in panel (d)), and the Alfvénn radius in the winds remains in the range of 15–50 $R_*$."
" Figure 4 shows results for the W3 case, which is the same as W2, except that the stellar wind outflow rate is 10 times less (x— 0.01)."," Figure \ref{fig_01b2000} shows results for the W3 case, which is the same as W2, except that the stellar wind outflow rate is 10 times less $\chi=0.01$ )."
" A lower wind mass outflow rate, if all else is equal, means a lower spin-down torque from the stellar wind."," A lower wind mass outflow rate, if all else is equal, means a lower spin-down torque from the stellar wind."
" Thus, the accretion-powered stellar wind in the W3 case is less effective at spinning down the star than the W2 case."," Thus, the accretion-powered stellar wind in the W3 case is less effective at spinning down the star than the W2 case."
" À comparison between Figures 4 and 2 reveals that the W3 case is qualitatively and quantitatively very similar to the W1 case, which has a higher wind outflow rate but a weaker magnetic field."," A comparison between Figures \ref{fig_01b2000} and \ref{fig_10b500} reveals that the W3 case is qualitatively and quantitatively very similar to the W1 case, which has a higher wind outflow rate but a weaker magnetic field."
 The similarity in the spin evolution of these two cases is a consequence of the dependence of the stellar wind torque on the parameters changed., The similarity in the spin evolution of these two cases is a consequence of the dependence of the stellar wind torque on the parameters changed.
" Specifically, equations (8)) and (9)) and m=0.223 results in ΤωοςΜΠΡΟΣ, which means that the factor of 10 difference in M,, is almost completely compensated by the factor of 4 difference in B,."," Specifically, equations \ref{eq_tw}) ) and \ref{eq_ra}) ) and $m=0.223$ results in $T_w \propto \dot M_w^{0.55} B_*^{0.89}$ , which means that the factor of 10 difference in $\dot M_w$ is almost completely compensated by the factor of 4 difference in $B_*$."
" While the stellar wind torques are similar in these two cases, the Alfvénn radius in the wind dependence on these parameters, and thus a “ave-W3 case have a much larger Alfvénn radii than the W1 case (compare the dash-triple-dotted lines in panel (d))."," While the stellar wind torques are similar in these two cases, the Alfvénn radius in the wind has a different dependence on these parameters, and thus the models in the W3 case have a much larger Alfvénn radii than the W1 case (compare the dash-triple-dotted lines in panel (d))."
" It is instructive to compare the cases presented here with the cases in Paper I. Case WI is identical to case O1 in Paper I, except that W1 includes the effect of an accretion-powered stellarwind, via a nonzero value of the mass-loss parameter x."," It is instructive to compare the cases presented here with the cases in Paper I. Case W1 is identical to case O1 in Paper I, except that W1 includes the effect of an accretion-powered stellarwind, via a nonzero value of the mass-loss parameter $\chi$ ."
" Likewise, cases W2 and"," Likewise, cases W2 and"
"We performed optical spectroscopy of LEDÀ 8127E aud TRAS 01250|2832 usiug the Nast Double Spectrograph on the Shaue 3-1 telescope at the Lick Observatory iu 2009 March and August. respectively,","We performed optical spectroscopy of LEDA 84274 and IRAS 01250+2832 using the Kast Double Spectrograph on the Shane 3-m telescope at the Lick Observatory in 2009 March and August, respectively."
 This spectrograph has two separate parallel chanucls - one optimized for the blue aud the other for the red., This spectrograph has two separate parallel channels - one optimized for the blue and the other for the red.
" The 600/1310 erizin and the 600/7500 erating were used in the blue and red spectrometers. respectively,"," The 600/4310 grism and the 600/7500 grating were used in the blue and red spectrometers, respectively."
 The wavelength coverage was 3700 - 8300 with a small gap at around 5500 produced bv the dichroic beamsplitter., The wavelength coverage was 3700 - 8300 with a small gap at around 5500 produced by the dichroic beamsplitter.
 The exposure times for both galaxies were 1500 sec., The exposure times for both galaxies were 1500 sec.
 The spectra of both LEDA 81271 and TRAS 01250|2832 are shown in Figure 2.., The spectra of both LEDA 84274 and IRAS 01250+2832 are shown in Figure \ref{fig:opt_spec}.
 We identify the stroug enuüssion lines in LEDA 81271 and TRAS 01250|2832 as shown in Table 3.., We identify the strong emission lines in LEDA 84274 and IRAS 01250+2832 as shown in Table \ref{tab:opt_leda}.
 Ou the basis of these identifications. we determine the redshitt of z=0.0377 for LEDA 81271L. which is consistent with the redshift previously detezuined by(1995).," On the basis of these identifications, we determine the redshift of z=0.0377 for LEDA 84274, which is consistent with the redshift previously determined by."
. The line ratios of the present results are consistent with previous ones. while the cluission line flux ratios relative to those of the Πα line in blucr part of the spectra are slightly larger than those in(1995).," The line ratios of the present results are consistent with previous ones, while the emission line flux ratios relative to those of the $\alpha$ line in bluer part of the spectra are slightly larger than those in."
. The spectzuii of IRÁS 01250|2832 shows a series of absorption lines iu the blue part of its spectrmu., The spectrum of IRAS 01250+2832 shows a series of absorption lines in the blue part of its spectrum.
 We identity these absorption lines as Call ID|Ix. HI. Cal e-band. H5. aud Fe Tas shown in Figure 2((b).," We identify these absorption lines as CaII H+K, $\delta$, CaI g-band, $\gamma$, and Fe I as shown in Figure \ref{fig:opt_spec}( (b)."
 Although Πα and [N I] emission lines are detected. no other emission lines are detected.," Although $\alpha$ and [N II] emission lines are detected, no other emission lines are detected."
 We use the absorption lines to estimate the redshift of +=0.013., We use the absorption lines to estimate the redshift of $z=0.043$.
 To examine the ealaxy stellar population. we measure the discontinuity.," To examine the galaxy stellar population, we measure the discontinuity."
 This discontinuity is originally defined by, This discontinuity is originally defined by
galaxy evolution.,galaxy evolution.
 The highest-z BAL quasar in their sample. SDSS J!04845.054+463718.3 (hereafter J1048+4637)) at =6.2 shows an extinction curve flat at wavelength A=1700 aand rising at Az1700À.," The $z$ BAL quasar in their sample, SDSS J104845.05+463718.3 (hereafter ) at $z=6.2$ shows an extinction curve flat at wavelength $\lambda \ga 1700$ and rising at $\lambda\la 1700$."
. Maiolinoetal.(2004b) show that the extinction curve of is in excellent agreement with the SN II dust models by TodiniFerrara (2001)., \citet{maiolino04b} show that the extinction curve of is in excellent agreement with the SN II dust models by \citet{todini01}.
.. Bianchi&Schneider(2007) consider. dust destruction by reverse shock in SNe. suggesting that of the initial dust mass survives. and that the extinction curve after the destruction is still consistent with that ofJ1048+4637.," \citet{bianchi07} consider dust destruction by reverse shock in SNe, suggesting that of the initial dust mass survives, and that the extinction curve after the destruction is still consistent with that of."
. More recently. Strattaetal.(2007). show that the dust extinction in the host galaxy of GRB 050904 at +=6.3 can be explained by the extinction curve ofJ10484-4637.. further supporting that the SNe II are the main sources of dust at 2>6.," More recently, \citet{stratta07} show that the dust extinction in the host galaxy of GRB 050904 at $z=6.3$ can be explained by the extinction curve of, further supporting that the SNe II are the main sources of dust at $z>6$."
 Also. Willottetal. find a similar extinction property for CFHQS 71509-1749 (2= 6.12) to that of J1048+4637.," Also, \citet{willott07}
 find a similar extinction property for CFHQS J1509-1749 $z=6.12$ ) to that of ."
. There are other series of theoretical papers on the extinction curves of high-z objects., There are other series of theoretical papers on the extinction curves of $z$ objects.
 Hirashitaetal.(2005.hereafterHOS) calculate the extinction curve based on the dust production calculation by NO3., \citet[][hereafter H05]{hirashita05} calculate the extinction curve based on the dust production calculation by N03.
 They also reproduce the extinction curve of bby using the the dust production in SNe Π. although the dust composition and size distribution are different from those of &Ferrara (2001).," They also reproduce the extinction curve of by using the the dust production in SNe II, although the dust composition and size distribution are different from those of \citet{todini01}."
.. Recently. Nozawaetal.(2007.hereafter have treated the dust destruction by the reverse shock as done by Bianchi&Schneider(2007)... but considering the motion of dust relative to gas caused by the drag force and the destruction of dust in the radiative phase as well as in the non-radiative phase of supernova remnants.," Recently, \citet[][hereafter N07]{nozawa07} have treated the dust destruction by the reverse shock as done by \citet{bianchi07}, but considering the motion of dust relative to gas caused by the drag force and the destruction of dust in the radiative phase as well as in the non-radiative phase of supernova remnants."
 Then. they show that the size distribution of grains supplied in the ISM is strongly modified by the reverse shock.," Then, they show that the size distribution of grains supplied in the ISM is strongly modified by the reverse shock."
 Grains smaller than ~0.02 j/m are efficiently destroyed if the ambient hydrogen number density is larger than 0.1 ?, Grains smaller than $\sim 0.02~\mu$ m are efficiently destroyed if the ambient hydrogen number density is larger than 0.1 $^{-3}$.
" Thus. it is important to reexamine the consistency between the observed extinction curve and the reverse shock destruction,"," Thus, it is important to reexamine the consistency between the observed extinction curve and the reverse shock destruction."
 In this paper. we calculate the extinction curves based on the dust properties calculated by NO7. who have focused on the effect of reverse shock destruction in SNe.," In this paper, we calculate the extinction curves based on the dust properties calculated by N07, who have focused on the effect of reverse shock destruction in SNe."
 Then. we compare the results with observed extinction curves at high z.," Then, we compare the results with observed extinction curves at high $z$."
 This paper is organized as follows., This paper is organized as follows.
 First. we describe our theoretical treatment to calculate the extinction curves of SN II and PISN dust in Section ??..," First, we describe our theoretical treatment to calculate the extinction curves of SN II and PISN dust in Section \ref{sec:model}."
 We show and examine our results in Section 22.., We show and examine our results in Section \ref{sec:results}.
 We discuss our results from the observational viewpoint in Section ??.. and finally give the conclusion of this paper in Section ??..," We discuss our results from the observational viewpoint in Section \ref{sec:obs}, and finally give the conclusion of this paper in Section \ref{sec:sum}."
 We derive the theoretical extinction. curves of dust grains produced in SNe II and PISNe and subsequently destroyed by the reverse shock., We derive the theoretical extinction curves of dust grains produced in SNe II and PISNe and subsequently destroyed by the reverse shock.
 Those grains are considered to be supplied in the interstellar spaces., Those grains are considered to be supplied in the interstellar spaces.
 The grain composition and size distribution in SNe before the destruction is calculated by NO3. whose results are adopted as the initial conditions for the calculations of reverse shock destruction by NO7.," The grain composition and size distribution in SNe before the destruction is calculated by N03, whose results are adopted as the initial conditions for the calculations of reverse shock destruction by N07."
 By using the results by NO7. the extinction curves are calculated by the same method as in HOS.," By using the results by N07, the extinction curves are calculated by the same method as in H05."
 The outline of our calculation is reviewed as follows., The outline of our calculation is reviewed as follows.
 N03 calculate the dust composition and size distribution in the ejecta of PopIII SNe II and PISNe based on the supernova model of Umeda&Nomoto(2002). carefully treating the radial density profile and the temperature evolution.," N03 calculate the dust composition and size distribution in the ejecta of PopIII SNe II and PISNe based on the supernova model of \citet{umeda02}, carefully treating the radial density profile and the temperature evolution."
 As mentioned in HOS. the resulting grain composition and size distribution are not sensitive to the metallicity of progenitor (NO3).," As mentioned in H05, the resulting grain composition and size distribution are not sensitive to the metallicity of progenitor (N03)."
 Thus. the assumption of zero-metallicity is not essential in this paper. and our results can be applicable to metal-enriched systems.," Thus, the assumption of zero-metallicity is not essential in this paper, and our results can be applicable to metal-enriched systems."
 Since it is still uncertain how efficiently the mixing of atoms within SNe occurs. NO3 treat two extreme cases for the mixing of elements: one is theumnmixed case in which the original onion-like structure of elements is preserved. and the other is themixed case in which the elements are uniformly mixed within the helium core.," Since it is still uncertain how efficiently the mixing of atoms within SNe occurs, N03 treat two extreme cases for the mixing of elements: one is the case in which the original onion-like structure of elements is preserved, and the other is the case in which the elements are uniformly mixed within the helium core."
 They show that the formed dust species depend largely on the mixing of seed elements within SNe. because the dominant reactions ehange depending on the ratio of available elements.," They show that the formed dust species depend largely on the mixing of seed elements within SNe, because the dominant reactions change depending on the ratio of available elements."
 The formed grain species in the calculation of 03 are listed in Table |. , The formed grain species in the calculation of N03 are listed in Table \ref{tab:species}. .
In the unmixed ejecta. a variety of grain species (Si. Fe. Ig» $10. MgSIO;. MgO. Al»O;. SiO». FeS. and C) condense. while in the mixed ejecta. only oxide grains (SiO».. MgSIO;. Tg» $10. Al5O;. and FeO.) form.," In the unmixed ejecta, a variety of grain species (Si, Fe, $_2$ $_4$, $_3$, MgO, $_2$ $_3$, $_2$, FeS, and C) condense, while in the mixed ejecta, only oxide grains $_2$, $_3$, $_2$ $_4$, $_2$ $_3$, and $_3$ $_4$ ) form."
 The species are summarized in Table |.., The species are summarized in Table \ref{tab:species}.
 Based on the results in NO3. NO7 treat the dust destruction by the reverse shock in the supernova remnant.," Based on the results in N03, N07 treat the dust destruction by the reverse shock in the supernova remnant."
 They tind that small-sized grains suffer dust destruction by the reverse shock and that the final grain size distribution is biased to larger grains than the original distribution calculated by NO3., They find that small-sized grains suffer dust destruction by the reverse shock and that the final grain size distribution is biased to larger grains than the original distribution calculated by N03.
 We adopt heir results as the properties of grains supplied to the interstellar space., We adopt their results as the properties of grains supplied to the interstellar space.
 Following HOS. we adopt the representative progenitor mass of SNe II as 20 A. and that of PISNe as 170 A7...," Following H05, we adopt the representative progenitor mass of SNe II as 20 $M_\odot$ and that of PISNe as 170 $M_\odot$."
 We also investigate the mixed and unmixed cases., We also investigate the mixed and unmixed cases.
 Therefore. we treat four Causes: (a) mixed SNe IT: (by unmixed SNe II: (c) mixedPISNe: and («d unmixed PISNe.," Therefore, we treat four cases: (a) mixed SNe II; (b) unmixed SNe II; (c) mixedPISNe; and (d) unmixed PISNe."
 All the formulation and the results can be seen in NO3 and N07., All the formulation and the results can be seen in N03 and N07.
 The grains are assumedto be homogeneous and spherical., The grains are assumedto be homogeneous and spherical.
enission in almost all cases.,emission in almost all cases.
 The relative antenna temperature increases wilh increasing cosmic ray ionisation rate., The relative antenna temperature increases with increasing cosmic ray ionisation rate.
" At C=10.! FI. the CO(2-1) transition just starts to dominate over the CO(1-0) line but bv ὁ=10P !, the CO(5-4) line is now the brightest CO line."," At $\zeta=10^{-14}$ $^{-1}$, the CO(2-1) transition just starts to dominate over the CO(1-0) line but by $\zeta=10^{-13}$ $^{-1}$, the CO(5-4) line is now the brightest CO line."
 This line can therefore be used to trace high density svstems at solar metallicity with extremelv high cosmic-ray [hixes such as sources with AGN or a starburst nucleus., This line can therefore be used to trace high density systems at solar metallicity with extremely high cosmic-ray fluxes such as sources with AGN or a starburst nucleus.
 In BOO ib was noted that svstems such as (his show some evidence lor having top-heavy IMES depending on the exact core temperature and magnetic fielcl within these galaxies., In B09 it was noted that systems such as this show some evidence for having top-heavy IMFs depending on the exact core temperature and magnetic field within these galaxies.
 In some active galaxies with phivsical conditions similar to that considered here. (hie magnetic pressure created through ambipolar diffusion max be sullicient to halt core collapse ancl formation of low-mass stars.," In some active galaxies with physical conditions similar to that considered here, the magnetic pressure created through ambipolar diffusion may be sufficient to halt core collapse and formation of low-mass stars."
 The FUV radiation field strength is varied between \=1.7 Draine and X —1.7x10! Draine with the rest of the input parameters fixed at n=10° 7. €=€. and (—10 ' !.," The FUV radiation field strength is varied between $\chi$ =1.7 Draine and $\chi$ $\times10^4$ Draine with the rest of the input parameters fixed at $n$ $^{5}$ $^{-3}$, $\xi=\xi_\odot$ and $\zeta$ $^{-17}$ $^{-1}$."
 The dependence of the CO SED on the FUV radiation field strength. is illustrated in Figure 12..," The dependence of the CO SED on the FUV radiation field strength, is illustrated in Figure \ref{fig:G}."
 lnereasing the FUV radiation field strength results in similar behaviour as increasing (he cosmic ray ionisation rate as in BOO., Increasing the FUV radiation field strength results in similar behaviour as increasing the cosmic ray ionisation rate as in B09.
 Higher-order lines become more aud more dominant and (he antenna temperature increases wilh 4., Higher-order lines become more and more dominant and the antenna temperature increases with $\chi$.
 At the maximum FUV field of €=1.7x10! considered here. it is the CO(3-7) line that is dominant but for à FUV radiation field that is an order of magnitude weaker. the CO(3-2) line dominates.," At the maximum FUV field of $\xi=1.7 \times 10^4$ considered here, it is the CO(8-7) line that is dominant but for a FUV radiation field that is an order of magnitude weaker, the CO(3-2) line dominates."
 The high-order CO transitions can therefore be used as (racers of extremelv high FUV radiation fields associated with the presence of massive stars., The high-order CO transitions can therefore be used as tracers of extremely high FUV radiation fields associated with the presence of massive stars.
 Such svstems may also have top-heavy IMES depending on whether (he magnetic pressure is sufficient to halt the collapse and formation of low-mass stars., Such systems may also have top-heavy IMFs depending on whether the magnetic pressure is sufficient to halt the collapse and formation of low-mass stars.
 Note that due to the large parameter space that we have tried to model. there will inevitable be degeneracies and the CO SEDs produced bx models with 4=1700 and C—10Hs ! may be observationally indistinguishable.," Note that due to the large parameter space that we have tried to model, there will inevitably be degeneracies and the CO SEDs produced by models with $\chi=1700$ and $\zeta=10^{-14}$ $^{-1}$ may be observationally indistinguishable."
 However. we have associated large values for both these parameters wilh active galaxies producing massive stars in this rather simplistic analvsis and it is probably fair to sav that the form of the CO SED observed in our models can be taken to be roughly representative of such active svstenms.," However, we have associated large values for both these parameters with active galaxies producing massive stars in this rather simplistic analysis and it is probably fair to say that the form of the CO SED observed in our models can be taken to be roughly representative of such active systems."
 We now consider (he relative velocity integrated CO antenna temperatures for the hieh, We now consider the relative velocity integrated CO antenna temperatures for the high
caused by observational limitations on intrinsically coherent light curves.,caused by observational limitations on intrinsically coherent light curves.
 We constructed 2000 light curve pairs to simulate simultaneous hard and soft light curves. using the method of ‘Timmer&Ixónig(1995).," We constructed 2000 light curve pairs to simulate simultaneous hard and soft light curves, using the method of \cite{Timmer}."
.. We used a double-bending power law model for the underlving PSD. with the parameters founcl by Mellardy et al. prop.):," We used a double-bending power law model for the underlying PSD, with the parameters found by McHardy et al. ):"
 where 44235105.04520.034;L2.0g=4.5.fui 'dlzand fog—2510/7 Hz.," where $A=3\times 10^{4}, \alpha_L=0, \alpha_M=1.2,
\alpha_H=4.5, f_{bL}=8.7\times 10^{-7} $ Hz and $f_{bH}=2 \times
10^{-3}$ Hz."
 By construction. the light curve pairs have a coherence of unity at all frequencies.," By construction, the light curve pairs have a coherence of unity at all frequencies."
" ‘Time-seale dependent lags were introduce by shifting the phase component of the Fourier. transform. of the ""hard simulatecd ight curves. by 2ἔτι)."," Time-scale dependent lags were introduce by shifting the phase component of the Fourier transform of the `hard' simulated light curves, by $2\pi f \tau(f)$."
 Appropriate Poisson noise was added to the resulting simulated light curves., Appropriate Poisson noise was added to the resulting simulated light curves.
 delata simuations were generated in 100 8 bins and sampled in exactly the same wav as the real 100 s binned light curve., data simulations were generated in 100 s bins and sampled in exactly the same way as the real 100 s binned light curve.
 They were subsequently re-binned in 5400 s evenly spaced bins. just as was done for the real data.," They were subsequently re-binned in 5400 s evenly spaced bins, just as was done for the real data."
 Unlike ddata. the real light curves are practically continuously sampled.," Unlike data, the real light curves are practically continuously sampled."
" ""Therefore. simulated light curves were simply generated in 24 s bins and then re-binned in 96 s bins."," Therefore, simulated light curves were simply generated in 24 s bins and then re-binned in 96 s bins."
 The cross spectrum for cach pair of simulated: light curves was computed. using the same binning used for the real data., The cross spectrum for each pair of simulated light curves was computed using the same binning used for the real data.
 The median of the distributions of coherence ancl ag values of the simulations. for each Fourier frequency. are joted in solid lines in Figs.," The median of the distributions of coherence and lag values of the simulations, for each Fourier frequency, are plotted in solid lines in Figs."
 3.2 and 5.., \ref{coh_test} and \ref{lags_test}.
 The dotted. lines in he same figures mark the spread of the distribution of simulated: values so that of the points lie above the top ine and lie below the bottom line., The dotted lines in the same figures mark the spread of the distribution of simulated values so that of the points lie above the top line and lie below the bottom line.
 The measured coherence in Fig., The measured coherence in Fig.
 3. follows the trend of he median of the distribution of simulations. for both data sets.," \ref{coh_test} follows the trend of the median of the distribution of simulations, for both data sets."
 The small coherence drop on the highest frequeney bins is probably partly due to inaccurate Poisson noise corrections. as suggested. by. the simulations (11ο ormula for this correction. derived. by Vaughan&Nowak(1997) is not strictly applicable when the variability signal-o-noise is low).," The small coherence drop on the highest frequency bins is probably partly due to inaccurate Poisson noise corrections, as suggested by the simulations (the formula for this correction derived by \citet{Vaughan_coh} is not strictly applicable when the variability signal-to-noise is low)."
 Above this frequency. the coherence drops slightly below the distribution of simulated cata. indicating a real but smal (10) loss of coicrence.," Above this frequency, the coherence drops slightly below the distribution of simulated data, indicating a real but small $<$ ) loss of coherence."
 Finally. our results suggest 1mt the strong drop a»»ve 10 Lz can be casily explained »v Poisson noise ellects. and hence is most probably not intrinsic.," Finally, our results suggest that the strong drop above $10^{-3}$ Hz can be easily explained by Poisson noise effects, and hence is most probably not intrinsic."
" As for the phase shi (in the Fourier components of the two light curves. initially we usec τς)Q.04f'"""". Le. the best fit to tje lag s»ectrum. over the entire frequency range. as the uncerlving lag spectrum."," As for the phase shift in the Fourier components of the two light curves, initially we used $\tau(f)=0.04f^{-0.9}$, i.e. the best fit to the lag spectrum over the entire frequency range, as the underlying lag spectrum."
 As seen in Fig., As seen in Fig.
 4. this assumed lag spectrum falls well below the best-celined lag measurements in the middle of the frequeney range probed., \ref{lags_tot} this assumed lag spectrum falls well below the best-defined lag measurements in the middle of the frequency range probed.
 ot surprisinglv. the four central data points remain above he top of the disribution of simulated. lags. while wo high-frequency. poins still fall below it. implying that he underline lag spectrum is inconsistent with the simple »ower law fitted to the data.," Not surprisingly, the four central data points remain above the top of the distribution of simulated lags, while two high-frequency points still fall below it, implying that the underlying lag spectrum is inconsistent with the simple power law fitted to the data."
" We repeated the test. using τι=0.5f""7 as the uncerlving lag spectrum. i.e. the fit othe ~107.57 Lz frequeney range."," We repeated the test using $\tau(f)=0.5f^{-0.7}$ as the underlying lag spectrum, i.e. the fit to the $\sim
10^{-5}-5\times 10^{-4}$ Hz frequency range."
 The resulting ag spectra distribution is plotted in Fig. 5.., The resulting lag spectra distribution is plotted in Fig. \ref{lags_test}.
 Phe drop in he lags at low and high frequencies is significant as the ags at the extreme frequencies fall far below the lower imit of the distribution of simulated. data., The drop in the lags at low and high frequencies is significant as the lags at the extreme frequencies fall far below the lower limit of the distribution of simulated data.
 We note. also. hat the artificial high-[requeney break in the lag spectrum. oduced by sampling ellects ancl noticed. by Craryetal.(L998) cannot explain the break we observe.," We note, also, that the artificial high-frequency break in the lag spectrum, produced by sampling effects and noticed by \citet{Crary} cannot explain the break we observe."
 The artificial ποσα should appear around 0.5fx where fy is the Nwquist requeney while. for the ddata we used. 0.5[x~2.610° Lz. and the break we observe in iis at an order of magnitude lower [requenev. al ~2.107 llz.," The artificial break should appear around $0.5\times f_{\rm
N}$ where $f_{\rm N}$ is the Nyquist frequency while, for the data we used, $0.5 \times f_{\rm N}
 \sim 2.6\times 10^{-3}$ Hz, and the break we observe in is at an order of magnitude lower frequency, at $\sim 2\times 10^{-4}$ Hz."
 At the low frequency end. the median of the distribution of simulated lag spectra does not reproduce the decreasing rend seen in the data and many data points fall. below he lower limit.," At the low frequency end, the median of the distribution of simulated lag spectra does not reproduce the decreasing trend seen in the data and many data points fall below the lower limit."
 The small scatter expected in the lag measurcments is partly due to the large laes intrinsic to the underlying model we assumed., The small scatter expected in the lag measurements is partly due to the large lags intrinsic to the underlying model we assumed.
 We repeated the same test. his time using the bending power law as the underlving lag spectrum.," We repeated the same test, this time using the bending power law as the underlying lag spectrum."
 The small lags at low frequencies. that this mocde ooduces. increases the low frequeney. scatter significantA," The small lags at low frequencies, that this model produces, increases the low frequency scatter significantly."
 'Phereore. a single bend at high frequencies canaccount for xh. bends. including the negative lag values of the data at he lowest frequencies.," Thereore, a single bend at high frequencies canaccount for both bends, including the negative lag values of the data at the lowest frequencies."
 Pherefore. a single-bend model. with constant time. lags below ~10.7n Hz is. consistent. with. the data.," Therefore, a single-bend model, with constant time lags below $\sim 10^{-4} $ Hz is consistent with the data."
 The median ancl extremes of the distribution of ag values for the bending power law mocel are shown in the xttom panel of Fig. 5.., The median and extremes of the distribution of lag values for the bending power law model are shown in the bottom panel of Fig. \ref{lags_test}.
 For completeness. we caleulated lag spectra using data rom the long monitoring campaign of performed: withINTE. described by Pounds 2001).," For completeness, we calculated lag spectra using data from the long monitoring campaign of performed with, described by \citet[e.g.][]{Pounds}."
. The lag spectra obtained from these delata are inconclusive however. as the low variability power low ~lo? llz produces a very weak signal in this requency range. making lag measurements too uncertain.," The lag spectra obtained from these data are inconclusive however, as the low variability power below $\sim 10^{-6}$ Hz produces a very weak signal in this frequency range, making lag measurements too uncertain."
" We conclude that the changes in slope of the lag spectrum at 2.10 ""and ~2.10+ Uz ave significant. so the spectrum is not consistent with a single power Law model."," We conclude that the changes in slope of the lag spectrum at $\sim
 2\times 10^{-5} $ and $\sim 2\times 10^{-4} $ Hz are significant, so the spectrum is not consistent with a single power law model."
 The spectrum is consistent with a bending power law but we cannot assess accurately the behaviour of the lags below the bend frequency., The spectrum is consistent with a bending power law but we cannot assess accurately the behaviour of the lags below the bend frequency.
 A model of constant lag up to 101 Ly bending to a f! dependence at high frequencies can reproduce the data well., A model of constant lag up to $10^{-4}$ Hz bending to a $f^{-4}$ dependence at high frequencies can reproduce the data well.
 The magnitude of the time lags between energy. bands ends to increase with energy separation. in both AGN and DIIXRDB (e.g.Papadakis.Nandra&Ixazanas2001:Nowaketal. 1999).," The magnitude of the time lags between energy bands tends to increase with energy separation, in both AGN and BHXRB \citep[e.g.][]{Papadakis_7469, Nowak_lags}."
. In 564. this ellect is clearly observable over he frequency range where significant lags can be measured.," In , this effect is clearly observable over the frequency range where significant lags can be measured."
 Figure 6. shows the lags between soft and medium and soft and hard bands. together with the corresponding bands.," Figure \ref{lags_bands} shows the lags between soft and medium and soft and hard bands, together with the corresponding bands."
 An increase of the energy. separation in the bands used to make the eross-spectrum. from a factor 3 to a factor of 6 dillerence in average energies. produces an increase in he lags by a factor of ~ 2. while preserving the shape of the ag spectra. within the uncertainties.," An increase of the energy separation in the bands used to make the cross-spectrum, from a factor 3 to a factor of 6 difference in average energies, produces an increase in the lags by a factor of $\sim 2$ , while preserving the shape of the lag spectra, within the uncertainties."
 ‘To investigate the energy dependence in more detail. we used the broader energy. bandpass of," To investigate the energy dependence in more detail, we used the broader energy bandpass of"
Hinode SOT Using MERLIN inversions we studied maps of inclination of the same region observed with IBIS a few hours later.,Hinode SOT Using MERLIN inversions we studied maps of inclination of the same region observed with IBIS a few hours later.
 Hinode SOT/SP Inversions were conducted at NCAR under the framework of the Community Spectro-polarimtetric Analysis Center (CSAC; http://www., Hinode SOT/SP Inversions were conducted at NCAR under the framework of the Community Spectro-polarimtetric Analysis Center (CSAC; ).
csac.hao.ucar.edu/)) The result is shown in fig. 1.., The result is shown in fig. \ref{inclination}.
 In fig., In fig.
" 3 (panel d) we show the chromospheric magnetic field inclination obtained from non-linear force-free extrapolations performed using the NLFF code (?),, after resolving the azimuth ambiguity by means of the code presented in ?!.. "," \ref{panels} (panel d) we show the chromospheric magnetic field inclination obtained from non-linear force-free extrapolations performed using the NLFF code \citep{2009ApJ...696.1780D}, after resolving the azimuth ambiguity by means of the code presented in \citet{2009ASPC..415..365L} ."
"Even though the low chromosphere is not a place where the force-free approximation holds, the observational results obtained in this work can easily be explained by our extrapolated model, suggesting the validity of such a scheme even in this context."," Even though the low chromosphere is not a place where the force-free approximation holds, the observational results obtained in this work can easily be explained by our extrapolated model, suggesting the validity of such a scheme even in this context."
" In the following we will focus on the analysis of the velocity field perturbations in two power spectral bands namely: 2.8—3.8 mHz and 4.8—5.8 mHz, hereafter the 5-minute and 3-minute bands respectively."," In the following we will focus on the analysis of the velocity field perturbations in two power spectral bands namely: $2.8-3.8$ mHz and $4.8-5.8$ mHz, hereafter the $5$ -minute and $3$ -minute bands respectively."
" Due to the limited extent of the time series data, the use of the FFT for the estimation of the power spectral density may result in a distorted estimation (?).."," Due to the limited extent of the time series data, the use of the FFT for the estimation of the power spectral density may result in a distorted estimation \citep{1970WRR.....6.1601E}."
" For this reason, we estimated the power spectral density using the Blackman-Tukey method and making use of a Barlett windowing function (2). "," For this reason, we estimated the power spectral density using the Blackman-Tukey method and making use of a Barlett windowing function \citep{blackman1958measurement}. ."
Fig.2 shows the power maps for the 3 and 5 minutes bands sampled by the Fe 617.3 nm line (panels a and c The power distribution in the two spectral bands look quite different., \ref{power} shows the power maps for the $3$ and $5$ minutes bands sampled by the Fe $617.3$ nm line (panels a and c The power distribution in the two spectral bands look quite different.
" The 5-minute oscillations tend to be spread throughout the FoV, while they are absent (absorbed) in the magnetic region."," The $5$ -minute oscillations tend to be spread throughout the FoV, while they are absent (absorbed) in the magnetic region."
 This is not surprising and a long series of similar results can be found in the literature (see ? and references therein)., This is not surprising and a long series of similar results can be found in the literature (see \citet{2009ApJ...706..909C} and references therein).
" The magnetic field is, in fact, able to scatter and absorb the acoustic field thus producing such an effect."," The magnetic field is, in fact, able to scatter and absorb the acoustic field thus producing such an effect."
" On theother hand, 3-minute oscillations are only found in the umbra of the pore."," On theother hand, $3$ -minute oscillations are only found in the umbra of the pore."
absorption line ancl another for the central emission.,absorption line and another for the central emission.
 In the case of emission lines with absorption cores. the absorption was also masked. before the Gaussian fit was attempted but only the emission ENLEM is given.," In the case of emission lines with absorption cores, the absorption was also masked before the Gaussian fit was attempted but only the emission FWHM is given."
 For comparison we have also included. the spectra of CGIx Per during its 1996 outburst (Morales-Itueda. Still lütoche 1999) and that of IP Pee. during its August 1994 outburst. (Morales-Itueda et al.," For comparison we have also included the spectra of GK Per during its 1996 outburst (Morales-Rueda, Still Roche 1999) and that of IP Peg, during its August 1994 outburst (Morales-Rueda et al."
 2000)., 2000).
 Phe spectrum. of LP Peg is very dillerent from that of the rest. of the cdwarf novae. the most striking difference being the strong {line in emission.," The spectrum of IP Peg is very different from that of the rest of the dwarf novae, the most striking difference being the strong line in emission."
Awe. As noted above. the red. continuum seen in the spectra of this system is probably the result of poor [lux calibration.," As noted above, the red continuum seen in the spectra of this system is probably the result of poor flux calibration."
Cyg. We took two spectra of this dwarf nova during the same outburst. one near maximum. the other in the descent to. quiescence.," We took two spectra of this dwarf nova during the same outburst, one near maximum, the other in the descent to quiescence."
 The spectra at. these two times show some major differences., The spectra at these two times show some major differences.
 At maximum the continuum increases significantlv., At maximum the continuum increases significantly.
 As a result of the cise becoming optically thick. we see absorption lines with an emission core.," As a result of the disc becoming optically thick, we see absorption lines with an emission core."
 The strength of the line is significantly larger relative to the strength. of the 3almer and lines compared to the quiescent. state., The strength of the line is significantly larger relative to the strength of the Balmer and lines compared to the quiescent state.
 In contrast. the spectrum taken near quiescence shows a lower continuum with Balmer and He emission lines.," In contrast, the spectrum taken near quiescence shows a lower continuum with Balmer and He emission lines."
 SS ςνο is one of the systems claimed to have spiral asvnunetries during outburst (Stcceghs 11996)., SS Cyg is one of the systems claimed to have spiral asymmetries during outburst (Steeghs 1996).
Gem. The spectrum shows Balmer absorption lines with faint emission cores., The spectrum shows Balmer absorption lines with faint emission cores.
 aappears strong in the spectrum and. couble peaked presumably due to the high inclination of the system., appears strong in the spectrum and double peaked presumably due to the high inclination of the system.
 There is also some emission in the Bowen blend., There is also some emission in the Bowen blend.
 In a recent study of U Gem during outburst Groot (2001) finds spiral structure in the accretion disc., In a recent study of U Gem during outburst Groot (2001) finds spiral structure in the accretion disc.
 Spectra taken on 2001 April 30. and May. 4.5. 9 and. LO when U Gom was in outburst. (first three nights) and decaving to quiescence (last two nights) show the presence of spiral structure in aand less clearly in wwhile the svstem was in outburst.," Spectra taken on 2001 April 30, and May 4, 5, 9 and 10 when U Gem was in outburst (first three nights) and decaying to quiescence (last two nights) show the presence of spiral structure in and less clearly in while the system was in outburst."
 The emission. line disappears almost completely from. the spectrum on May 9. and LO. when the system was half way down to its quiescent state (Stecghs. private Communication).," The emission line disappears almost completely from the spectrum on May 9 and 10, when the system was half way down to its quiescent state (Steeghs, private communication)."
 This indicates that the origin of iis either irracliation caused by a thickening of the accretion, This indicates that the origin of is either irradiation caused by a thickening of the accretion
"the maps are even higher than in maps that hadn't been rotated, indicating the presence of even stronger correlations.","the maps are even higher than in maps that hadn't been rotated, indicating the presence of even stronger correlations."
" This suggests that, at least for small rotations off the axis, the correlations are just as significant, if not more so."," This suggests that, at least for small rotations off the axis, the correlations are just as significant, if not more so."
" As an aside, the colour plots of the phase differences for Bianchi maps rotated by a number of different 0 in the range 0 to 27 were generated."," As an aside, the colour plots of the phase differences for Bianchi maps rotated by a number of different $\theta$ in the range 0 to $\pi$ were generated."
" These plots have been condensed together into which show that the correlations in the VIIo and V maps are visible across all 0 and for the VII, map are visible within about 7/3 of the preferred axis.", These plots have been condensed together into which show that the correlations in the $_0$ and V maps are visible across all $\theta$ and for the $_h$ map are visible within about $\pi/3$ of the preferred axis.
". Now to investigate the effect of noise, we considered three different types of noise."," Now to investigate the effect of noise, we considered three different types of noise."
 Firstly we tried the simplest form by just adding white noise to the Bianchi map., Firstly we tried the simplest form by just adding white noise to the Bianchi map.
 A map of random Gaussian noise (white noise) was generated., A map of random Gaussian noise (white noise) was generated.
 Using Healpix the spherical mode resolution was reduced to @ < 20., Using Healpix the spherical mode resolution was reduced to $\ell$ $\le$ 20.
 Then the “noise” map was modified to have zero mean and variance half that of the Bianchi map., Then the “noise” map was modified to have zero mean and variance half that of the Bianchi map.
 The second “noise” map was derived from a product available on the WMAP website which provides the effective number of observations per pixel., The second “noise” map was derived from a product available on the WMAP website which provides the effective number of observations per pixel.
 A map of random Gaussian noise was again generated., A map of random Gaussian noise was again generated.
 The variance was modified per pixel so that it was inversely proportional to the square of the number of observations in that pixel., The variance was modified per pixel so that it was inversely proportional to the square of the number of observations in that pixel.
 Using Healpix the spherical mode resolution was reduced to 4 x 20., Using Healpix the spherical mode resolution was reduced to $\ell$ $\le$ 20.
 Then the noise map was modified to have zero mean and variance half that of the Bianchi map., Then the noise map was modified to have zero mean and variance half that of the Bianchi map.
 The final “noise” map used a simulation of ACDM fluctuations of the CMB (as performed by Eriksenetal. (2005)))., The final “noise” map used a simulation of $\Lambda$ CDM fluctuations of the CMB (as performed by \cite{Eriksen2005}) ).
 Again the noise map was modified to reduce the spherical mode resolution to / < 20 and have variance half that of the Bianchi map., Again the noise map was modified to reduce the spherical mode resolution to $\ell$ $\le$ 20 and have variance half that of the Bianchi map.
 Each of these “noise” maps was added to each of the rotated Bianchi maps., Each of these “noise” maps was added to each of the rotated Bianchi maps.
 We see from the example in Figure 6 that the spherical harmonic coefficients derived still have visible correlations in the phases for the Bianchi V map., We see from the example in Figure \ref{figRotNoisePhaseDif} that the spherical harmonic coefficients derived still have visible correlations in the phases for the Bianchi V map.
 The results of, The results of
axis and one quarter-wave rhomb) are employed to perform a very achromatie polarimetric analysis over the whole spectral domain.,axis and one quarter-wave rhomb) are employed to perform a very achromatic polarimetric analysis over the whole spectral domain.
 They are followed by a Wollaston prism which splits the incident light into two beams. respectively containing light linearly polarized perpendicular/parallel to the axis of the prism.," They are followed by a Wollaston prism which splits the incident light into two beams, respectively containing light linearly polarized perpendicular/parallel to the axis of the prism."
 The two beams produced by the Wollaston prism are imaged onto the two optical fibres that carry the light to the spectrograph., The two beams produced by the Wollaston prism are imaged onto the two optical fibres that carry the light to the spectrograph.
 Each Stokes V spectrum is obtained from a combination of four sub-exposures taken with the half-wave rhombs oriented at different azimuths (Semel et al., Each Stokes V spectrum is obtained from a combination of four sub-exposures taken with the half-wave rhombs oriented at different azimuths (Semel et al.
 1993)., 1993).
 The data reduction is performed by Libre-Esprit. a dedicated. fully automated software described by Donati et al. (," The data reduction is performed by Libre-Esprit, a dedicated, fully automated software described by Donati et al. ("
1997) and implementing the optimal spectral extraction principle of Horne (1986) and Marsh (1989).,1997) and implementing the optimal spectral extraction principle of Horne (1986) and Marsh (1989).
 A total of 838 spectra was recorded from July 2008 to October 2009. during 4 different telescope campaigns.," A total of 838 spectra was recorded from July 2008 to October 2009, during 4 different telescope campaigns."
 The first data set. taken with NARVAL during 4 consecutive nights in July 2008. is described by L09.," The first data set, taken with NARVAL during 4 consecutive nights in July 2008, is described by L09."
 To complement this first time-series. we have recorded another set of 80 NARVAL spectra in. June/July 2009. followed by 316 ESPaDOnS spectra in September 2009 and 146 NARVAL spectra in October/November 2009.," To complement this first time-series, we have recorded another set of 80 NARVAL spectra in June/July 2009, followed by 316 ESPaDOnS spectra in September 2009 and 146 NARVAL spectra in October/November 2009."
 The observations of July 2008 and September 2009 dominate all other available data due to their dense temporal sampling over consecutive nights. by their high and homogeneous signal-to-noise ratio and by the large number of spectra collected at these two epochs (about of the observing material at our disposal).," The observations of July 2008 and September 2009 dominate all other available data due to their dense temporal sampling over consecutive nights, by their high and homogeneous signal-to-noise ratio and by the large number of spectra collected at these two epochs (about of the observing material at our disposal)."
 For this reason. these two data sets will be preferentially used in the rest of our analysis.," For this reason, these two data sets will be preferentially used in the rest of our analysis."
 To ensure an optimal data quality in our study. we have discarded from our data sets all spectra in which a low rreveals significant atmospheric absorption or tracking problems.," To ensure an optimal data quality in our study, we have discarded from our data sets all spectra in which a low reveals significant atmospheric absorption or tracking problems."
 We have also ignored all spectra with significant solar contamination in the Stokes | parameter (observations collected close to sunrise)., We have also ignored all spectra with significant solar contamination in the Stokes I parameter (observations collected close to sunrise).
 After cleaning up the data sets. we end up with a total of useful 799 spectra (Table 2.1)).," After cleaning up the data sets, we end up with a total of useful 799 spectra (Table \ref{tab:obs}) )."
 The integration time adopted for the four individual subexposures constituting the Stokes V sequences is varying from one observing run to the next. with values ranging from 4 sec in September 2009 to 16 sec in June 2009.," The integration time adopted for the four individual subexposures constituting the Stokes V sequences is varying from one observing run to the next, with values ranging from 4 sec in September 2009 to 16 sec in June 2009."
 To this shutter time. we must add another 120 sec for each Stokes V sequence. including readout time and rotation of the polarimetric optics.," To this shutter time, we must add another 120 sec for each Stokes V sequence, including readout time and rotation of the polarimetric optics."
 The total time spent to obtain a Stokes V spectrum is therefore comprised between 136 sec and 184 sec. depending on the observing run.," The total time spent to obtain a Stokes V spectrum is therefore comprised between 136 sec and 184 sec, depending on the observing run."
 The Least-Squares Deconvolution technique (LSD: Donati et al., The Least-Squares Deconvolution technique (LSD; Donati et al.
 1997) was applied to all spectra. extracting from each of the 799 spectra a mean line profile with enhanced.," 1997) was applied to all spectra, extracting from each of the 799 spectra a mean line profile with enhanced."
S/N.. Different line lists were employed to check the robustness of our results. but unless specifically mentioned hereafter. the line list used in this paper is identical to that presented by L09.," Different line lists were employed to check the robustness of our results, but unless specifically mentioned hereafter, the line list used in this paper is identical to that presented by L09."
 The mask is based on a Kurucz atmospheric model with an effective temperature Tay=10.000 K. a surface gravity log(g)=4.0 and a solar metallicity. yielding a total of about 1.100 atomic lines in the spectral window of NARVAL and ESPaDOnS. The resulting oof Stokes V LSD cross-correlation profiles is listed in Table 2.]..," The mask is based on a Kurucz atmospheric model with an effective temperature $_{\rm eff} = 10,000$ K, a surface gravity $\log(g)=4.0$ and a solar metallicity, yielding a total of about 1,100 atomic lines in the spectral window of NARVAL and ESPaDOnS. The resulting of Stokes V LSD cross-correlation profiles is listed in Table \ref{tab:obs}."
 We do not detect any Zeeman signature in any of the individual Stokes V LSD profiles. in agreement with the previous analysis of LOY.," We do not detect any Zeeman signature in any of the individual Stokes V LSD profiles, in agreement with the previous analysis of L09."
 As a strategy to further improve the, As a strategy to further improve the
the case of M8BG. we sce a strong signature of boxy Al components in the inner parts of the ealaxy. aud restrict our study of substructure to regious at larger radius.,"the case of M86, we see a strong signature of boxy A4 components in the inner parts of the galaxy, and restrict our study of substructure to regions at larger radius."
 A great deal of the area near M89 is covered iu substructure., A great deal of the area near M89 is covered in substructure.
 Easily visible are the Πο feature to the west (Reeion 1) and the extremely bright shell to the south (Region 2)., Easily visible are the “jet” feature to the west (Region 1) and the extremely bright shell to the south (Region 2).
 We also find plumes to the northwest (Region 3). Ες (Reeiou L). aud southeast (Reeiou 5) of M89.," We also find plumes to the northwest (Region 3), southwest (Region 4), and southeast (Region 5) of M89."
 Fainter aud further away are two shells: oue (C12 kpe) to the east (Region 6). aud another (57 kpc) to the northeast (Region 7).," Fainter and further away are two shells: one (42 kpc) to the east (Region 6), and another (57 kpc) to the northeast (Region 7)."
 As with the previous galaxies. we have confirmed that all these features ave visible in the unsubtracted image as well.," As with the previous galaxies, we have confirmed that all these features are visible in the unsubtracted image as well."
 Photometric measurements were made of cach of the detected features and compared with directly adjacent sky regions., Photometric measurements were made of each of the detected features and compared with directly adjacent sky regions.
 Table 6 gives the photometric properties for cach feature.," Table \ref{m89tab}
 gives the photometric properties for each feature."
 The variety of tidal structure around M89 argues for a complicated accretion historv., The variety of tidal structure around M89 argues for a complicated accretion history.
 The W Tail (Region 1) is peculiar in that it has a high. relatively coustaut surface brightuess across its lensth. aud may be a long tidal tail seen curving back onu itself iun projection.," The W Tail (Region 1) is peculiar in that it has a high, relatively constant surface brightness across its length, and may be a long tidal tail seen curving back on itself in projection."
 Indeed there is some hint of a plume extending to the cast of M89 along the same axis as the W Tail. projecting across the outer shells if real. this could be the extension of the W tail curving— back across the face of the galaxy.," Indeed there is some hint of a plume extending to the east of M89 along the same axis as the W Tail, projecting across the outer shells – if real, this could be the extension of the W tail curving back across the face of the galaxy."
 Uulike the shell svstem of ΑΔΕΙΟ. MS9's shells do not have the classic aligned auc iuterleaved structure expected from phase wrapping from a small accretion eveut Quinn 1981).," Unlike the shell system of M49, M89's shells do not have the classic aligned and interleaved structure expected from phase wrapping from a small accretion event Quinn 1984)."
 hDustead. the shells occur at a variety of position angles aud radii. suggesting material that has conie into the ealaxy with a range of angular moneuta.," Instead, the shells occur at a variety of position angles and radii, suggesting material that has come into the galaxy with a range of angular momenta."
 Such features could arise either from inultiple accretions. or Via a major merger of two disk galaxies. where tidal material can spatially wrap around the remnant without any co-aliguimeut of shells Heruquist. Sperecl 1992: IHibbard Aihos 1995).," Such features could arise either from multiple accretions, or via a major merger of two disk galaxies, where tidal material can spatially wrap around the remnant without any co-alignment of shells Hernquist Spergel 1992; Hibbard Mihos 1995)."
 The sharpness of the outer shells (Reeious 6 and 7) also argues that. like M19. AISO9 1uust not have expericuced sieuificaut tidal strippiug from the Vireo cluster euvirounieut over the past Cr or so. otherwise the shells would have been disturbed or ciszuptect.," The sharpness of the outer shells (Regions 6 and 7) also argues that, like M49, M89 must not have experienced significant tidal stripping from the Virgo cluster environment over the past Gyr or so, otherwise the shells would have been disturbed or disrupted."
 Iu sununary. we lave used deep. wide-Bekl surface photometry to study the extended cuvelopes of five huninous Vireo ellipticals.," In summary, we have used deep, wide-field surface photometry to study the extended envelopes of five luminous Virgo ellipticals."
 We use the IRAF ELLIPSE task to fit elliptical isophotes to the galaxies! surface brightness profiles., We use the IRAF ELLIPSE task to fit elliptical isophotes to the galaxies' surface brightness profiles.
 Analytic fits to these isoplotal models compare well to fits published elsewhere in the literature., Analytic fits to these isophotal models compare well to fits published elsewhere in the literature.
 Subtracting the isophotal models from the nuages. we identify a varicty of low surface brightuess tidal features streams. plumes. aud shells in the outer envelopes of these elliptical galaxies. which we lave quantified iu terms of their total luninosity aud peals surface brightness.," Subtracting the isophotal models from the images, we identify a variety of low surface brightness tidal features – streams, plumes, and shells – in the outer envelopes of these elliptical galaxies, which we have quantified in terms of their total luminosity and peak surface brightness."
 We fud that ALS? is characterized by an extended diffuse halo with broad plumes aud radial streams. but no sharp tidal shells or loops.," We find that M87 is characterized by an extended diffuse halo with broad plumes and radial streams, but no sharp tidal shells or loops."
 M86 shows a umber of small. relatively high surface strenuis. while ALSlt shows no evidence for significant substructure )evond its μποτ] elliptical isophotes.," M86 shows a number of small, relatively high surface streams, while M84 shows no evidence for significant substructure beyond its smooth elliptical isophotes."
" Iu contrast. botl ΑΤΙ and M89 show complex (and in ALLO. hitherto unudiscovered) systems of distinct shells aud other tidal catures,"," In contrast, both M49 and M89 show complex (and in M49, hitherto undiscovered) systems of distinct shells and other tidal features."
 The variety of structure we sec in these Vireo ellipticals uav well reflect differences iu their history of accretion of snaller galaxies. their accretion history iuto the Vireo Cluster itself. or a combination of the two.," The variety of structure we see in these Virgo ellipticals may well reflect differences in their history of accretion of smaller galaxies, their accretion history into the Virgo Cluster itself, or a combination of the two."
 To oen our discussion of environmental influcuces. we show in Figure ο three orthogonal projections of the hnree dimensional positious of our sample galaxies iu he Virgo cluster.," To begin our discussion of environmental influences, we show in Figure \ref{virgogeom} three orthogonal projections of the three dimensional positions of our sample galaxies in the Virgo cluster."
 We use the distances to cach galaxy derived by Mei (2007) using the surface briglhtuess Huctuation technique., We use the distances to each galaxy derived by Mei (2007) using the surface brightness fluctuation technique.
 We adopt this dataset largolv )ecause it has cousisteutlv derived distances for all of our ealaxies. with very low internal errors.," We adopt this dataset largely because it has consistently derived distances for all of our galaxies, with very low internal errors."
 However. there are systematic uncertainties in the absolute cdistauces: other techniques vield somewhat different results.," However, there are systematic uncertainties in the absolute distances; other techniques yield somewhat different results."
 Most notably. the position of M86 is quite uncertain — while the Mei distances place M86 in the cluster core. in frout of M81. distance determinatious derived from planetary nebula put both M86 and ALS αἲ the same distance. 1 Mpe belund M87 (Jacoby 1990).," Most notably, the position of M86 is quite uncertain – while the Mei distances place M86 in the cluster core, in front of M84, distance determinations derived from planetary nebula put both M86 and M84 at the same distance, 1 Mpc behind M87 (Jacoby 1990)."